{
    "0207/astro-ph0207371_arXiv.txt": {
        "abstract": "The characteristic of the solar acoustic spectrum is such that mode lifetimes get shorter and spatial leaks get closer in frequency as the degree of a mode increases for a given order. A direct consequence of this property is that individual $p$-modes are only resolved at low and intermediate degrees, and that at high degrees, individual modes blend into ridges. Once modes have blended into ridges, the power distribution of the ridge defines the ridge central frequency and it will mask the true underlying mode frequency.  An accurate model of the amplitude of the peaks that contribute to the ridge power distribution is needed to recover the underlying mode frequency from fitting the ridge.  We present the results of fitting high degree power ridges (up to $\\ell = 900$) computed from several two to three-month-long time-series of full-disk observations taken with the Michelson Doppler Imager (MDI) on-board the Solar and Heliospheric Observatory between 1996 and 1999.  We also present a detailed discussion of the modeling of the ridge power distribution, and the contribution of the various observational and instrumental effects on the spatial leakage, in the context of the MDI instrument. We have constructed a physically motivated model (rather than some {\\em ad hoc} correction scheme) resulting in a methodology that can produce an unbiased determination of high-degree modes, once the instrumental characteristics are well understood.  Finally, we present changes in high degree mode parameters with epoch and thus solar activity level and discuss their significance. ",
        "introduction": "Since from a single vantage point we can only observe a bit less than half of the solar surface, helioseismic power spectra computed for a specific target mode with degree $\\ell$ and azimuthal order $m$ also contains power from modes with different -- and usually nearby -- $\\ell$ and $m$ values. The presence of these unwanted modes, or spatial leaks, complicates the fitting of the resulting observed spectra and degrade the mode parameter estimates, especially when the leaks have frequencies similar to that of the target mode. As mode lifetimes get shorter and spatial leaks get closer in frequency (\\ie, $d\\nu/d\\ell$ becomes small), individual p-modes can no longer be resolved. This mode blending occurs around $\\ell=150$ for p-modes with frequency near $3.3$ mHz and around $\\ell=250$ for the f-modes. Once individual modes blend into ridges, as illustrated in Figure~\\ref{fig:spectra}, the power distribution of the ridge masks the true underlying mode frequency and the ridge central frequency is not a good estimate of the target mode frequency. Moreover the amplitudes of the spatial leaks have been shown to be asymmetric \\cite[see for example][]{korzennik99}. A direct consequence of this leakage asymmetry is to offset the power distribution of the ridges. Such offset results in a significant difference between the central frequency of the ridge and the frequency of the targeted individual mode. To recover the underlying mode frequency from fitting the ridge, an accurate model of the amplitude of the peaks that contribute to the ridge power distribution (\\ie, the leakage matrix) is needed.  The lack of unbiased determinations of mode frequencies and frequency splittings at high degrees has so far limited the use of such high-degree data in helioseismic inversions for constraining the near-surface structure and dynamics of the sun. Since experiments like the Michelson Doppler Imager (MDI) on board the Solar and Heliospheric Observatory (SOHO), with a two-arcsec-per-pixel spatial resolution in full-disk mode, allows us to detect oscillation modes up to $\\ell \\approx 1500$ \\citep{scherrer95}, only a small fraction of the observed modes are currently used.  High-degree modes are trapped near the solar surface: for example, the lower turning point of a mode of degree $\\ell = 500$ and frequency around 3 mHz is at 0.99 of the solar radius. This makes them exceptional diagnostic tools to probe the near-surface region of the Sun, a region of great interest. Indeed, it is there that the effects of the equation of state are felt most strongly, and that dynamical effects of convection and processes that excite and damp the solar oscillations are predominantly concentrated. \\citet{rabello-soares00} have shown that the inclusion of high-degree modes (\\ie, $\\ell$ up to 1000) has the potential to improve dramatically the inference of the sound speed in the outermost 2 to 3\\% of the solar radius.  Furthermore, inversion of artificial mode frequency differences resulting from models computed with two different equations of state (\\ie, MHD and OPAL) recovered the intrinsic difference in $\\Gamma_1$, the adiabatic exponent, throughout the second helium ionization zone and well into the first helium and hydrogen ionization zones, with error bars far smaller than the differences resulting from using two different equations of state.  These tests were carried out using the relatively large observational uncertainties resulting from ridge fitting, but with the implicit assumption that systematic errors were not present.  These tests show that we can probe subtle effects in the thermodynamic properties of this region, but only when including such high-degree modes.  The first estimates of high-degree mode frequencies used $m$-averaged Big Bear Solar Observatory data \\citep{libbrecht88}. To recover the underlying mode frequency from fitting a given ($n, \\ell$) ridge, they used a simple Gaussian profile as an approximation for the $m$-averaged leakage matrix. Namely: \\begin{equation} C_r^2(\\ell, \\ell') = \\exp({-(\\frac{\\Delta\\ell - \\epsilon\\,\\ell}{2s})^2}) \\label{eq:klCr} \\end{equation} where $\\Delta\\ell = \\ell' - \\ell$ while the $\\epsilon\\,\\ell$ term represents the leakage asymmetry introduced by an image scale error of a fraction $\\epsilon$ of the image size.  The ridge centroid frequency was estimated using a simple weighted average: \\begin{equation} \\tilde{\\nu}_{n,\\ell} = \\frac{\\sum_{\\ell'} C_r^2(\\ell, \\ell') \\, A_{n,\\ell'} \\, \\nu_{n,\\ell'}} {\\sum_{\\ell'} C_r^2(\\ell, \\ell') \\, A_{n,\\ell'}} \\end{equation} where $A_{n,\\ell'}$ is the individual mode power amplitude and $\\nu_{n,\\ell'}$ is the mode frequency. The frequency difference between the ridge and the mode frequency, $\\Delta \\nu_{n,\\ell} = \\tilde{\\nu}_{n,\\ell}(\\mbox{ridge}) - \\nu_{n,\\ell}(\\mbox{mode})$, can thus be estimated using a parametric representation of the leakage coefficients.  For intermediate degree modes ($50 < \\ell < 150$) this frequency difference can be directly measured by reducing the frequency resolution of the observed power spectra as to force individual modes to blend into ridges. Using observed $\\Delta \\nu$, the parameters $s$ and $\\epsilon$ can be calibrated at these intermediate degrees and the correction extrapolated to high-degree modes.  As \\citet{libbrecht88} state in their paper, this is only a first step and there is substantial room for improvement particularly at high $\\ell$. Following the same idea, but improving on the method, \\citet{korzennik90} and \\citet{rhodes99} estimated high-degree mode frequencies using Mount Wilson and MDI data respectively. In contrast, \\citet{bachmann95}, using data from the High-L Helioseismometer at Kitt Peak, calculated the $m$-averaged leakage matrix, but neglecting the horizontal components, in order to estimate high-degree mode frequencies.  In this paper, we present the results of an extensive study of the various elements that contribute to the precise value of the effective leakage matrix, that in turn is key to the precise determination of high degree modes.  We have attempted to construct a physically motivated model --- rather than an {\\em ad hoc} correction scheme --- in order to produce an unbiased determination of the high-degree modes. Since \\citet{korzennik99} has shown that the inclusion of the horizontal component of the leakage matrix calculation partially explains its observed asymmetry, we have included the horizontal component in all our leakage matrix calculations.  In Section 2, we describe the data we used and how we computed and fitted the power spectra used in this work.  In Section 3 we present and discuss ridge modeling for high degree modes while in Section 4 we address the issue of estimating the ratio between the radial and horizontal components. In Section 5 we discuss the instrumental effects specific to MDI that must be included to properly model the ridge power distribution. Finally, in Section 6, we present the results of our analysis followed by our conclusions. ",
        "conclusions": "We believe that we have shown that one can construct a physically motivated model (rather than some {\\em ad hoc} correction scheme) of the ridge power distribution that results in a methodology that can produce an unbiased determination of high-degree modes, once the instrumental characteristics are well understood and precisely measured.  We have gained substantial insight in the understanding and modeling of the instrumental effects of the Michelson Doppler Imager, including plate scale error, image distortion, point spread function and image orientation.  Now that we better understand the MDI instrument imaging imperfections, we can -- and will in the near future -- reprocess the {\\em Dynamics} data, and include the correct image scale and the known image distortion in the spatial decomposition.  We have produced an error budget for the characterization of high degree mode parameters. Beyond MDI {\\em Dynamics} observations, this has direct applications on the characterization of the GONG+ instruments and the reduction of the GONG+ observations as well as on the characteristics of the Helioseismic/Magnetic Imager (HMI) instrument planned for the Solar Dynamics Observatory (SDO) mission, if one wishes to make effective use of the high degree modes accessible to these experiments.  Last but not least, we also see changes with solar activity in the high degree frequencies and the asymmetry of the $f$-mode. Our attempt to see changes in splitting coefficients and mode linewidths remained inconclusive."
    },
    "0207/gr-qc0207100_arXiv.txt": {
        "abstract": " ",
        "introduction": "In its \\wfs\\ \\GR\\ predicts that, among other things,  the orbit of a test particle freely falling in the gravitational field of a central rotating body is affected by the so called \\grc\\ dragging of the inertial frames or \\leti\\ \\ef. More precisely,  the longitude of the ascending \\nd\\ $\\Omega$ and the argument of the \\pg\\ $\\omega$ of the orbit  undergo tiny \\pc s according to [\\textit{Lense and Thirring}, 1918] \\eqia \\dot\\Omega \\lt & = & \\frac{2GJ}{c^{2}a^{3}(1-e^{2})^{\\frac{3}{2}}},\\\\ \\dot\\omega \\lt & = & -\\frac{6GJ\\cos{i}}{c^{2}a^{3}(1-e^{2})^{\\frac{3}{2}}},\\eqfa in which $G$ is the Newtonian gravitational constant, $J$ is the proper angular momentum of the central body, $c$ is the speed of light $in\\ vacuum$, $a,\\ e$ and $i$ are the \\sa, the \\ec\\ and the \\ic, respectively, of the orbit of the test particle. The \\leti\\ precessions for the \\lg\\ satellites amount to \\eqia \\dot\\Omega\\lt^{\\rm LAGEOS}& = & 31\\ \\textrm{mas/y},\\\\ \\dot\\Omega\\lt^{\\rm LAGEOS\\ II} & = & 31.5\\ \\textrm{mas/y},\\\\ \\dot\\omega\\lt^{\\rm LAGEOS} & = & 31.6\\ \\textrm{mas/y},\\\\ \\dot\\omega\\lt^{\\rm LAGEOS\\ II} & = & -57\\ \\textrm{mas/y}. \\eqfa  The first measurement of this \\ef\\ in the gravitational field of the \\et\\ has been obtained by analyzing a suitable combination of the laser-ranged data to the existing passive geodetic \\st s \\lg\\ and \\lgg\\ [\\textit{Ciufolini et al.,} 1998]. The observable [{\\it Ciufolini}, 1996] is a linear trend with a slope of 60.2 milliarcseconds per year (mas/y in the following) and includes the residuals of the nodes of \\lg\\ and \\lgg\\ and the \\pg\\ of \\lgg\\footnote{The \\pg\\ of \\lg\\ was not used because it introduces large observational errors due to the smallness of the \\lg\\ \\ec\\ [{\\it Ciufolini}, 1996] which amounts to 0.0045.}. The total relative accuracy of the measurement of the solve-for parameter $\\mlt$, introduced in order to account for this \\grl\\ \\ef, is of the order of $2\\times 10^{-1}$ [{\\it Ciufolini et al.}, 1998].  In this kind of satellite--based space experiments the major source of \\se s is represented by the aliasing trends due to the classical secular precessions [\\textit{Kaula}, 1966] of the \\nd\\ and the \\pg\\ induced by the mismodelled \\zh\\ of the \\gp\\ $J_2,\\ J_4,\\ J_6,...$ Indeed, according to the present knowledge of the \\et's gravity field based on the EGM96 model [\\textit{Lemoine et al.}, 1998], they amount to a large part of the \\grc\\ precessions of interest, especially for the first two even zonal harmonics. In the performed LAGEOS--LAGEOS II Lense--Thirring experiment the adopted observable allows for the cancellation of the static and dynamical effects of $J_2$ and $J_4$. The remaining higher degree even zonal harmonics affects the measurement at a $13\\%$ level.  In order to achieve a few percent accuracy, in [\\textit{Ciufolini}, 1986] it was proposed to launch a passive geodetic laser-ranged \\st- the former {\\rm LAGEOS} III - with the same orbital parameters of \\lg\\ apart from its inclination which should be supplementary to that of \\lg.  This orbital configuration would be able to cancel out exactly the classical \\nl\\ \\pc s, which are proportional to $\\cos i$, provided that the observable to be adopted is the sum of the residuals of the \\nl\\ \\pc s of {\\rm LAGEOS} III and LAGEOS \\eqi \\delta\\dt\\Omega^{{\\rm III}}+\\delta\\dt\\Omega^{{\\rm I}}=62\\mlt.\\lb{lares}\\eqf Later on the concept of the mission slightly changed. The area-to-mass ratio of {\\rm LAGEOS} III was reduced in order to make less relevant the impact of the non-gravitational perturbations  and the eccentricity  was enhanced in order to be able to perform other \\grl\\ tests: the LARES was born [\\textit{Ciufolini}, 1998]. The orbital parameters of \\lg, \\lgg\\ and LARES are in Table 1.  \\begin{table}[ht!] \\caption{Orbital parameters of \\lg, \\lgg\\, LARES and POLARES.} \\label{para} \\begin{center} \\begin{tabular}{llllll} \\noalign{\\hrule height 1.5pt} Orbital parameter & \\lg & \\lgg & LARES & POLARES\\\\ \\hline $a$ (km) & 12,270 & 12,163 & 12,270 & 8,378\\\\ $e$ & 0.0045 & 0.014 & 0.04 & 0.04\\\\ $i$ (deg) & 110 & 52.65 & 70 & 90\\\\ \\noalign{\\hrule height 1.5pt} \\end{tabular} \\end{center} \\end{table}  Recent developments of the concept of the twin satellites in supplementary orbits have led to the discover of new, possible gravitomagnetic observables based on the use of the perigees as well [{\\it Iorio}, 2002a] and of unexpected connections with  the gravitomagnetic clock effect [{\\it Iorio and Lichtenegger}, 2002].  Unfortunately, at present we do not know if the LARES mission will be approved by any space agency. Although much cheaper than other proposed and/or approved  space--based missions, funding is the major obstacle in implementing the LARES project. The most expensive part is the launching segment.  Very recently the possibility of launching the LARES satellite into an orbit with $a=8,378$ km, $i=90$ deg, $e=0.04$ and using as observable its node has been considered [{\\it Lucchesi and Paolozzi}, 2001].  In the following we will name POLARES the LARES satellite in such proposed polar orbit. The Lense--Thirring secular rate of the POLARES node would amount to 96.9 mas/y. Then, the observable would be \\eqi\\delta\\dot\\Omega^{\\rm PL}=96.9\\mu_{\\rm LT},\\eqf i.e. a linear secular trend with a slope of 96.9 mas/y. The choice of a so low altitude is motivated by the need of using a cheap rocket launcher; so, it must certainly be considered as an admirable and important further effort towards the practical realization of the LARES project. The polar orbit would allow to prevent the aliasing effects of the mismodelled classical secular precessions of the node induced by the even $l=2n$ zonal $m=0$ coefficients of the multipolar expansion of the static part of the geopotential. The non--gravitational perturbations [{\\it Lucchesi}, 2001; 2002; {\\it Lucchesi and Paolozzi, } 2001] would represent a minor problem.  In this paper we wish to critically analyze such important and interesting evolution of the LARES concept. ",
        "conclusions": "In this paper the proposal of putting the LARES satellite into a polar, elliptical orbit with an altitude of 2,000 km in order to look at its gravitomagnetic secular node shift has been critically analyzed.  The key point is that the mismodelled classical nodal secular precessions induced by the even zonal coefficients of the multipolar expansion of the terrestrial gravitational field vanished if and only if the inclination of the satellite would be exactly $90$ deg. Of course, mainly due to possible orbital injection errors, this could never happen. It turns out that the low altitude of the proposed orbital configuration, and the consequent high sensitivity to the higher even degree zonal terms of the geopotential, would greatly enhance the impact of even small departures of the real values of the inclination from the nominal value of 90 deg. For example, for just $i_{\\rm PL}=90\\pm 0.2$ deg the systematic gravitational error would be of the order of 5$\\%$--10$\\%$, according to the EGM96 Earth gravity model up to $l=20$, which is of the same order of magnitude of the present LAGEOS--LAGEOS II experiment (Its total error, including various systematic gravitational and non--gravitational perturbations, is of the order of 20$\\%$--30$\\%$). Of course, such a situation would be further made critical by the possible use of a low--cost launcher which, unavoidably, would induce not negligible orbital injection errors. Moreover, the tesseral $K_1$ tidal perturbation, which has the same period of the satellite's node, would induce a secular aliasing trend over an observational time span of a few years because its period would amount to several tens of years for near polar orbits.  Another important point is that the proposed POLARES would not yield substantial improvements also in the context of the combined residuals approach which allows to cancel out the contribution of the first even zonal coefficients of the geopotential irrespectively of the inclination of the satellites. Indeed, it turns out that a combination including the nodes of LAGEOS, LAGEOS II, POLARES and the perigees of LAGEOS II and POLARES would be not defined for $i_{\\rm PL}=90$ deg because the coefficient weighing the node of the LARES would go to infinity. For very small deviations of $i_{\\rm PL}$ from such a critical value the systematic error induced by the remaining even zonal harmonics would amount to 10$\\%$--20$\\%$. Moreover, the coefficient with which the node of POLARES would enter the combination would be much more larger than unity and would greatly enhance the impact of the gravitational and non--gravitational time--dependent perturbations. Even the use of the nodes of LAGEOS and LAGEOS II and the perigees of LAGEOS II and POLARES, in which case $i_{\\rm PL}=90$ deg would not create problems, would not yield benefits because the systematic gravitational error would be of the order of 20$\\%$--30$\\%$.  Moreover, the accuracy of the practical data reduction  from such version of LARES would be affected by the atmospheric drag.  Finally, we can conclude that the proposed low--cost version of the LARES mission would not yield any significant improvements in the measurement of the elusive Lense--Thirring effect, according to the present--day level of knowledge of the Earth's gravitational field. Perhaps, the situation could improve to a certain extent with the new, more accurate gravity models from CHAMP and GRACE missions which should become available in the next few years. On the contrary, the original concept of the couple of supplementary satellites would deserve grater attention thanks to its much richer spectrum of high accuracy relativistic observables."
    },
    "0207/astro-ph0207147_arXiv.txt": {
        "abstract": "{Recent efforts to account for the observed $L_X - T_X$ relation of galaxy clusters has led to suggestions that the ICM has an apparent ``entropy floor'' at the level of $K_0 \\gtrsim 300$ keV cm$^2$.  Here, we propose new tests based on the thermal SZ effect and on the $M_{gas} - T_X$ trend (from X-ray data) to probe the level of the excess entropy in the ICM.  We show that these new tests lend further support to the case for a high entropy floor in massive clusters.}   \\addkeyword{Cosmology: Theory} \\addkeyword{X-rays: Galaxies: Clusters}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Relationships between the global X-ray properties of clusters have proven to be important probes of the intracluster medium (ICM).  Case in point are studies of the X-ray luminosity ($L_X$) - mean emission-weighted gas temperature ($T_X$) relation. Theoretical models that include only the effects of gravity and shock heating (self-similar models) predict $L_X \\propto T_X^2$, yet the observed relation is $L_X \\propto T_X^{2.6-3.0}$ (e.g., Markevitch 1998; Allen \\& Fabian 1998).  This discrepancy has prompted a number of theorists to consider alternative models.  Both the effects of heating (e.g., Kaiser 1991; Evrard \\& Henry; Wu et al.\\@ 2000; Babul et al.\\@ 2002) and radiative cooling (e.g., Bryan 2000; Voit \\& Bryan 2001; Dav\\'{e} et al.\\@ 2002) have been examined.  These studies find that heating and/or cooling introduces a core into the entropy profiles of clusters (an ``entropy floor'') which, in turn, results in a steepening of the $L_X - T_X$ relation (as required).  Ponman and collaborators (Ponman et al.\\@ 1999; Lloyd-Davies et al.\\@ 2000) have presented direct evidence for an entropy floor in nearby groups.  Investigations of the ICM are not limited to the $L_X - T_X$ relation, however.  Other cluster observables can be used as alternative probes of the ICM.  For example, McCarthy et al.\\@ (2002) have studied the effects of an entropy floor on the cluster gas mass ($M_{gas}$) - $T_X$ relation and, through a detailed comparison with observations, have placed stringent limits on the entropy floors in nearby massive clusters.  Because the $M_{gas} - T_X$ relation is derived from X-ray data, the test provides a valuable self-consistency check of the $L_X - T_X$ results.  The main results of that study are presented below.  Ultimately, however, scaling relations that are independent of the $L_X - T_X$ and $M_{gas} - T_X$ relations are desirable.  The thermal Sunyaev-Zeldovich (SZ) effect, which has a different dependence on the entropy of the ICM than does the X-ray emission, can be used for such a purpose (McCarthy et al.\\@ in preparation).  Here, we derive a relation between the central and integrated cluster Compton parameters (which are both proportional to the SZ effect), analyze how this relation is affected by an entropy floor, and compare the results to recent SZ effect observations.  This is the first time the SZ effect has been used as a probe of the entropy floors of clusters.  \\begin{figure}[!t] \\includegraphics[width=\\columnwidth]{ababul_fig1.eps} \\caption{A comparison of $M_{gas}(r_{500}) - T_X$ relations.  The squares represent the observations of Mohr et al. (1999).  The dotted line is the self-similar result. The short-dashed, long-dashed, dot-dashed, and solid lines represent the models with entropy floor constants of $K_0$ = 100, 200, 300, and 427 keV cm$^2$, respectively.} \\label{fig1} \\end{figure} ",
        "conclusions": "\\begin{figure}[!t] \\includegraphics[width=\\columnwidth]{ababul_fig2.eps} \\caption{A comparison of $y_0 - S_{\\nu}(r < 150 $kpc$)/f_{\\nu}$ relations.  The squares ($0.14 \\leq z \\leq 0.3$) and triangles ($z > 0.3$) represent the data of Reese et al. (2002).  The thick, thin lines are the $z$ = 0.2, 0.5 predictions, respectively.  For clarity, we plot the $z = 0.5$ lines for the self-similar and $K_0 = 100$ keV cm$^2$ models only.  The integrated Compton parameters (both data and models) have been arbitrarily rescaled for $z = 0.2$.} \\label{fig2} \\end{figure}  The relations that we have described above demonstrate that a high entropy floor ($K_0 \\gtrsim 300$ keV cm$^2$) is required to match the X-ray and SZ effect observations of massive clusters.  This is consistent with previous investigations of the $L_X - T_X$ relation (e.g., Tozzi \\& Norman 2001).  More work \\adjustfinalcols is required to determine the origin of the entropy floor. \\vskip-0.15in"
    },
    "0207/astro-ph0207237_arXiv.txt": {
        "abstract": "We study the angular power spectra of the polarized component of the Galactic synchrotron emission in the 28--deg$^2$ Test Region of the Southern Galactic Plane Survey at 1.4 GHz. These data were obtained by the Australia Telescope Compact Array and allow us to investigate angular power spectra down to arcminute scales. We find that, at this frequency, the polarization spectra for $E$-- and $B$--modes seem to be affected by Faraday rotation produced in compact foreground screens. A different behavior is shown by the angular spectrum of the polarized intensity $PI=\\sqrt{Q^2+U^2}$. This is well fitted by a power law ($C_{PI\\ell}\\propto\\ell^{-\\alpha_{PI}}$) with slope $\\sim1.7$, which agrees with higher frequency results and can probably be more confidently extrapolated to the cosmological window. ",
        "introduction": "In recent years the measurement of the Cosmic Microwave Background (CMB) polarization has become one of the major aims of a large number of planned experiments. Its detection is however a technological challenge: so far, we have only upper limits on the CMB polarization level [see Staggs et al. (1999) for a review and the recent measurements by PIQUE (Hedman et al. 2001) and POLAR (Keating et al. 2001) experiments]. Since the CMB polarization signal is expected to be less than 10$\\%$ of the temperature anisotropies, instrumental sensitivities of a few $\\mu$K or less are required. These will be probably reached by the forthcoming experiments: three space missions, MAP (Wright 1999), Planck (De Zotti et al. 1999) and SPOrt (which is completely devoted to the study of sky polarized emissions; Carretti et al. 2002) will measure the polarization on nearly the full sky, while several ground--based or balloon--borne experiments are planned to observe small sky areas with high spatial resolution [for instance, AMIBA (Kesteven et al. 2002), BOOMERanG 2K2 (Masi et al. 2002) and BaR--SPOrt (Zannoni et al. 2002); see De Zotti (2002) for a short review].  The possibility of extracting information on cosmological parameters from CMB experiments is strictly related to the computation of the angular power spectra (APS), which in the presence of Gaussian statistics give a complete statistical description of the CMB emission. Although for the temperature fluctuations the APS is easily defined through the ordinary spherical-harmonic expansion, for polarization we need two different components, the ``electric'' ($E$) and ``magnetic'' ($B$) modes, in order to describe the polarization intensity and orientation (Kamionkowski, Kosowsky \\& Stebbins 1997, Zaldarriaga \\& Seljak 1997). In section 2.2 we will discuss the definition of the polarization APS, with particular attention to the small scale limit.  The detection of the CMB polarization is constrained by the presence of foreground emissions. Different techniques have been worked out to separate the cosmological signal from the Galactic and extra--galactic emissions; all the methods exploit the differences in the frequency and spatial behaviors. For this reason, the analysis of the APS has become a common tool for studying the foregrounds: in fact, its knowledge allows us to estimate the foreground contamination to the CMB signal at different frequencies and angular scales (or equivalently, spherical--harmonic index $\\ell$). For the total intensity emission, information on the spatial properties of foregrounds is limited only to a small interval of frequencies and angular scales (see, e.g., Tegmark et al. 2000), and it is completely unsatisfactory between 20 and 90 GHz. For the polarization, the situation is even worse because of the lack of high resolution surveys covering large areas.  In this paper we consider the synchrotron emission, which is intrinsically highly polarized and is expected to be the dominant foreground at low frequencies. Several authors have estimated the APS of the polarized synchrotron using data obtained with different resolutions and from limited sky regions at various latitudes (for a review of the up--to--date surveys, see Tucci et al. 2000). On small sky patches, Fourier analysis has been applied to the Stokes parameters $Q$ and $U$ (from which APS of the $E$--, $B$--modes can be computed) and to the polarized intensity $PI=\\sqrt{Q^2+U^2}$ (Tucci et al. 2000; Tucci et al. 2001; Bruscoli et al. 2002). Scalar and spin--weighted harmonic expansions have been respectively used for $PI$ (Baccigalupi et al. 2001; Giardino et al. 2002) and for $Q$ and $U$ (Giardino et al. 2002). These computations are not equivalent in an important respect. Fourier analysis, being adequate for small sky patches, necessarily provides {\\it local} effective spectra which for highly non-Gaussian fields (such as the Galactic synchrotron distribution) may be widely different from the global angular spectra. In practice such local spectra may be more important than the global ones for the separation of Galactic foregrounds from CMB. Further, we note that the sum $C_{P\\ell}=C_{E\\ell}+C_{B\\ell}$, which is the quantity usually considered in CMB analyses, and $C_{PI\\ell}$ provide different information. This important point has been noticed in Tucci et al. (2002), and will be discussed in Section 2.2.  Here we observe that it helps to explain some discrepancies appearing in the literature. Estimates of the synchrotron polarization spectra were performed from the Parkes surveys of the Southern (Duncan et al. 1995, 1997, hereafter D97) and Northern (Duncan et al. 1999, hereafter D99) Galactic Plane. For both surveys, sampling more than half of the Galactic Equator, the APS are nearly independent of longitude, and could be modeled by power laws with slopes $\\alpha_E\\simeq\\alpha_B=1.4\\div1.5$ and $\\alpha_{PI}=1.6\\div1.8$ in the $\\ell$--range $100\\div800$ [Tucci et al. 2000, 2001, Baccigalupi et al. 2001, Bruscoli et al. 2002, Giardino et al. 2002 (although the latter obtained $\\alpha_{PI}=2.37\\pm0.21$ in the $\\ell$--range $40\\div250$ from D97 data)]. Out of the Galactic Plane five patches are available at intermediate latitudes, $5^{\\circ}\\le|b|\\le20^{\\circ}$, from the survey by Uyaniker et al. (1999). The APS vary significantly there, with  slopes ranging from 1 to $\\sim2.5$ (Baccigalupi et al. 2001, Bruscoli et al. 2002). At latitudes far from the Galactic plane the only available information comes from the low-resolution survey of Brouw \\& Spoelstra (1976), covering about 40\\% of the sky at five frequencies in the range $408\\div1411$ MHz. From this survey, Bruscoli et al. (2002) found values of $\\alpha_{E,B,\\,PI}$ between $1\\div2$ at scales $\\ell<100$, in agreement with the results from higher resolution data. Using the same survey, Baccigalupi et al. (2001) studied the $PI$ field and found steeper spectra, with $\\alpha_{PI} \\simeq 3$.  The present paper extends the analysis of the APS of the polarized Galactic synchrotron to arcminute scales ($\\ell=10^3\\div10^4$), i.e., to angular scales smaller than in previous works by nearly one order of magnitude. The study of the synchrotron contribution on these scales is relevant for CMB observations. In fact, the angular scales $300\\la\\ell<2000$ are expected to be those where CMB should exhibit the highest level of polarized signal. Moreover, at $\\ell>3000$ non-linear effects on CMB become important, producing polarized signal stronger than the primary spectrum. These include the Vishniac, patchy reionization and kinetic Sunyaev-Zeldovich effects (Hu 2000, Liu et al. 2001).  We make use of high--resolution polarization data taken from a test region for the Southern Galactic Plane Survey (McClure--Griffiths et al. 2001, Gaensler et al. 2001, hereafter G01) consisting of 1.4--GHz observations carried out with the Australia Telescope Compact Array (ATCA). We find that the $E$ and $B$ spectra can be well approximated by a power law at $600\\la\\ell\\la6000$, with a steep slope $\\alpha_{E,B}\\simeq 2.7\\div 2.9$. Moreover, we compute the spectrum of the polarized intensity, $PI$ and find that it is remarkably different from the above spectra, following a power law with $\\alpha_{PI}\\simeq1.7$ on the whole $\\ell$--range. We compare the ATCA APS with the APS computed in the corresponding patch from the 2.4 GHz Parkes survey. We find a noticeable agreement for $C_{PI,\\ell}$, but not for $C_{E,B,\\ell}$. The different slopes found in the APS are interpreted in section 2.4 as due to effects of Faraday rotation produced by foreground screens. No evidence is found for a contribution of extragalactic point sources.  The behaviors of synchrotron APS on arcminute scales at GHz frequencies may be interesting for information on Galactic structure that spectra contain. They can tell us, in fact, about both the magnetohydrodynamic turbulence in the emitting region and the electron density fluctuations in the intervening medium (shown by Faraday rotation). In this connection, the comparison between 1.4\\,GHz and 2.4\\,GHz data may be useful to separate the transverse structure of the magnetic field in the emitting region and the longitudinal field in the foreground screens. This point is open to future studies. ",
        "conclusions": "In this paper, for the first time, we extend the study of the angular power spectrum for the polarized component of the Galactic synchrotron emission to arcminute scales, i.e. up to $\\ell\\sim10^4$. To reach such scales we needed high--resolution data, which were provided by the ATCA observations of a small patch of the Galactic Plane at 1.4 GHz.  In the paper we compute the polarization spectra for ``electric'' and ``magnetic'' modes, plus the spectrum for the polarized intensity. We find that, in the range $600\\le\\ell\\le6000$, both $C_{E\\ell}$ and $C_{B\\ell}$ can be well approximated by power laws with slopes $\\alpha_E\\simeq\\alpha_B\\sim 2.7\\div 2.9$. Such spectra are significantly steeper than those arising at $\\ell\\le800$ from low--resolution data. Moreover, their amplitude, if compared to the spectra obtained by the Parkes telescope in the same sky area, turn out to be higher by nearly one order of magnitude at angular scales between $20^{\\prime}$ and 10$^{\\prime}$. These peculiar behaviors are well interpreted as due to the small--scale modulation of a relatively uniform polarized background by Faraday rotation along the line of sight. On the contrary, we believe that our estimates of $C_{PI\\ell}$, whose slope ($\\alpha_{PI}\\sim1.7$) is in agreement with D97 data at 2.4 GHz, are not affected by Faraday effects and fairly describe the intrinsic spatial distribution of the polarized emission.  An interesting point, which arises from our analysis, regards the distinctive meaning of $C_{PI\\ell}$ with respect to $C_{E,B\\ell}$. As we have discussed in section 2.2, $PI$ is a scalar quantity and refers only to the intensity of the polarization without any information on its direction. We then expect that the APS for $E$-- and $B$--modes, that provide a complete description of the polarization field, do not have the same shape as the $PI$ spectrum. The differences should be greater when the direction of polarization changes very rapidly. Deviations between $C_{E\\ell}$ and $C_{PI\\ell}$ are found in the CMB (see the results of the simulations in Fig. \\ref{f2}): these are not unexpected, because of the geometry of the polarization angle in the $E$--mode spots. In the case of synchrotron emission, the polarization direction for the diffuse component is quite smooth on large scales following the Galactic magnetic field. From low--resolution surveys the estimates of $C_{E,B\\ell}$ and $C_{PI\\ell}$ give only moderate differences in the spectral shape (see Bruscoli et al. 2002). However, when small scales are considered, fluctuations in magnetic fields, discrete sources and also Faraday effects contribute to amplify the variations in the polarization spectra, as the present results highlight.  The extrapolation of our results for $C_{E\\ell }$ and $C_{B\\ell }$ to higher frequencies should not be regarded as a trivial matter, since Faraday effects are substantial at 1.4 GHz while they become negligible at a few tens of GHz. The electric and magnetic spectra are strongly affected by the Faraday rotation along the line of sight, showing a steep slope ($\\alpha_{E,B}\\sim2.8$). The polarized intensity spectrum, instead, can be more reliably extrapolated to the ``cosmological'' frequencies, because, as discussed in section 2.4, it is not affected by Faraday rotation, except in severe cases when significant depolarization is occurring. The power index $\\alpha_{PI}$ is less than 2 independent of the region analysed, with a value of $1.66\\pm0.05$ in the $4^{\\circ}\\times4^{\\circ}$ box.  We have seen that analyses on the synchrotron polarization spectrum in the literature indicate a moderate slope ($\\alpha_{X}\\la2$ with $X=E,B,\\,PI$) on angular scales $\\ell<10^3$. Now ruling out a significant contribution from point sources, we confirm this result also at $\\ell\\la10^4$ for $C_{PI\\ell}$, and we show that synchrotron emission is rather rich in small scale structures. Hence, contrary to what is usually assumed, it might be a relevant contaminant in CMB polarization measurements at very small scales. It remains to be seen if the APS as deduced by ATCA observations of the Galactic plane is a common feature in regions at high galactic latitudes. In general we expect the polarized synchrotron emission to be fainter in regions far out of the Galactic plane, except in very bright areas. In one such region Bruscoli et al. (2002) estimate $C_{X\\ell}$ for $\\ell<100$, finding a spectrum behavior consistent with those of the Galactic plane ($1<\\alpha_{X}<2$). Even if very bright regions are not typical at high latitudes, these play an important role in the process of foreground subtraction in CMB experiments; this is why we have to put great care into the study of the APS of these regions."
    },
    "0207/astro-ph0207001_arXiv.txt": {
        "abstract": "We present the results of deep X-ray and $\\gamma$-ray observations of the Geminga pulsar obtained in the final years of the \\asca\\ and {\\it CGRO} missions, and an upper limit from {\\it RXTE\\/}.  A phase-connected ephemeris from the $\\gamma$-rays is derived that spans the years 1973--2000, after allowing for a minor glitch in frequency of $\\Delta f/f = 6.2 \\times 10^{-10}$ in late 1996.  \\asca\\ observations of the hard X-ray pulse profile in 1994 and 1999 confirm this glitch.  An improved characterization of the hard X-ray pulse profile and spectrum from the long \\asca\\ observation of 1999 confirms that there is a non-thermal X-ray component that is distinct from the $\\gamma$-ray spectrum as measured by EGRET.  It can be parameterized as a power-law of photon index $\\Gamma = 1.72 \\pm 0.10$ with a flux of $2.62 \\times 10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ in the $0.7-5$ keV band and pulsed fraction $0.54 \\pm 0.05$, similar to, but more precise than values measured previously.  An extrapolation of this spectrum into the energy band observed by the {\\it RXTE\\/} PCA is consistent with the non-detection of pulsed emission from Geminga with that instrument.  These results are interpreted in the context of outer-gap models, and motivations for future X-ray observations of Geminga are given. ",
        "introduction": "Discovered in 1972 by the SAS-2 satellite \\citep{fi75,th77}, Geminga is the second brightest $\\gamma$-ray source in the sky above 100 MeV \\citep{sw81}. It was known only as a $\\gamma$-ray source until a promising candidate was detected in X-rays by the Einstein Observatory \\citep{bi83}, and associated with an optical counterpart \\citep{bi87,ht88,bi88}.  Subsequently, Geminga was found to be a rotation-powered pulsar with a period of 237~ms in X-rays by \\ro\\ \\citep{hh92}, and in $\\gamma$-rays by the Energetic Gamma Ray Experiment Telescope (EGRET) on the {\\it Compton Gamma-ray Observatory} \\citep{be92}. Prior to the discovery of the 237~ms spin period of Geminga, claims had been made for various periods in the range 59--60~s in $\\gamma$-rays and in X-rays \\citep{th77,ma77,zm83,bi84,zy88,ka85}, but no such detections have been made in high quality X-ray and $\\gamma$-ray observations during the past decade.  The optical spectrum of Geminga is predominantly non-thermal, with possible ion cyclotron features \\citep{mhs98,mcb98}.  \\cite{sh98} reported optical modulation from Geminga that resembles its $\\gamma$-ray light curve.  Geminga is unusual as a rotation-powered pulsar because it is not a strong radio source. In 1997, three groups \\citep{mm97,kl97,sp97} claimed detection of pulsed radio emission at 102~MHz, but observations at other radio frequencies have so far been negative \\citep{rdi98,mcl99,bfb99,kl99}.  A phase-connected ephemeris covering the first 27 years of $\\gamma$-ray observations of Geminga was presented and updated by Mattox, Halpern, \\& Caraveo (1998, 2000).  In this paper we present the results of a long observation with the {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca ), which allows us to better constrain the hard X-ray spectrum of Geminga and perform pulse-phase spectroscopy.  X-ray pulse times of arrival are compared with the latest ephemeris from EGRET.  Additional constraints on the hard X-ray emission are derived from an observation by the {\\it Rossi X-ray Timing Explorer} Proportional Counter Array ({\\it RXTE\\/} PCA). ",
        "conclusions": ""
    },
    "0207/astro-ph0207284_arXiv.txt": {
        "abstract": "{ We present 86 GHz ($v = 1, J = 2 \\rightarrow 1$) SiO maser line observations with the IRAM 30-m telescope of a  sample of 441 late-type stars in the Inner Galaxy ($ -4 \\degr < l < +30 \\degr$). These stars were selected on basis of their infrared magnitudes and colours from the ISOGAL and MSX catalogues. SiO maser emission was detected in 271 sources, and their line-of-sight velocities indicate that the stars are located in the Inner Galaxy. These new detections double the number of line-of-sight velocities  available from previous SiO and OH maser observations in the area covered by our survey and are, together with other samples of e.g.\\ OH/IR stars, useful for kinematic studies of the central parts of the Galaxy. ",
        "introduction": "There has been a growing interest in studies characterizing the kinematics and the spatial distribution of stars in the Inner Galaxy ($ 30 \\degr < l < -30 \\degr$).  Many  recent studies attempt to determine the parameters that describe the dynamics and structure of the Inner Galaxy, i.e.\\ its central bar and/or its bulge tri-axial mass distribution.  One approach is to map the spatial density of a stellar population. This has been done, e.g., for  stars detected by IRAS toward the Galactic bulge \\citep{nakada91,weinberg92}, bulge Mira variables \\citep{whitelock92}, bulge red clump stars \\citep{stanek94} and giant stars seen in  fields at symmetric longitudes with respect to the Galactic centre \\citep{unavane98}.  Optical studies of the Inner Galaxy are much hindered by the high interstellar extinction, which can exceed $A_V \\approx 30$ \\citep[e.g.][]{schultheis99}, and thus are limited to small optical windows \\citep{holtzman98,zhao94}. At infrared and radio wavelengths however, interstellar extinction is much less severe, or even absent. Extensive infrared  point source catalogues have recently become available from the ground based DENIS \\citep{epchtein94} and 2MASS \\citep{beichman98} near-infrared (nIR) surveys,  the mid-infrared (mIR)  ISO satellite survey  \\citep[ISOGAL:][]{omont99,omont02}, and the Midcourse Space Experiment \\citep[MSX:][]{price97,egan99}. These data have given new insights into the  {\\em spatial} stellar density distribution in the Inner Galaxy. To interpret the information given by the recent observations, detailed models all include some kind of tri-axiality: a tri-axial Galactic bulge or  bar \\citep[e.g.][]{debattista02,ortwin02,lopezcorreidora01,lopezcorreidora01b, alard01b}. However, the bar characteristics such as length,  pattern speed, and position angle, are still poorly constrained.  Spatial density studies often  neglect an important measurable dimension of phase space: the stellar line-of-sight velocity. In contrast to the large number of data points in the spatial domain of phase-space, the available data on the line-of-sight velocities of the stars is  sparse because it is still   difficult to  measure velocities from optical or infrared studies. Asymptotic Giant Branch (AGB) stars with large mass-loss are a valuable exception, since their envelopes often harbour masers which are strong enough to be detected throughout the Galaxy and  thereby reveal the line-of-sight velocity of the star to within a few \\kms; frequently detected maser lines are from OH at 1.6 GHz, H$_2$O at 22 GHz, and SiO at 43 GHz and 86 GHz \\citep[for a review see][]{habing96}. Previous SiO and OH maser surveys in the Galaxy have demonstrated that locating the circumstellar masers is an effective way to measure line-of-sight velocities of the AGB stars \\citep[e.g.][]{baud79,lindqvist92,blommaert94,sevenster01, sevenster97a,sevenster97b,sjouwerman98a,izumiura99,deguchi00a, deguchi00b}.  Until recently, only a few hundred stellar line-of-sight velocities were known toward the inner regions of the Milky Way  ($ 30 \\degr < l < -30 \\degr$ and $|b| < 1$), mainly from OH/IR stars, AGB stars with OH maser emission in the 1612 MHz line, mostly undetected at visual wavelengths. This number is too small to allow for a good quantitative multicomponent analysis of the Galactic structure and dynamics \\citep{vauterin98}. Obtaining more line-of-sight velocities therefore remains an issue of prime importance. However, masers are rare among stars, because sustaining a maser requires a special physical environment. Most of the mid-infrared brightest OH/IR stars  close to the Galactic plane were probably already detected in the blind OH surveys or in the targeted OH or 43 GHz SiO maser observations of colour-selected sources from the IRAS survey \\citep[e.g.][]{vanderveen88}. H$_2$O surveys  \\citep[e.g.][]{levine95} are  probably incomplete because the H$_2$O masers are strongly variable.  SiO maser emission is detected  from several transitions towards oxygen-rich AGB stars and red supergiants. On the basis of  the shape and the amplitude of  their visual light curve AGB stars  have been classified as  semi-regular (SR) stars  and Mira stars. Variable AGB stars may also  be classified as long period variable (LPV) stars,  when their periods are longer than 100 days \\citep{habing96}. Almost all OH/IR stars are  variable and have periods longer than 500 days. In the IRAS color-color diagram the oxygen-rich AGB stars are distributed on a well-defined sequence of increasing shell opacity and stellar mass-loss rate \\citep[e.g.][]{vanderveen88,olnon84}, which goes from Miras with the bluest colors and the 9.7 $\\mu$m silicate feature in emission, to OH/IR stars with the reddest colors and the 9.7 $\\mu$m silicate feature in absorption. The relative strengths of different SiO maser lines are observed to vary with AGB type \\citep{nyman86,nyman93,bujarrabal96}, indicating that the SiO maser properties depend on the stellar mass loss rate and  on the stellar variability. The ratio of the SiO maser intensities of 43 over 86 GHz is found to be much lower in Mira stars and in supergiants than in OH/IR stars. This implies that the 86 GHz ($v = 1$) SiO maser transition is a good tool to measure stellar line-of-sight velocities  of Mira-like stars. Another advantage is that Mira stars are far more numerous than  OH/IR stars. However, these conclusions are based on small number statistics, and have neglected effects of variability.  To significantly enlarge the number of known stellar line-of-sight velocities we have conducted a targeted survey for the 86 GHz SiO ($v = 1, J = 2 \\rightarrow 1$) maser line toward  an infrared selected sample of late-type stars. Here we describe the selection of sources and the observational results. A detailed discussion of the kinematic and physical properties of the detected stars will be addressed in a forthcoming paper \\citep[][ in preparation]{messineo02b}. All velocities in this paper refer to line-of-sight velocities, measured with respect to the Local Standard of Rest (LSR). ",
        "conclusions": "We have observed 441 colour-selected ISOGAL and MSX sources in the Inner Galaxy ($30^\\circ < l < -4^\\circ$ and $|b|$ mostly $<1$), in the SiO ($v = 1, J = 2 \\rightarrow 1$) maser transition and  detected 271  lines. We thereby obtained 255 new line-of-sight velocities which doubles the number of maser line-of-sight velocities  known in  the region we surveyed. To search for 86 GHz ($v = 1$) SiO maser lines in colour-selected mIR sources has proven to be an efficient way to  obtain stellar radial velocities in the Inner Galaxy. In the central 2 degrees we notice some confusion with interstellar \\hcn\\ emission, but usually the interstellar \\hcn\\  and the stellar SiO line can be separated  well  by using their radial velocities and line widths. The SiO maser emission was detected towards  61 \\%  of our sources, objects which lie in a transition region of the IR-colour space between Mira and OH/IR stars. The SiO maser detectability decreases with decreasing mIR flux density. We observed 15 sources from the sample of LPV stars by \\citet{glass01} and found  86 GHz SiO maser emission in 11 of them (73 \\%), while only 23 \\% of the LPV stars which follow our selection criteria show OH maser emission. Therefore 86 GHz SiO maser emission is more frequent than OH maser emission. In a later study we will use our new catalogue of stellar line-of-sight velocities for a quantitative analysis of stellar kinematics and SiO maser properties  in the Inner Galaxy."
    },
    "0207/astro-ph0207551_arXiv.txt": {
        "abstract": "{\\small Radio emission from X-ray binary systems (XRBs) has developed in recent years from being peculiar phenomenon to being recognised as an ubiquitous property of several classes of XRBs. In this scenario the synchrotron emission is interpreted as the radiative signature of jet-like outflows, some or all of which may possess relativistic bulk motion. We have analysed a collection of quasi-simultaneous radio/X-ray observations of Black Holes in the Low/Hard X-ray state, finding evidence of a clear correlation between their fluxes over many orders of magnitude in luminosity. Given that the correlation extends down to GX 339-4 and V404 Cyg in quiescence, we can confidently assert that even at accretion rates as low as $\\sim 10^{-5}$ $\\dot{m}_{Edd}$ a powerful jet is being formed. The normalisation of the correlation is very similar across a sample of nine sources, implying that it is nearly independent of jet inclination angle. Remarkably, V 404 Cyg is the second source (after GX 339-4) to show the correlation $S_{radio}\\propto S_{X}^{+0.7}$ from quiescent level up to close to the High/Soft state transition. Moreover, assuming the same physics and accretion:outflow coupling for all of these systems, the simplest interpretation for the observed scenario is that outflows in Low/Hard state do not have large bulk Lorentz factors.} ",
        "introduction": "Radio emission is often observed from X-ray binaries, particularly transient systems and especially Black Hole (BH) candidates (for detailed reviews see \\eg \\cite{hh}, \\cite{mr}, \\cite{rob1}, \\cite{rob2}). The Low/Hard state is one of the five `canonical' X-ray states observed from BH X-ray binaries in our Galaxy. It is charaterized by a hard (spectral index $\\sim$ 1.5) X-ray spectrum, associated with a steady, self-absorbed jet which emits synchrotron radiation; the Off state may simply be the Low state `turned down' to lower accretion rates. X-ray spectra of High/Soft state BHs are dominated by thermal radiation, while the radio emission drops below detectable levels, probably corresponding to the physical disappearance of the jet. For the so called Intermediate and Very High states the connection between X-ray and radio properties is not yet completely established. Transitions between states are often associated with multiple ejections of synchrotron emitting material, possibly with high Lorentz factors.\\\\ There is an extremely strong correlation between radio and X-ray emission: the former has been directly observed to arise in outflows and to produce synchtrotron emission from a population of high energy electrons. The latter has been inferred to arise via Comptonisation by a thermal electron distribution. All the evidence points to the corona in these systems being physically related to the presence of a jet: by far the simplest interpretation therefore is that the Comptonising region is just the base of the relativistic outflow. ",
        "conclusions": "The results of our analysis can be summarized as follows: \\begin{itemize} \\item In Low/Hard state BHs the observed radio and X-ray fluxes are correlated over more than three orders of magnitude in accretion rate, implying a strong jet/corona coupling; no lower limit to the relation has been found. \\item We can confidently assert that, even at accretion rates as low as $\\sim$ $10^{-5} \\dot m_{Edd}$ a powerful jet is being formed. \\item V 404 Cyg is the second source to display $S_{radio} \\propto S_{X}^{+0.7}$, from quiescence up to Soft state transition. A physical explanation for this relation is proposed by Markoff \\etal (\\cite{sera}). \\item Above $\\sim 10^{-2} \\dot m_{Edd}$ the jet disappears within a factor of a few, probably in all sources (observed in three sources and no exceptions) \\item Comparison of different sources may indicate that jets in Low/Hard have a low velocity compared to those in transient outbursts. Of course better distance and inclination determinations are strongly required to probe this conjecture.\\\\ Finally, we would like to mention that the possible detection of the Soft X-ray Transient A0620-00 at the predicted -- extrapolating the relation we found -- radio level would demonstrate that jet production is an ubiquitous mechanism between $10^{-7}$ and $10^{-2} \\dot m_{Edd}$. \\end{itemize}"
    },
    "0207/astro-ph0207621_arXiv.txt": {
        "abstract": "name{Samenvatting}% \\def\\bibname{Bibliografie}\\def\\chaptername{Hoofdstuk}% \\def\\appendixname{Bijlage}\\def\\contentsname{Inhoudsopgave}% \\def\\listfigurename{Lijst van figuren}\\def\\listtablename{Lijst van tabellen}% \\def\\indexname{Index}\\def\\figurename{Figuur}\\def\\tablename{Tabel}% \\def\\partname{Deel}\\def\\enclname{Bijlage(n)}\\def\\ccname{Ter attentie van}% \\def\\headtoname{Aan}\\def\\headpagename{Pagina}% \\def\\today{\\number\\day\\space\\ifcase\\month\\or januari\\or februari\\or maart\\or% april\\or mei\\or juni\\or juli\\or augustus\\or september\\or oktober\\or% november\\or december\\fi \\space\\number\\year}% \\typeout{ >>>>> use hlatex209 for Dutch hyphenation <<<<< }} \\hyphenation{Schrij-ver Krij-ger Kuij-pers Bal-le-gooij-en}   \\def\\warningoverprint #1{\\special{!userdict begin /bop-hook{gsave 100 600 translate -45 rotate /Times-Roman findfont 100 scalefont setfont 0 0 moveto 0.9 setgray (#1) show grestore}def end}}  \\newcounter{onefig} \\newcounter{fignumber} \\newcount\\nocaptions \\newcount\\nofigures \\newcount\\figwidth \\newcount\\viewgraphs \\def\\figsfile{} \\def\\figspath{} \\def\\paper{} \\def\\format{} \\def\\figlabel{} \\long\\def\\nextfig#1{\\setcounter{figure}{\\value{fignumber}} \\addtocounter{fignumber}{1} \\ifnum \\viewgraphs=1 \\newpage \\pagestyle{empty} \\fi \\ifnum\\value{onefig}=0 #1 \\fi \\ifnum\\value{onefig}=\\value{fignumber} #1 \\fi} \\def\\figwidths#1#2{\\ifnum \\nocaptions=1 #2mm \\else #1mm \\fi} \\def\\paper#1{}  % \\long\\def\\plotfig#1#2{\\ifnum \\nofigures=1 \\else #2 \\fi} \\long\\def\\captiontext#1{\\ifnum \\nofigures=1 \\raggedright \\fi \\ifnum \\nocaptions=1 \\paper \\ifnum \\viewgraphs=0 \\newline  \\mbox{}\\hrulefill\\mbox{} \\newline \\newline label:~\\{\\figlabel\\} \\fi \\else \\ifnum \\nofigures=0 \\fi #1 \\fi}  \\newcount\\panelwidth \\newcount\\panelheight \\newcount\\bxmin \\newcount\\bymin \\newcount\\bxmax \\newcount\\bymax \\newcount\\tbxmin \\newcount\\tbymin \\newcount\\tpanelwidth \\newcount\\tpanelheight \\newcount\\tpdif \\panelwidth=70 \\panelheight=70  % \\def\\panelsize #1,#2;{\\panelwidth=#1 \\panelheight=#2} \\def\\setbb #1,#2;#3,#4;#5,#6;{% \\tbxmin=#1 \\tbymin=#2    % \\bxmin=#3 \\bymin=#4      % \\bxmax=#5 \\bymax=#6}     % \\def\\barepanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,clip=,% width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}} \\def\\labelxpanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelxypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}}  \\newcommand{\\panellabel}[3]{\\makebox[0pt][r]{\\raisebox{#2} #3 \\hspace{#1}}}  \\def\\topfraction{0.99}                      % \\def\\dbltopfraction{0.99} \\def\\textfraction{0.01} \\def\\floatpagefraction{0.99}  \\def\\CC{\\par \\vspace*{-2ex} \\footnotesize \\baselineskip=8pt \\begin{verbatim}} \\def\\EE{\\par \\vspace*{-2ex} \\normalsize}  \\long\\def\\startignore #1\\stopignore{}   %   \\newenvironment{task}{% \\begin{list}{$\\bullet$}{\\itemsep=1ex \\parindent=3ex \\parskip=1ex \\topsep=0ex}} {\\end{list}}  \\def\\koprr{\\list{}{ \\itemindent=0pt \\listparindent=0pt \\itemsep=0pt \\leftmargin=5ex}\\item[]}                     % \\let\\endkoprr=\\endlist  \\def\\kolom{\\list{}{\\rightmargin=5cm      % \\topsep=0ex \\itemsep=0cm \\parsep=0ex \\parskip=0ex \\partopsep=0ex \\listparindent=3ex \\itemindent=0pt \\leftmargin=0pt} \\item[] \\vspace*{-0ex}}                        % \\let\\endkolom=\\endlist  \\renewcommand{\\labelitemi}{{\\bf --}}                % \\def\\setlistparams{ \\topsep=0.7ex                 % \\itemsep=0.7ex                % \\leftmargini=3ex}             % \\setlistparams                  %  \\newcounter{alistindex}  \\newcounter{alistvalue}     % \\newenvironment{alist}{% \\begin{list}{\\alph{alistindex})}{\\usecounter{alistindex}} \\itemsep=1.5ex \\parskip=1ex}{\\end{list}}  \\def\\spacing #1{\\small \\renewcommand{\\baselinestretch}{#1} \\normalsize \\setlistparams}             %  \\newenvironment{itemrr}{\\vspace*{-\\parskip} \\leftmargini=3ex \\topsep=0.6ex \\begin{itemize} \\itemsep=0.5ex \\parindent=0ex \\parsep=0ex}{ \\end{itemize} \\vspace*{-\\parskip}}  \\newenvironment{conditem}{ \\leftmargini=6ex \\topsep=0.2ex \\begin{itemize} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\newenvironment{enumerr}{ \\leftmargini=3ex \\topsep=1.5ex \\begin{enumerate} \\itemsep=0.7ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newenvironment{condenum}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{enumerate} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newcounter{romenumnr} \\newenvironment{romanenum}{ \\leftmargini=5ex \\topsep=1.5ex \\begin{list}{(\\roman{romenumnr}) --}{\\usecounter{romenumnr} \\itemsep=1ex \\listparindent=3ex \\parsep=0ex}}{ \\end{list}}  \\newcommand{\\posterpage}[2]{{\\Large \\sf \\newpage \\setlength{\\minipagewidth}{\\textwidth} \\addtolength{\\minipagewidth}{-2cm} \\fboxrule=0.3ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[\\textwidth]{ \\begin{minipage}[t]{\\minipagewidth} \\parskip=3ex \\vspace*{1ex} \\begin{center} \\LARGE \\bf #1 \\end{center} \\vspace{2ex} #2 \\vspace*{3ex} \\mbox{ } \\end{minipage}} \\hspace*{\\textwidth} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.2ex} \\framebox[\\textwidth]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{\\textwidth}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}}  \\newcommand{\\vwgraph}[3]{\\newpage \\begin{center} \\shadowframe{#1}{ \\begin{center} \\vspace*{2ex} {\\Large \\bf #2}\\\\[4ex] \\end{center} {\\Large \\sf #3}}   % \\end{center}} \\newenvironment{itemvg}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{itemize} \\vspace{-1ex} \\itemsep=0.5ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\long\\def\\lectvw#1#2{\\newpage \\vspace*{-15mm} \\begin{center} \\begin{center} {\\LARGE \\sc \\underline{#1}}\\\\[4ex] \\end{center} {\\Large \\sf #2}   % \\end{center}}  \\newlength{\\minipagewidth} \\newcommand{\\shadowframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{1.0ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{4.5ex} \\rule{#1}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newsavebox{\\boxcontent} \\newcommand{\\ovalhead}[1]{ \\unitlength=1cm \\sbox{\\boxcontent}{\\mbox{~~{#1}~~}} \\begin{center} \\ifdim\\wd\\boxcontent>6ex \\ifdim\\wd\\boxcontent<8cm \\begin{picture}(8,3) \\thicklines \\put(4.0,0.8){\\oval(8,1.6)} \\put(0.0,0.7){\\parbox{8cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\ifdim\\wd\\boxcontent<12cm \\begin{picture}(12,3) \\thicklines \\put(6.0,0.8){\\oval(12,1.6)} \\put(0.0,0.7){\\parbox{12cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\begin{picture}(16,3) \\thicklines \\put(8.0,0.8){\\oval(16,1.6)} \\put(0.0,0.7){\\parbox{16cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\fi \\fi \\fi \\end{center}}  \\newcommand{\\colloqframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\parskip=3ex \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{#1}{2ex} \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newcommand{\\soloval}[1]{\\fboxrule=0.5mm \\unitlength=1cm \\hspace*{5mm} \\sol{20mm}\\\\[-20mm] \\begin{picture}(15,3) \\thicklines \\put(7.5,2){\\oval(15.60,3.60)} \\put(7.5,2){\\oval(15.55,3.55)} \\put(7.5,2){\\oval(15.50,3.50)} \\put(7.5,2){\\oval(15.45,3.45)} \\put(7.5,2){\\oval(15.40,3.40)} \\put(7.5,2){\\oval(15.35,3.35)} \\put(7.5,2){\\oval(15.00,3.00)} \\put(2.0,1.9){\\parbox{13cm}{#1}} \\end{picture}}  \\setcounter{secnumdepth}{3} \\setcounter{tocdepth}{3}  \\def\\chapterrr#1{\\chapter{#1} \\label{chap:#1} \\vskip-\\parskip} \\def{ The fine calibration of the ISO-SWS detectors (Infrared Space Observatory - Short Wavelength Spectrometer) has proven to be a delicate problem. We therefore present a detailed spectroscopic study in the 2.38 -- 12\\,\\mic\\ wavelength range of a sample of 16 A0 -- M2 stars used for the calibration of ISO-SWS. By investigating the discrepancies between the ISO-SWS data of these sources, the theoretical predictions of their spectra, the high-resolution FTS-KP (Kitt Peak) spectrum of $\\alpha$ Boo and the solar FTS-ATMOS (Atmospheric Trace Molecule Spectroscopy) spectrum, both {\\it calibration} problems and problems in {\\it computing the theoretical models and the synthetic spectra} are revealed. The underlying reasons for these problems are sought for and the impact on the further calibration of ISO-SWS and on the theoretical modelling is discussed extensively. ",
        "introduction": "For the astronomical community analysing ISO-SWS data \\citep[Infrared Space Observatory, Short-Wavelength Spectrometer,][]{deGraauw1996A&A...315L..49D}, a first point to assess when judging and qualifying their data concerns the flux calibration accuracy. Since the calibration process is not straightforward, knowledge on the {\\it{full}} calibration process and on the still remaining calibration problems is crucial when processing the data.  One way to detect calibration problems is by comparing observed data with theoretical predictions of a whole sample of standard calibration sources.  But, as explained in \\citet{Decin2000A&A...364..137D} (hereafter referred to as Paper~I) a full exploitation of the ISO-SWS data may only result from an iterative process in which both new theoretical developments on the computation of stellar spectra  --- based on the MARCS and Turbospectrum code \\citep{Gustafsson1975A&A....42..407G, Plez1992A&A...256..551P, Plez1993ApJ...418..812P}, version May 1998 --- and more accurate instrumental calibration are involved.   Precisely because this research entails an iterative process, one has to be extremely careful not to confuse technical detector problems with astrophysical issues. Therefore, the analysis in its entirety encloses several steps. Some steps have already been demonstrated in the case of $\\alpha$ Tau in \\citetalias{Decin2000A&A...364..137D}. They will be summarised in Sect.\\ \\ref{summary}. Other points will be introduced in Sect.\\ \\ref{summary} and will be elaborated on in the first sections of this article (Sect.\\ \\ref{sample} -- \\ref{highres}). Having described the method of analysis, the general discrepancies between observed and synthetic spectra are subjected to a careful scrutiny in order to elucidate their origin.  At this point, a distinction can be made between discrepancies typically for {\\it{warm}} stars and those typical for {\\it{cool}} stars. For this research, {\\it{warm}} stars are defined as being hotter than the Sun (T$_{\\mathrm{eff},\\odot} = 5770$~K) and their infrared spectra are mainly dominated by atomic lines, while molecular lines are characteristic of cool star spectra. A description on the general trends in the discrepancies for {\\it{warm}} and {\\it{cool}} stars will be made in this paper, while each star of the sample will be discussed individually in two forthcoming papers in which also an overview of other published stellar parameters will be given.  As stated in \\citetalias{Decin2000A&A...364..137D}, the detailed spectroscopic analysis of the ISO-SWS data has till now been restricted to the wavelength region from 2.38 to 12\\,\\mic. So, if not specified, the wavelength range under research is limited to band 1 (2.38 -- 4.08\\,\\mic) and band 2 (4.08 -- 12.00\\,\\mic). Band 3 (12.00 -- 29.00\\,\\mic) will be elaborated on by Van Malderen (Van Malderen et al., 2001, in prep.).  This paper is organised as follows: in Sect.\\ \\ref{summary} the general method of analysis is summarised. The sample of ISO-SWS observations is described in Sect.\\ \\ref{sample}, while the data reduction procedure is discussed in Sect.\\ \\ref{datareduction}. The observations of two independent instruments are introduced in Sect.\\ \\ref{highres}. In Sect.\\ \\ref{results}, the results are elaborated on. In the last section, Sect.\\ \\ref{impact}, the impact on the calibration of ISO-SWS and on the theoretical modelling is given.   The appendix of this article is published electronically. Most of the grey-scale plots in the article are printed in colour in the appendix, in order to better distinguish the different spectra. ",
        "conclusions": ""
    },
    "0207/astro-ph0207403_arXiv.txt": {
        "abstract": "Oscillations detected on the solar surface provide a unique possibility for investigations of the interior properties of a star. Through major observational efforts, including extensive observations from space, as well as development of sophisticated tools for the analysis and interpretation of the data, we have been able to infer the large-scale structure and rotation of the solar interior with substantial accuracy, and we are beginning to get information about the complex subsurface structure and dynamics of sunspot regions, which dominate the magnetic activity in the solar atmosphere and beyond. The results provide a detailed test of the modeling of stellar structure and evolution, and hence of the physical properties of matter assumed in the models. In this way the basis for using stellar modeling in other branches of science is very substantially strengthened; an important example is the use of observations of solar neutrinos to constrain the properties of the neutrino. ",
        "introduction": "\\plabel{sec:introduc}  By the standards of astrophysics, stars are relatively well understood. Modelling of stellar evolution has explained, or at least accounted for, many of the observed properties of stars. Stellar models are computed on the basis of the assumed physical conditions in stellar interiors, including the thermodynamical properties of stellar matter, the interaction between matter and radiation and the nuclear reactions that power the stars. By following the changes in structure as the stars evolve through the fusion of lighter elements into heavier, starting with hydrogen being turned into helium, the models predict how the observable properties of the stars change as they age. These predictions can then be compared to observations. Important examples are the distributions of stars in terms of surface temperature and luminosity, particularly for stellar clusters where the stars, having presumably been formed in the same interstellar cloud, can be assumed to share the same age and original composition. These distributions are generally in reasonable agreement with the models; the comparison between observations and models furthermore provides estimates of the ages of the clusters, of considerable interest to the understanding of the evolution of the Galaxy. Additional tests, generally quite satisfactory, are provided in the relatively few cases where stellar masses can be determined with reasonable accuracy from the motion of stars in binary systems. Such successes give some confidence in the use of stellar models in other areas of astrophysics. These include studies of element synthesis in late stages of stellar evolution, the use of supernova explosions as `standard candles' in cosmology, and estimates of the primordial element composition from stellar observations.  An important aspect of stellar astrophysics is the use of stars as physics laboratories. Since the basic properties of stars and their modeling are presumed to be relatively well established, one may hope to use more detailed observations to provide information about the physics of stellar interiors, to the extent that it is reflected in observable properties. This is of obvious interest: conditions in the interiors of stars are generally far more extreme, in terms of temperature and density, than achievable under controlled circumstances in terrestrial laboratories. Thus sufficiently detailed stellar data might offer the hope of providing information on the properties of matter under these conditions.  Yet in reality there is little reason to be complacent about the status of stellar astrophysics. Most observations relevant to stellar interiors provide only limited constraints on the detailed properties of the stars. Where more extensive information is becoming available, such as determinations of detailed surface abundances, the models often fail to explain it. Furthermore, the models are in fact extremely simple, compared to the potential complexities of stellar interiors. In particular, convection, which dominates energy transport in parts of most stars, is treated very crudely while other potential hydrodynamical instabilities are generally neglected. Also stellar rotation is rarely taken into account, yet could have important effects on the evolution. These limitations could have profound effects on, for example, the modeling of late stages of stellar evolution, which depend sensitively on the composition profile established during the life of the star.  The Sun offers an example of a star that can be studied in very great detail. Furthermore, it is a relatively simple star: it is in the middle of its life, with approximately half the original central abundance of hydrogen having been used, and, compared to some other stars, the physical conditions in the solar interior are relatively benign. Thus in principle the Sun provides an ideal case for testing the theory of stellar evolution.  In practice, the success of such tests was for a long time somewhat doubtful. Solar modeling depends on two unknown parameters: the initial helium abundance and a parameter characterizing the efficacy of convective energy transport near the solar surface. These parameters can be adjusted to provide a model of solar mass, matching the solar radius and luminosity at the age of the Sun. Given this calibration, however, the measured surface properties of the Sun provide no independent test of the model. Furthermore, two potentially severe problems with solar models have been widely considered. One, the so-called faint early Sun problem, resulted from the realization that solar models predicted that the initial luminosity of the Sun, at the start of hydrogen fusion, was approximately 70 per cent of the present value, yet geological evidence indicated that there had been no major change in the climate of the Earth over the past 3.5 Gyr ({\\eg}, Sagan and Mullen, 1972).% \\footnote{The change in luminosity was noted by Schwarz\\-schild (1958) who speculated about possible geological consequences.} This change in luminosity is a fundamental effect of the conversion of hydrogen to helium and the resulting change in solar structure; thus the attempts to eliminate it resorted to rather drastic measures, such as suggestions for changes to the gravitational constant. As noted by Sagan and Mullen, a far more likely explanation is a readjustment of conditions in the Earth's atmosphere to compensate for the change in luminosity. A more serious concern was the fact that attempts to detect the neutrinos created by the fusion reactions in the solar core found values far below the predictions. This evidently raised doubts about the computations of solar models, and hence on the general understanding of stellar evolution, and led to a number of suggestions for changing the models such as to bring them into agreement with the neutrino measurements.  The last three decades have seen a tremendous growth in our information about the solar interior, through the detection and extensive observation of oscillations of the solar surface. Analyses of these oscillations, appropriately termed helioseismology, have resulted in extremely precise and detailed information about the properties of the solar interior, rivaling or in some respects exceeding our knowledge about the interior of the Earth. ",
        "conclusions": ""
    },
    "0207/astro-ph0207635_arXiv.txt": {
        "abstract": "When modeling the three-dimensional hydrodynamics of interstellar material rotating in a galactic gravitational potential, it is useful to have an analytic expression for gravitational perturbations due to stellar spiral arms.  We present such an expression for which changes in the assumed characteristics of the arms can be made easily and the sensitivity of the hydrodynamics to those characteristics examined.  This analytic expression also makes it easy to rotate the force field at the pattern angular velocity with little overhead on the calculations. ",
        "introduction": "\\label{section:Intro} In this paper we present analytic expressions for the perturbation of the galactic axisymmetric gravitational potential due to redistribution of part of the stellar disk mass into spiral arms, and for the distribution of density responsible for that potential.  Adjustable parameters include the number of arms, $N$, the pitch angle, $\\alpha$, the radial scale length of the dropoff in density amplitude of the arms, $R_{\\rm s}$, the midplane arm density, $\\rho_{\\rm o}$ at fiducial radius $r_{\\rm o}$, and the scale height of the stellar arm perturbation, $H$.  The amplitude of the spiral density distribution whose gravitational potential we set out to find is given by: \\begin{equation} \\rho_A(r, z) = \\rho_{\\rm o}  \\exp\\left(-\\frac {r-r_{\\rm o}}{R_{\\rm s}}\\right) {\\rm sech}^2\\left(\\frac{z}{H}\\right). \\label{equation:One} \\end{equation} Modulating this by a simple sinusoidal pattern in $\\phi$, following a logarithmic spiral with a pitch angle $\\alpha$, the overall density perturbation is \\begin{equation} \\rho(r, \\phi, z) = \\rho_A(r, z)\\cos(\\gamma) \\label{equation:Two} \\end{equation} where \\begin{equation} \\gamma = N \\left[\\phi - \\phi_p(r_{\\rm o}) - \\frac {\\ln(r/r_{\\rm o})}{\\tan(\\alpha)}\\right]. \\label{equation:Three} \\end{equation}  More complicated azimuthal arm structures can be constructed with linear combinations of these solutions of the form \\begin{equation} \\rho(r, \\phi, z) = \\rho_A(r, z) \\sum_n {\\rm C_n}\\cos( n \\gamma). \\label{equation:Four} \\end{equation} In a particularly interesting example, the density behaves approximately as a cosine squared in the arms but is separated by a flat interarm region occupying half the volume.  It has three terms in its sum, with $C_1 = 8/(3 \\pi) $, $C_2 = 1/2$, and $C_3 = 8/(15 \\pi)$.  The resulting phase pattern is compared with that of a simple sinusoid in Figure~\\ref{fig1}.  An important feature of such a perturbation is that its average density is zero.  Its potential can thus be added to observationally constrained models for the azimuthally averaged potential without altering the latter.  In addition, because the assumed arm perturbation extends over all radii, it is important that the average density be zero at both large and small radii where the arms do not actually exist.  The radial exponential damping introduced in Equation~\\ref{equation:One} is also useful in this regard.  Thus the gradient of the perturbation potential in the calculation region is provided predominantly by the local distribution of material. ",
        "conclusions": ""
    },
    "0207/astro-ph0207290_arXiv.txt": {
        "abstract": "\\noindent We present temperature and metallicity maps of the Perseus cluster core obtained with the {\\chandra} X-ray Observatory.  We find an overall temperature rise from $\\sim3.0$ keV in the core to $\\sim5.5$ keV at 120 kpc and a metallicity profile that rises slowly from $\\sim0.5$ solar to $\\sim0.6$ solar inside 60 kpc, but drops to $\\sim0.4$ solar at 120 kpc. Spatially resolved spectroscopy in small cells shows that the temperature distribution in the Perseus cluster is not symmetrical. There is a wealth of structure in the temperature map on scales of $\\sim10$ arcsec (5.2 kpc) showing swirliness and a temperature rise that coincides with a sudden surface brightness drop in the X-ray image. We obtain a metallicity map of the Perseus cluster core and find that the spectra extracted from the two central X-ray holes as well as the western X-ray hole are best-fit by gas with higher temperature and higher metallicity than is found in the surroundings of the holes. A spectral deprojection analysis suggests, however, that this is due to a projection effect; for the northern X-ray hole we find tight limits on the presence of an isothermal component in the X-ray hole, ruling out volume-filling X-ray gas with temperatures below 11 keV at 3$\\sigma$. ",
        "introduction": "\\label{intro}  The Perseus cluster, Abell\\,426, at a redshift of $z=0.0183$ or distance about 100~Mpc is the the nearest high luminosity cluster with a high central surface brightness \\citep[e.g., ][]{Fabian81,Allen01,Fabian94}. This makes it the brightest cluster in the X-ray sky.  The first {\\chandra} subarcsecond-resolution X-ray images of the cluster core around the central dominant galaxy NGC\\,1275 were presented by \\citet[][ F00]{Fabian00a}. Using X-ray colours F00 showed that the temperature of the gas decreases from about 6.5~keV to 3~keV inward to the central galaxy NGC\\,1275, which is surrounded by a spectacular low-ionization, emission line nebula \\citep[][ see also \\citealt*{Conselice01}]{Lynds70}.  The nucleus powers the radio source 3C84 (Pedlar et al 1990) which has structures on various scales. The 0.5 arcmin-sized radio lobes coincide with holes in the soft X-ray emission (\\citealt{Boehringer93,McNamara96}; F00). {\\chandra} has revealed that the holes are not due to absorption and have X-ray bright rims, which are cooler than the surrounding gas (F00). It was found that the rims are not distinguishable as sharp features on a 3--7~keV image and are therefore not shock features, contrary to the early prediction of \\citet*{Heinz98}, and are not expanding supersonically. The simplest interpretation of the low surface brightness is that they are devoid of X-ray gas and have pressure support from cosmic rays and magnetic fields. However they may contain some hotter gas at the virial temperature, or even above, of the cluster with the radio plasma having a low filling factor. In this scenario the rims consist of cooler gas which has been swept aside. The two outer holes, one of which was previously known, were identified with recently-found spurs of low-frequency radio emission by \\citet{Blundell00}.  The core of the Perseus cluster offers the rare opportunity to study the radio source/intracluster gas interaction on arcsecond scales. In the present study we have carried out a spectral analysis of the {\\chandra} images which provides us detailed information about the cluster physics. In Sect.~\\ref{observations} we describe the observations including the full 29.0 ks {\\chandra} image of the Perseus cluster core. In Sect.~\\ref{specanalysis} we describe our method of spectral analysis and show temperature and metallicity profiles, as well as maps. In Sect.~\\ref{specdeproj} we carry out a spectral deprojection analysis of the cluster and search for the presence of X-ray gas in the northern X-ray hole. In Sect.~\\ref{conclusions} we conclude with a summary and a discussion.  We use $H_0=50\\kmpspMpc$ and $q_0=\\frac{1}{2}$ throughout. Unless otherwise stated, quoted error bars are 1\\,$\\sigma$ (68.3\\% confidence). ",
        "conclusions": "\\label{conclusions}  We have presented in this paper the results from a detailed temperature and metallicity mapping of the Perseus cluster core using {\\chandra} observations.  We work with the complete useful ACIS-S data set comprising a total of 23.9 ks. Beginning by averaging over large annuli around the nucleus, we find that the temperature averaged in such annuli rises smoothly from $\\sim3.0$ keV to $\\sim5.5$ keV at 120 kpc. A small kink is found, however, in this profile at $\\sim80$ kpc, which corresponds to a surface brightness drop seen in the X-ray image. Using solar abundances according to \\citet{Anders89}, the metallicity rises from $Z\\sim0.4$ solar with decreasing radius to a maximum $Z\\sim0.6$ solar at $\\sim60$ kpc. There may be a peak of the metallicity profile at $\\sim50-60$ kpc, which is reminiscent of such peaks in other clusters \\citep[e.g., ][]{Sanders01}. However, inside 60 kpc the metallicity varies little, with a central metallicity of $Z\\sim 0.5$ solar. We have also carried out a spectral deprojection analysis of the radially averaged profile, which yields a very similar picture.  Spatially resolved spectroscopy in small cells shows that the temperature distribution in the Perseus cluster is not symmetrical. In fact, the distribution of cold ($2-3$ keV) gas in the centre appears to spiral outward and corresponds to the swirly appearance of the X-ray emission (see Figs.~\\ref{combined} and~\\ref{500counts}). This may suggest the presence of angular momentum of the intracluster gas (F00). There are, however, other possibilities, such as the model by \\citet{Churazov00} (see also their adaptively smoothed ROSAT image) who propose that rising radio bubbles (such as the X-ray holes) are responsible for the overall spiral structure. The outer surface brightness drop is traced by a rise in temperature (with increasing radius). This is reminiscent, although not as large, of the temperature drop at the cold front seen by \\citet{Markevitch00}.  If thermal conduction is an important effect in the intracluster medium, such fronts could reflect the magnetic field structure.  A comparison of the metallicity distribution with the galaxy distribution from an optical image from the Digital Sky Survey does not show any obvious correlation between the galaxies and regions of higher metallicity. Projection effects and the resolution of the {\\chandra} metallicity map, however, do not allow to rule out any correlation yet.  We find that the spectra extracted from the two central X-ray holes as well as the western X-ray hole are best-fit by hotter and more metal rich gas than their immediate surroundings. Using a spectral deprojection under the assumption of spherical symmetry in a sphere segment containing the northern inner X-ray hole, we have tested whether this could be due to a projection effect. Interestingly, we find that most of the X-ray emission in the hole can be explained by the projected emission of the shells further out. This directly addresses the issue of the X-ray gas content of the X-ray holes mentioned in the introduction; we find tight limits on the presence of an isothermal component in the X-ray hole, ruling out volume-filling X-ray gas with temperatures below 11 keV at 3$\\sigma$.  The temperature distribution of the gas in the vicinity of the X-ray holes is of great importance for theories of radio source heating to counter the short cooling time of the X-ray gas in this cluster (e.g., F00). Longer exposure times are needed to reveal if there is cool gas associated with the wake of the bubble, as, for example, in the model by \\citet{Churazov00}, who proposed that the radio bubbles would transport cooler gas from the cluster centre `upwards'. The bubbles may also transport or drag metals upwards. The subject of radio bubbles and their interaction with the intracluster gas has recently become the subject of much work (\\citealt*{Churazov01,Brueggen01,Quilis01,McNamara01}; F02). The {\\chandra} results provide essential observational input to test these theories."
    },
    "0207/astro-ph0207223_arXiv.txt": {
        "abstract": "Latest developments in theoretical computations since the international Opacity Project (OP), under the new the Iron Project (IP) and extensions, are described for applications to a variety of objects such as stellar atmospheres, nebulae, and active galactic nuclei. The primary atomic processes are: electron impact excitation (EIE), photoionization, electron-ion recombination, and bound-bound transitions, all considered using the accurate and powerful R-matrix method including relativistic effects. As an extension of the OP and the IP, a self-consistent and unified theoretical treatment of photoionization and recombination has been developed. Both the radiative and the dielectronic recombination (RR and DR) processes are considered in a unified manner. Photoionization and recombination cross sections are computed with identical wavefunction expansions, thus ensuring self-consistency in an ab initio manner. The new unified results differ from the sum of previous results for RR and DR by up to a factor of 4 for the important but complex atomic systems such as Fe~I~-~V. The fundamental differences are due to quantum mechanical intereference and coupling effects neglected in simpler approximations that unphysically treat RR and DR separately, which can not be independently measured or observed. The electronic, web-interactive, database, TIPTOPBASE, to archive the OP/IP data in a readily accessible manner is also described. TIPTOPBASE would include electron-ion recombination data and new fine structure transition probabilities. Efficient codes developed by M.J. Seaton to calculate `customized' mixture opacities and radiative accelerations ('levitation') in stars will also be available. ",
        "introduction": "At densities and temperatures in stellar atmospheres, many atomic levels are excited under non-local thermal equilibrium (NLTE) conditions. NLTE models and other applications such as stellar opacities require large amount of accurate atomic parameters for collisional and radiative processes to describe radiation transfer, spectral analysis etc. The collisional process is primarily electron impact excitation (EIE), while radiative processes are photoionization, electron-ion recombination, and bound-bound transitions. These four basic dominant atomic processes in the plasmas can be described as follows.  \\noindent i) Electron-impact excitation (EIE) of an ion $X^{+z}$ of charge z:  $$e + X^{+z} \\rightarrow e' + X^{+z*}.$$  \\noindent ii) Radiative bound-bound transitions:  $$X^{+z} + h\\nu \\rightleftharpoons X^{+z*},$$  \\noindent iii) Photoionization (PI) by absorption of a photon:  $$X^{+z} + h\\nu \\rightleftharpoons X^{+z+1} + \\epsilon.$$  \\noindent The inverse process of PI is electron-ion radiative recombination (RR).  \\noindent iv) Autoionization (AI) and dielectronic recombination (DR):  $$e + X^{+z} \\rightarrow (X^{+z-1})^{**} \\rightarrow \\left\\{ \\begin{array}{ll} e + X^{+z} & \\mbox{AI} \\\\ X^{+z-1} + h\\nu & \\mbox{DR} \\end{array} \\right. $$  \\noindent The inverse process of DR is photoionization via the intermediate doubly excited autoionizing states, i.e. resonances in atomic processes. At prevailing densities and temperatures in stellar atmospheres, the role of metastable states and low-lying fine structure levels in photoionization and recombination of ions bears special emphasis.  Collisional and radiative atomic process have been studied in ab initio manner under the OP ({\\it The Opacity Project} 1995, 1996), and the IP (Hummer et. al 1993). The close coupling R-matrix methodology employed under the OP and IP enables the computation of self-consistent sets of atomic parameters, thereby reducing uncertainties in applications involving different processes and approximations. Sample results obtained under the two projects are presented. ",
        "conclusions": "The current status of large-scale ab initio close coupling R-matrix calculations for radiative and collisional processes is reported. The Iron Project Breit Pauli R-matrix radiative calculations include large numbers of dipole allowed and intercombination transitions. Self-consistent sets of atomic data for photoionization and unified (electron-ion) recombination (including RR and DR) are obtained, and should yield more accurate photoionization models. Work is in progress for heavy ions of the iron group elements."
    },
    "0207/astro-ph0207153_arXiv.txt": {
        "abstract": "We present the results of a systematic study of the formation and evolution of binaries containing black holes and normal-star companions with a wide range of masses. We first reexamine the standard formation scenario for close black-hole binaries, where the progenitor system, a binary with at least one massive component, experienced a common-envelope phase and where the spiral-in of the companion in the envelope of the massive star caused the ejection of the envelope. We estimate the formation rates for different companion masses and different assumptions about the common-envelope structure and other model parameters. We find that black-hole binaries with intermediate- and high-mass secondaries can form for a wide range of assumptions, while black-hole binaries with low-mass secondaries can only form with apparently unrealistic assumptions (in agreement with previous studies).  We then present detailed binary evolution sequences for black-hole binaries with secondaries of 2 to 17\\Msun\\ and demonstrate that in these systems the black hole can accrete appreciably even if accretion is Eddington limited (up to 7\\Msun\\ for an initial black-hole mass of 10\\Msun) and that the black holes can be spun up significantly in the process.  We discuss the implications of these calculations for well-studied black-hole binaries (in particular GRS 1915+105) and ultra-luminous X-ray sources of which GRS 1915+105 appears to represent a typical Galactic counterpart.  We also present a detailed evolutionary model for Cygnus X-1, a massive black-hole binary, which suggests that at present the system is most likely in a wind mass-transfer phase {\\em following} an earlier Roche-lobe overflow phase. Finally, we discuss how some of the assumptions in the standard model could be relaxed to allow the formation of low-mass, short-period black-hole binaries which appear to be very abundant in Nature. ",
        "introduction": "There are currently 17 binary systems containing black holes for which dynamical mass estimates are available (see e.g. Table 1 of Lee, Brown \\& Wijers 2002 [LBW], and references therein; Orosz et al.\\ 2002). According to conventional wisdom, these systems formed from primordial binaries where at least one of the stars was quite massive (i.e. $M \\ga 20$\\,--\\,$25\\Msun$).  If mass transfer from the primary to the secondary commences at an orbital period in the range of $\\sim$ 1\\,--\\,10\\yr, a common envelope may form during which the hydrogen-rich envelope of the primary is expelled (Paczy\\'nski 1976). If the secondary and the core of the primary avoid a merger, then the massive core may evolve to core collapse and the formation of a black hole in a close binary.  For 9 of the 17 black-hole binaries (see e.g. LBW), the current-epoch companion mass is $\\la 1\\Msun$ and the orbital periods are $\\la 1$ day.  For reasons discussed later in the text, these systems probably had primordial secondaries whose mass was not substantially greater than $\\sim 1.5 \\Msun$ (but see \\S~4.4).  One quantitative difficulty with the common-envelope scenario for forming this type of black-hole binary is that the amount of orbital energy that can be released by the spiral-in of a low-mass secondary may not be sufficient to eject the massive envelope of the primary.  It has long be recognized that this is energetically challenging even if the common-envelope ejection mechanism is very efficient (Podsiadlowski, Cannon \\& Rees 1995; Portegies Zwart, Verbunt \\& Ergma 1997; Kalogera 1999; see, however, also Romani 1992). Furthermore, recent determinations of the binding energy of the envelopes of massive supergiants by Dewi \\& Tauris (2000, 2001) suggest that all studies so far may have significantly underestimated how tightly bound these envelopes actually are, which seriously aggravates the problem.  On the other hand, it has been estimated that there may be up to several thousand low-mass black-hole transients in the Galaxy (Wijers 1996; Romani 1998). This has led to several alternative formation scenarios for low-mass black-hole binaries, where either the low-mass companion is a third star in a triple system being captured into a tight orbit when the two massive components merge (Eggleton \\& Verbunt 1986), or where the low-mass star forms after the black hole -- out of a collapsed massive envelope (Podsiadlowski et al.\\ 1995). In the present work we reexamine the standard formation scenarios for low-mass black hole binaries with plausible modifications to some of the usual assumptions  By contrast, for 4 of the black-hole binaries, the mass of the companion is substantially larger (i.e. $\\ga 6 \\Msun$), and the availability of orbital binding energy for ejecting the common envelope is greatly enhanced.  The remaining 4 systems (4U 1543-47, GRO J1655-40, GS 2023+338, and GRS 1915+105) have either intermediate-mass donor stars (i.e. $2 \\la M_{\\rm d} \\la 5\\Msun$) or orbital periods longer than 2.5 days, thereby allowing for primordial secondaries of at least intermediate mass, and substantial mass loss or evolution of the secondary to its present status as the donor star. It is on the evolution of these latter two categories, with particular emphasis on GRS 1915+105, that we focus this work (for other recent discussions of intermediate-mass black-hole binaries see Kalogera 1999; Brown et al.\\ 2000; LBW).  In addition to the common-envelope ejection mechanism, another major uncertainty in the modelling of black-hole binaries is the initial mass of the black hole which is caused by uncertainties in the theory of both single and binary stellar evolution. Some of the key factors that determine the maximum initial black-hole mass are (1) the minimum initial mass above which a star leaves a black-hole remnant (mostly believed to be in the range of 20\\,--\\,25\\Msun; Maeder 1992; Woosley \\& Weaver 1995; Portegies Zwart et al.\\ 1997; Ergma \\& Fedorova 1998; Ergma \\& van den Heuvel 1998; Brown et al.\\ 2000; Fryer \\& Kalogera 2001; Nelemans \\& van den Heuvel 2001; cf Romani 1992), (2) the minimum mass above which a single star loses its envelope in a stellar wind and becomes a helium/Wolf-Rayet star, (3) the maximum radius of a single star before and after helium core burning, (4) the amount of wind mass loss in the Wolf-Rayet phase and (5) the fraction of the mass that is ejected when the black hole forms (for detailed recent discussions see Brown et al.\\ 2000; Fryer \\& Kalogera 2001; Nelemans \\& van den Heuvel 2001). Generally one expects the most massive black holes to form from stars that have an initial mass close to the minimum mass above which a star loses its hydrogen-rich envelope in a stellar wind and becomes a Wolf-Rayet star, and where the common-envelope phase occurs near the end of the evolution of the massive primary (i.e. experiences case C mass transfer; Brown, Lee \\& Bethe 1999; Wellstein \\& Langer 1999).  This avoids a long phase where the mass of the helium star, the black-hole progenitor, is reduced by a powerful stellar wind, as typically seen from Wolf-Rayet stars, which would reduce the final helium-star mass and hence the maximum black-hole mass (see e.g. Woosley, Langer \\& Weaver 1995)\\footnote{It should also be noted that in the formation of some black holes (e.g. the black hole in Nova Scorpii) a significant fraction of the mass of the helium star is ejected in the supernova explosion in which the black hole formed (Podsiadlowski et al.\\ 2002; LBW). Thus the final helium-star mass strictly provides only an upper limit on the black-hole mass.}. Unfortunately, the evolutionary tracks for massive post-main-sequence stars and in particular the maximum radius a star attains after helium core burning are rather uncertain (and generally inconsistent with observed distributions of stars in the Hertzsprung-Russell diagram; see e.g. Langer \\& Maeder 1995).  Fryer \\& Kalogera (2001) have shown that initial black-hole masses as high as 15\\Msun\\ can be obtained if either the parameter range for case C mass transfer is increased or the wind mass-loss rate in the helium-star phase of the black-hole progenitor is reduced (see also Brown et al.\\ 2001; Nelemans \\& van den Heuvel 2001; Belczynski \\& Bulik 2002; Pols \\& Dewi 2002)  Only a few of the previous studies (e.g. LBW) have considered the possibility that the black hole may increase its mass substantially since its formation by mass transfer from the companion star. It is one of the purposes of this paper to demonstrate that accretion from a companion star can substantially increase the mass of a black hole and spin it up in the process and that the present mass may not be representative of the initial black-hole mass.  A closely coupled result is that the observed donor star masses may be substantially lower than their initial mass.  The paper is structured as follows. In \\S~2 we present detailed binary population synthesis calculations to show how the formation rate of black-hole binaries and the distribution of the secondary masses depend on the structure of massive supergiant envelopes and the modelling of common-envelope ejection. In \\S~3 we discuss the results of extensive binary evolution calculations for black-hole binaries with intermediate-/high-mass secondaries, which we then apply in \\S~4 to observed systems, in particular GRS 1915+105, ultraluminous X-ray sources and Cyg X-1.  Finally, in \\S~4 we reexamine the standard formation scenario for low-mass black-hole binaries to understand why such systems appear to be so plentiful in Nature. ",
        "conclusions": ""
    },
    "0207/astro-ph0207479_arXiv.txt": {
        "abstract": "We examine the gas kinematics of the LMC revealed by the high spatial and velocity resolution H\\,I data cube of the combined ATCA and Parkes telescope surveys. We adopt an approach designed to facilitate comparisons with quasar absorption line observations, in particular restricting our analysis to pointings with H\\,I column density satisfying the damped \\lya (DLA) criterion.  We measure velocity widths for $\\approx 5000$ random pointings to the LMC and find a median value of $\\approx 40 \\mkms$ which modestly exceeds the value predicted by differential rotation. This median is significantly lower than the median value observed for the metal-line profiles of high $z$ DLA. Therefore, assuming the metal-line profiles track H\\,I kinematics, the velocity fields of high $z$ DLA are inconsistent with the kinematics of low-mass galaxies like the LMC. We also investigate the kinematic characteristics of the giant H\\,I shells which permeate the LMC.  These shells impart an additional 10 to 20~km/s to the velocity widths in $\\approx 20\\%$ of random pointings to the LMC. These non-gravitational motions are insufficient to explain the DLA kinematics even if the shells had significanlty larger covering fraction in the early universe. To account for the DLA with processes related to SN feedback requires winds with a qualitatively different nature than those observed in the LMC. ",
        "introduction": "\\label{sec-intro}  The introduction of echelle spectrographs on 10m class telescopes has revealed the velocity fields of gas at high $z$ with unprecedented precision.  In the low density \\lya forest, for example, observations have investigated the thermal history and physical nature of the gas comprising the intergalactic medium \\citep[e.g.][]{kirkman97,rauch98,bryan00,theuns02}. Similarly, studies of optically thick absorbers probe the velocity fields of gas in non-linear, collapsed or collapsing structures \\citep[hereafter PW97]{charlton96,rauch97,pw97}. Together these observations constrain the nature of galaxy formation and examine the processes of structure formation in the early universe.  Although every quasar sightline is rich in detail, their 'pinhole' nature complicates interpretation. Quasar sightlines produce a core-sample of an absorption system with data acquired along a single dimension and with resolution along the redshift axis only. In general, one must assume a specific geometry in order to impose a spatial delineation to the observations. In a few rare cases, lensed quasars or quasar pairs allow a multi-dimensional analysis \\citep[e.g.][]{dinshaw95,lopez99,rauch99,dodorico02}, but even these observations are difficult to unravel. In turn, one frequently finds that several models can be introduced to explain specific sets of observations.  A good example of this degeneracy is found in the damped \\lya systems (DLA), quasar absorption line (QAL) systems believed to represent the progenitors of present galaxies \\citep{kauff96,steinmetz01}. PW97 and \\cite{pw98} analysed the velocity fields of over 30 DLA and ruled out several plausible morphologies for these high $z$ galaxies, but their true physical nature remains uncertain. The DLA low-ion\\footnote{The term 'low-ion' refers to the dominant ionization state of an element in an H\\,I gas (e.g.\\ Fe$^+$, Si$^+$).} kinematics are remarkably well described by a thick, rotating disk (PW97) but also by CDM scenarios involving the merging of multiple 'clumps' bound to individual dark matter halos \\citep{hae98,mcd99,maller01}. Furthermore, several authors have claimed the damped \\lya systems might be explained by outflows from SN winds \\citep{nulsen98,schaye01}. Owing to the 'pinhole' nature of QAL studies and the absence of spatial resolution, it is difficult to distinguish between these very different models.  One approach toward resolving the ambiguities of QAL research is to repeat these experiments on present-day galaxies whose physical properties and gas kinematics are well understood. Comparisons with velocity profiles obtained from high $z$ protogalaxies could provide fresh insight into the physical nature of young galaxies. Regarding the damped \\lya systems, any galaxy with large H\\,I surface density is an analog. Surveys of H\\,I-selected galaxies suggest the damped \\lya systems should exhibit a range of luminosity including both dwarf galaxies and large spiral galaxies \\citep{zwaan99,zwaan02,ryan02,rosenberg02}. This includes the Milky Way, the Magellanic Clouds, and many other Local Group galaxies. The Milky Way aside, it is difficult to probe even a single absorption sightline through these galaxies because present technology limits UV spectroscopy to magnitudes $V<16$, restricting the number of potential background targets to $\\approx 0.01$ per square degree \\citep[e.g.][]{shull02}. Without a major advance in UV space telescopes, one must consider alternative approaches.  In this paper we study the H\\,I velocity fields of the Large Magellanic Cloud (LMC). Optically, the LMC has a bar but otherwise has a chaotic appearance being peppered with bright H\\,II regions. A stellar color-magnitude study by \\cite{smecker02} indicates that the LMC bar was formed in a major episode around 5~Gyr ago, whilst the disk component of the LMC was more gradually formed. In some respects, the LMC may have similar properties to the compact galaxies seen in the Hubble deep field \\citep{colley97}. The smooth disk component of the LMC is more easily seen in  the 2MASS/DENIS infrared study of \\cite{vdmarel01}. In terms of gas kinematics, the ATCA HI study of \\cite{kim98b} reveals the disk to be in regular rotation, though many sightlines in the eastern half display multiple velocity components. Deviations from regular rotation appear to be a result of star formation activity \\citep{kim99}, tidal forces \\citep{weinberg00} and, possibly, interaction with the halo of the Milky Way \\citep{deboer98}.  The following analysis examines the high resolution 21cm data cubes constructed from the combined ATCA survey of \\cite{kim98a} and the Parkes telescope survey of \\cite{staveley02}. We perform a statistical investigation designed to facilitate comparisons with QAL observations and shed new light on their interpretation, in particular those of the damped \\lya systems. Furthermore, our analysis examines the kinematic impact of the H\\,I shells which permeate the LMC \\citep[e.g.][]{kim99}. We assess the contribution of their non-gravitational motions to the observed velocity fields and consider the implications for interpretations of the damped \\lya kinematics.  Finally, we propose future observations which will build upon our treatment. ",
        "conclusions": "H$\\alpha$ emission maps of the LMC first revealed the presence of H\\,I shells throughout the galaxy \\citep[e.g.][]{meaburn80}.  These shells and the holes they encompass are a generic feature of many nearby galaxies as cataloged through H\\,I observations (e.g.\\ M33: Deul \\& den Hartog 1990; SMC: Staveley-Smith et al.\\ 1997).  Their origin is linked to the combined radiative and kinematic pressure of stellar winds from massive stars and SN feedback \\citep{weaver77}.  \\cite{kim99} has extensively surveyed the H\\,I shells in the LMC and presented a classification scheme which we adopt: giant shells are H\\,I shells confined to the main H\\,I layer of the LMC (GS; $\\ell < 360$~pc) and supergiant shells have sizes which significantly exceed the H\\,I layer (SGS; $\\ell > 360$~pc).  The shells have sizes ranging from radii of 40~pc to 1.2~kpc following a power law with slope $\\alpha = -1.5 \\pm 0.4$ over the interval 100 to 1000~pc. The expansion velocity of the shells is well correlated with radius: $v_{exp} \\approx 15 \\mkms$ for the smallest shells and $v_{exp} \\approx 20-35 \\mkms$ for the largest.  Of primary interest to the current analysis is addressing the impact of these expanding H\\,I shells on the gas kinematics of the LMC. In terms of the DLA, if gas-rich protogalaxies correspond primarily to dwarf galaxies at high $z$, then one must introduce non-gravitational motions to explain the larger velocity widths. Several authors have hypothesized that SN winds, for example, may explain the DLA kinematics \\citep{nulsen98} and the velocity profiles observed in some Lyman limit systems \\citep{rauch99,bond01}. In the LMC, we can examine the effects of SN feedback on the gas kinematic characteristics.   \\begin{figure*} \\begin{center} \\includegraphics[height=6.0in, width=4.0in, angle=-90]{f3.eps} \\caption{ Representative pointings toward the expanding H\\,I shell surrounding 30~Dor in the LMC.  Many profiles toward expanding shells exhibit the double-peaked characteristic evident in the left-hand panels while the remaining are similar to the profiles in the right-hand panels. } \\label{fig:bubble} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[height=4.4in, width=4.0in]{f4.eps} \\caption{ Velocity width distributions for the complete LMC sample (top panel) compared against random pointings to a series of H\\,I shells identified by \\cite{kim99}.  In general, the expanding H\\,I shells exhibit systematically larger velocity widths, with median offsets of 10 to 20~km/s.  The $P_{KS}$ probabilities indicate that these shells exhibit significantly different kinematics than the typical LMC sightline. } \\label{fig:lmcdelv} \\end{center} \\end{figure*}   We have investigated this aspect of the LMC velocity fields as follows. We selected $\\approx 300$ random pointings toward several H\\,I shells identified by \\cite{kim99} and maintained the $\\N{HI} = 2 \\sci{20} \\cm{-2}$ threshold. We then measured the velocity width of each pointing with exactly the same technique as the complete LMC sample. Figure~\\ref{fig:bubble} presents a representative sample of 4 sightlines through the shell surrounding 30~Dor.  In the left panels, one observes the double-peaked profiles characteristic of an expanding (or infalling) shell.  A significant fraction of the sightlines share this characteristic. The remaining profiles are more similar to the right panels; these show nearly uniform optical depth or some mild asymmetry. Figure~\\ref{fig:lmcdelv} summarizes the velocity width distributions for 6 H\\,I shells.  For comparison, we plot the complete LMC $\\delv$ distribution (top panel). The $P_{KS}$ values give the probability that the kinematic characteristics of the gas comprising the H\\,I shells are consistent with the general LMC.  In most cases, the distributions are inconsistent with this hypothesis at $< 0.01$ probability owing to the larger velocity widths observed through the expanding H\\,I shells.  Toward 30~Dor, for example, the median $\\delv$ value is $\\approx 15 \\mkms$ larger than the LMC sample.  Our analysis demonstrates that the kinematics associated with expanding H\\,I shells can significantly influence the H\\,I kinematics within low-mass galaxies.  These non-gravitational velocity fields impart an additional 10 to 20~km/s to the velocity widths which corresponds to $33-50 \\%$ of the median LMC velocity width. At the same time, the significant difference between the $\\delv$ distributions of the H\\,I shells and the LMC sample indicate their filling factor and, therefore, their overall kinematic impact is relatively small. We found that the fraction of the H\\,I data cube covered by H\\,I shells satisfying the DLA threshold is only $20 \\%$. This fraction is large enough to comprise most of the high $\\delv$ tail observed in the LMC distribution, yet too small to grossly affect the LMC H\\,I kinematics.   The results presented in Figure~\\ref{fig:lmcdelv} suggest that the H\\,I shells produced via SN feedback are unlikely to explain the damped \\lya kinematics.  Because the H\\,I shells typically contribute $< 20 \\mkms$ to the observed velocity widths, these non-gravitational motions would comprise only a small fraction of even the median $\\delv$ value observed in the DLA.  Even if the porosity of these H\\,I shells was significantly larger in high $z$ dwarf galaxies, the H\\,I kinematics would still be inconsistent with the majority of damped systems. As the porosity reached unity, however, the velocity and scale of the shells might have a qualitatively different nature, i.e., a superwind. Indeed, if SN feedback is to explain the DLA kinematics, the implied velocity fields must have a very different nature than the expanding H\\,I shells observed in the LMC. These superwinds may be more prevalent at high redshift owing to higher star formation rates, feedback from Population III supernovae, or even differences in the relative contribution of Type~Ia and Type~II SN."
    },
    "0207/astro-ph0207365_arXiv.txt": {
        "abstract": "High resolution rest frame UV quasar absorption spectra covering low and high ionization species, as well as the Lyman series lines, provide remarkably detailed information about the gaseous phases of galaxies and their environments.  For redshifts less than 1.5, many important chemical transitions remain in the observed ultraviolet wavelength range.  I present examples of absorption that arises from lines of sight through a variety of structures, drawn from UV spectra recently obtained with STIS/HST.  Even with the greater sensitivity of COS/HST there will be a limit to how many systems can be studied in detail. However, there is a great variety in the morphology of the phases of gas that we observe, even passing through different regions of the same galaxy.  In order to compile a fair sample of the gaseous structures present during every epoch of cosmic history, hundreds of systems must be sampled.  Multiple lines of sight through the same structures are needed, as well as some probing nearby structures whose luminous hosts have been studied with more standard techniques.  Combined with high resolution optical and near--IR ground--based spectra, it will be possible to uniformly study the gaseous morphologies of galaxies of all types through their entire evolutionary histories. ",
        "introduction": "The tool of quasar absorption line spectroscopy has several distinct advantages over more traditional imaging studies of distant galaxies. Spectra covering absorption lines from numerous chemical elements in various states of ionization can yield detailed information about the physical conditions in the gaseous components of the universe.  This is not limited to only the most luminous components.  Dwarf galaxies, low surface brightness galaxies, and even intergalactic structures are probed by quasar lines of sight.  Quasar absorption lines can be used to study gas during the birth and death of stars and of entire galaxies.  The kinematic information contained in high resolution spectra allows us to study processes.  This is not just a still picture snapshot; it is more like a short movie.  Finally, the same level of detail is available for our study at all redshifts because the same method can be applied using optical and near--IR spectroscopy.  However, the study of quasar absorption lines also presents some challenges if we are to extract from the quasar spectra all of the detailed information about physical conditions of gas. Imagine if we had to classify a galaxy according to its standard morphological type by zooming in on an image of just a small part of the galaxy.  We have to consider carefully what a single line of sight tells us about the global conditions in galaxies. In fact, it should be possible to learn a great deal if we consider the evolution of the ensemble of gaseous structures probed by quasar lines of sight.  In order to realize this potential, we must learn to connect absorption signatures to the physical conditions of the phases of gas that produce them, and to the processes that give rise to such signatures in local galaxies.  For the optimal study of the detailed physical conditions along a quasar line of sight we need high resolution spectroscopy covering all rest frame UV transitions.  For redshifts less than one, a significant fraction of the key transitions still appear in the ultraviolet. In this proceedings, I review examples of various systems that have been studied in detail, focusing on the question of what additional data would yield a significant advance.  The goal is to define capabilities for the optimal UV spectrograph and telescope to be used for such a program.   \\begin{figure} \\plotone{charlton.fig1.eps} \\caption {Selected transitions for the $z=0.9902$ strong \\hbox{{\\rm Mg}\\kern 0.1em{\\sc ii}} system along the PG~$1634+706$ line of sight, presented in velocity space.  The velocity zero--point corresponds to the apparent optical depth centroid for the \\hbox{{\\rm Mg}\\kern 0.1em{\\sc ii}} profile.  The \\hbox{{\\rm Mg}\\kern 0.1em{\\sc i}} and \\hbox{{\\rm Mg}\\kern 0.1em{\\sc ii}} profiles were obtained with HIRES/Keck, with $R=45,000$ (P.I. Churchill), and the other transitions were observed with STIS/HST, with $R=30,000$ (P.I.'s Jannuzi and Burles).} \\end{figure} ",
        "conclusions": "The small sample of systems presented here demonstrates great promise for learning about the detailed physical conditions in a variety of gaseous environments.  For these systems, the data have some limitations and, as a result, some questions remain unanswered. These highlight the need for additional UV spectroscopy capabilities. I conclude by presenting a ``wish list'' if we aim to construct the dynamical history of the ensemble of galaxies and gaseous structures:  \\begin{itemize} \\item{hundreds of systems, or even thousands!} \\item{cover all key chemical species from the rest frame ultraviolet} \\item{study many systems that are at low enough redshift to image their galaxy hosts; these serve as a calibrator for higher redshift systems} \\item{higher spectral resolution ($R=100,000$ or even higher) for a subset of the systems in order to resolve interstellar medium features} \\item{high signal--to--noise ($>20$) for a subset to enable abundance pattern studies} \\item{multiple lines of sight through the same objects} \\item{a separate, detailed study of the interstellar medium of the Milky Way and of nearby galaxies using the same techniques} \\end{itemize}  Support for this work was provided by the NSF (AST--9617185) and by NASA (NAG 5--6399 and HST--GO--08672.01--A), the latter from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS5--26555."
    },
    "0207/astro-ph0207614_arXiv.txt": {
        "abstract": "A spatially unresolved velocity feature, with an approaching radial velocity of $\\approx$~100 \\kms\\ with respect to the systemic radial velocity, in a position--velocity array of \\oiii\\ line profiles is identified as the kinematical counterpart of a jet from the proplyd LV~5 (158--323) in the core of the Orion Nebula. The only candidate in HST imagery for this jet appears to be a displaced, ionized knot. Also an elongated jet projects from the proplyd GMR~15 (161--307). Its receding radial velocity difference appears at $\\approx$~80~\\kms\\ in the same position--velocity array.  A `standard' model for jets from young, low mass stars invokes an accelerating, continuous flow outwards with an opening angle of a few degrees. Here an alternative explanation is suggested which may apply to some, if not all, of the proplyd jets. In this, a `bullet' of dense material is ejected which ploughs through dense circumstellar ambient gas. The decelerating tail of material ablated from the bullet's surface would be indistinguishable from a continuously emitted jet in current observations. ",
        "introduction": "The nature of the compact gaseous knots, dubbed `proplyds' by O'Dell, Wen \\& Hu (1993), in the close vicinity of the O6 star \\thetaC\\ in the Orion Nebula, is becoming increasingly clear. They were first discovered (LV 1-6) by Laques \\& Vidal (1979) in  the optical emission lines with many more identified later as sub-arcsec diameter thermal radio sources (Churchwell et al. 1987; Garay 1987; Garay, Moran \\& Reid 1987 - hereafter GMR; Felli et al. 1993a; Felli et al. 1993b). Each proplyd was shown to contain a low mass star (Meaburn 1988; McCaughrean \\& Stauffer 1994) whose youth is suggested by this partial cocoon of primaeval material (see the dramatic HST images in O'Dell, Wen \\& Hu 1993; O'Dell \\& Wong 1996; Bally et al. 1998). The $\\approx$ 50 \\kms\\ photoevaporated flows, driven by the intense Lyman flux of \\thetaC, from the ionized proplyd surfaces (Meaburn 1988) and their interactions with the particle wind from \\thetaC\\ to form stand-off bow-shocks (Hayward, Houck \\& Miles 1994) have most recently been considered by Henney et al. (2002) and Graham et al. (2002) and references therein.  Jets from young stellar objects (YSOs) were considered as one plausible explanation of the $\\geq$ 100 \\kms, highly collimated ($\\leq$ 1 \\arcsec\\ wide) velocity `spikes' on longslit position-velocity (pv) arrays of \\oiii\\ line profiles in ground-based observations of  proplyds (Meaburn et al. 1993; Massey \\& Meaburn 1995). The ubiquity of such jets from the proplyds then became immediately apparent in the HST imagery of Bally, O'Dell \\& McCaughrean (2000). The HST spectral observations, with STIS at 0.1\\arcsec\\ resolution of the \\ciii\\ profiles from the jet of LV~2, showed the jet outflow to have a radial velocity extent of $\\approx$~120 \\kms\\ (Henney et al. 2002). This is consistent with the extent of the velocity spike in the ground-based 1\\arcsec\\ resolution observations of the \\oiii\\ profiles from LV 2 (fig. 5 in Meaburn et al. 1993; Henney et al. 2002). Consequently, it is safe to assume that the narrow velocity spikes in all of the pv arrays of  \\oiii\\ profiles observed from the ground in other proplyds are a measure of their jet velocities.  In the present paper this assumption has been applied to two proplyds, LV~5 (158--323) and possibly GMR~15 (161--307), where high-speed, collimated, velocity features (`spikes') have been found in the pv arrays of \\oiii\\ profiles observed from the ground and where convincing HST images of their jets exist. The bracketed identifications are from O'Dell \\& Wen (1994). ",
        "conclusions": "An elongated jet in an HST image projecting from proplyd GMR 15 appears as a velocity `spike' in the corresponding pv array obtained with MES. Also, the only jet candidate to explain the high-speed unresolved feature in an MES pv array is found to be a displaced ionized knot in the HST imagery of the proplyd LV~5.  Two distinctly different dynamical explanations are considered for these jets. In many cases standard jet models (suitably modified for the conditions within the Orion Nebula) may be able to describe proplyd jets. However, we suggest that in some instances it may be more useful to model them as being due to the passage of discrete bullets of ejecta that are promptly ionized and that are ablated as they travel through the ambient medium. The trailing ablated material is then identified as the jet, with the bullet at its head. Such a dynamical mechanism has the attraction that the `Hubble-type' velocity law, if found in future observations of the proplyd jets, is naturally explained since material that was ablated earlier has slowed down more than material more recently incorporated into the flow. Kinematic differences between this model and a continous jet model could be investigated using STIS spectroscopy with its 0.1\\arcsec resolution."
    },
    "0207/astro-ph0207422_arXiv.txt": {
        "abstract": "We report on BeppoSAX and Chandra observations of three Hard X--Ray Transients in quiescence containing fast spinning ($P<5$ s) neutron stars: A~0538--66, 4U~0115+63 and V~0332+53. These observations allowed us to study these transients at the faintest flux levels thus far. Spectra are remarkably different from the ones obtained at luminosities a factor $>10$ higher, testifying that the quiescent emission mechanism is different. Pulsations were not detected in any of the sources, indicating that accretion of matter down to the neutron star surface has ceased. We conclude that the quiescent emission of the three X--ray transients likely originates from accretion onto the magnetospheric boundary in the propeller regime and/or from deep crustal heating resulting from pycnonuclear reactions during the outbursts. ",
        "introduction": "Young magnetic neutron stars orbiting a Be star companion occasionally show transient X--ray emission. Their spectra are relatively hard up to energies of tens of keV (power law with photon indexes $\\sim 1$) hence their name of Hard X--ray transients (HXRTs; White, Kaluzienski \\& Swank 1984). Be stars are characterised by equatorial mass loss episodes likely originating from their high (nearly break up) rotational velocities. When the neutron star along its (eccentric) orbit enters this equatorial disk an X--ray outburst episode is observed (Stella et al. 1986; Bildsten et al. 1997). Part of the material outflowing from the Be star is captured by the gravitational field of the neutron star, accretion onto its magnetic polar caps takes place, generating an intense pulsed X--ray flux. The spin periods of HXRTs range from 69 ms (A 0538--66; Skinner et al. 1982) to 25 min (RX J0146.9+6121; Mereghetti, Stella \\& De Nile 1993). Magnetic field strengths, when accurately inferred from cyclotron line features, lie in the range of $B\\sim (1-10)\\times 10^{12}$ G (Dal Fiume et al. 2000).  The powering mechanism and the accretion regime that pertains to the quiescent state of these sources is still uncertain. A change in the accretion regime should take place at lower luminosities in the case of fast spinning neutron stars as their magnetosphere expands beyond the corotation radius, therefore halting the infalling material at the magnetospheric boundary (i.e. the propeller regime). A contribution from matter leaking through the centrifugal barrier may still be present (Campana et al. 2001). A further emission mechanism is represented by cooling of the neutron star surface (deep crustal heating) due to pycnonuclear reactions occurring during outbursts (Brown, Bildsten \\& Rutledge 1998).  In this paper we present the first detailed investigation of the quiescent state of three of the fastest spinning (accreting) neutron stars in HXRTs. A 0538--66 and 4U 0115+63 were detected in our BeppoSAX observations while V 0332+53 remained undetected. A Chandra pointing revealed also the latter source (Section 2). We discuss these results in the light of the regimes experienced by a neutron star subject to a range of matter inflow. By using the neutron star parameters deduced from observations in outburst, we infer that these sources are likely in the propeller regime (Section 3). ",
        "conclusions": "We observed a sample of fast spinning neutron stars in HXRTs during quiescence with BeppoSAX and Chandra. The quiescent luminosities observed in the fast HXRTs of our sample are very low (especially in the case of V 0332+53, see Table 2) and one can ask if the inflowing matter can reach the neutron star surface. If accretion onto the neutron star surface took place in quiescence, then the mass inflow rate has to decrease by a large factor (up to $10^6$ in the case of V 0332+53) from outburst to quiescence, posing severe limitations to the Be wind and disk characteristics. A different way out is represented by the limited efficiency of the accretion process because matter is halted at the neutron star magnetosphere ($r_{\\rm m}$) when the magnetic field lines rotate locally at super-Keplerian speed. This process is often referred to as the centrifugal barrier or propeller mechanism (Illarionov \\& Sunyaev 1975; Stella et al. 1986). These systems have well known spin periods and magnetic field strengths thus allowing us to estimate the luminosity at which the centrifugal barrier starts operating. Using simple spherical accretion theory(which also provides a good approximation in the case of disk accretion, e.g. Wang 1995, 1996), one can work out the limiting mass inflow rate $\\mdot_{\\rm lim}$ and in turn the limiting accretion luminosity for the onset of the propeller: \\begin{eqnarray} L_{\\rm lim}(R)&=&G\\,M\\,\\mdot_{\\rm lim}/R \\\\ \\nonumber &\\simeq& 3.9\\times 10^{37} \\,\\xi^{7/2}\\,B_{12}^2\\,P_0^{-7/3}\\,M_{1.4}^{-2/3}\\,R_6^{5}\\ergs \\end{eqnarray} (where the neutron star magnetic field, spin period, mass and radius are scaled as $B=B_{12}\\,10^{12}$ G\\footnote{The magnetic field is obtained from the magnetic dipole moment $\\mu=B\\,R^3/2$.}, $P=P_0$ 1 s, $M=M_{1.4}\\,1.4\\msole$ and $R=R_6\\,10^6$ cm, respectively, e.g. Stella et al. 1986). $\\mdot$ indicates the mass accretion rate and $G$ is the gravitational constant. The factor $\\xi$ accounts for the deviations of $r_{\\rm m}$ as computed in spherical symmetry from the case of an accretion disk. In general $\\xi$ is in the range 0.5--1.5 (here we use $\\xi=1$). Values in the range of $\\sim 10^{34}-10^{37}\\ergs$ are derived for the onset of the centrifugal inhibition in the fast HXRTs of our sample (see below and Table 2).  For lower mass inflow rates than those in Eq. 1, the great majority of accreting matter can no longer reach the neutron star surface and a sharp drop off of the accretion luminosity is expected. The corresponding luminosity jump mainly depends on the spin period of the neutron star \\be \\Delta=\\bigl({{G\\,M\\,P^2}\\over {4\\,\\pi^2\\,R^3}}\\bigr)^{1/3} = 170\\, M_{1.4}^{1/3}\\,R_6^{-1}\\,P_0^{2/3} \\en (e.g. Corbet 1996; Campana \\& Stella 2000) and is a factor of 30--500 for the HXRTs in our sample (see Table 2). Therefore, the maximum accretion luminosity that can be emitted in the propeller regime is \\begin{eqnarray} L_{\\rm min}(r_{\\rm cor})&=&L_{\\rm min}(R)/\\Delta=G\\,M\\, \\mdot_{\\rm lim}/r_{\\rm cor}= \\\\ \\nonumber &=&2.4\\times 10^{35}\\,\\xi^{7/2}\\,B_{12}^2\\,P_0^{-3}\\,M_{1.4}^{-1}\\,R_6^{6}\\ergs \\end{eqnarray} Clearly these luminosities are all bolometric. While in the case of accretion onto the neutron star surface most of the emission goes into X--rays, in the propeller regime this is not clear and these numbers should be referred as upper limits. Moreover, the physics of the propeller regime is poorly understood and a fraction of matter may still leak through the barrier.  As can be noted from Table 2, the observed X--ray luminosities are all below the threshold for the onset of the centrifugal barrier and below the maximum expected luminosity in the propeller regime (this is true even for $\\xi=0.5$). This testifies that the HXRTs in our sample are all likely detected in the propeller regime and the observed luminosity derives from the mass inflow releasing its gravitational energy down to the magnetospheric radius.  The luminosity level pertaining to quiescence in the propeller regime clearly depends on the unknown quiescent mass inflow. An additional and independent luminosity can derive from the cooling of the neutron star made hot during the events of intense accretion: the inner crust compressed by the loaded material becomes the site of pycnonuclear reactions that may deposit enough heat into the core (Brown, Bildsten \\& Rutledge 1998; Colpi et al. 2001; see also Campana et al. 1998). In the last few years V 0332+53 and A 0538--66 did not show any outburst activity and therefore it is hard to estimate a mean accretion rate. This is instead possible for 4U 0115+63 which showed two strong outbursts and a number of small outbursts during the RXTE lifetime. Based on the observed outbursts one can derive a time-average rate of $\\sim 4\\times10^{15}\\gs$, resulting in a deep crustal heating luminosity of $4\\times10^{33}\\ergs$. This value has to be compared with the inferred black body luminosity of $\\sim 10^{33}\\ergs$, which is a factor $\\sim 4$ lower, this luminosity however could be hidden in the lower energy part of the spectrum. Similar luminosity levels (if not lower) apply to V 0332+53 and A 0538--66. Being the quiescent luminosity A 0538--66 $\\sim 5\\times 10^{35}\\ergs$ it cannot be supported by this emission mechanism only."
    },
    "0207/astro-ph0207108_arXiv.txt": {
        "abstract": "{ We derive age-metallicity relations (AMRs) and orbital parameters for the 1658 solar neighbourhood stars to which accurate distances are measured by the {\\it HIPPARCOS} satellite. The sample stars comprise 1382 thin disc stars, 229 thick disc stars, and 47 halo stars according to their orbital parameters. We find a considerable scatter for thin disc AMR along the one-zone Galactic chemical evolution (GCE) model. Orbits and metallicities of thin disc stars show now clear relation each other. The scatter along the AMR exists even if the stars with the same orbits are selected. We examine simple extension of one-zone GCE models which account for inhomogeneity in the effective yield and inhomogeneous star formation rate in the Galaxy. Both  extensions of one-zone GCE model cannot account for the scatter in age - [Fe/H] - [Ca/Fe] relation simultaneously. We conclude, therefore, that the scatter along the thin disc AMR is an essential feature in the formation and evolution of the Galaxy. The AMR for thick disc stars shows that the star formation terminated 8 Gyr ago in thick disc. As already reported by \\citet{Gratton_et.al.2000} and \\citet{Prochaska_et.al.2000}, thick disc stars are more Ca-rich than thin disc stars with the same [Fe/H]. We find that thick disc stars show a vertical abundance gradient. These three facts, the AMR, vertical gradient, and [Ca/Fe]-[Fe/H] relation, support monolithic collapse and/or accretion of satellite dwarf galaxy as thick disc formation scenario. ",
        "introduction": "The individual ages for solar neighbourhood stars are indispensable in the research of star formation history of the Galaxy. \\citet{Twarog1980} first derived the age-metallicity relation (AMR) for the disc in the neighbourhood of the Sun from {\\it ubvy} and H$\\beta$ photometry of a large sample of field stars. Theoretical isochrones used in the age determination were taken from \\citet{Ciardullo_Demarque1977}. In the Twarog's AMR, the metallicity increases from [Fe/H]$=-1.0$ at 13 Gyr to [Fe/H]$=-0.03$ at the age of the Sun. The mean metallicity has increased more slowly since then to a present value of [Fe/H]$=+0.01$ for the youngest stars. The dispersion in [Fe/H] is as small as $\\pm 0.1$ dex at any given age. \\citet{Carlberg_et.al.1985} used stellar models of \\citet{VandenBerg1983} and a revised metallicity calibration that takes into account a temperature dependence. Both ages and metallicities were estimated in a photometric manner, and the resulting AMR is qualitatively similar to that of \\citet{Twarog1980}, but [Fe/H] increases more gradually, showing an increase of only 0.3 dex over the past 15 Gyrs \\citep[cf.][]{Nissen_Schuster1991}. The metallicity dispersion decreases from $\\pm 0.15$ dex for the oldest stars ($13-20$ Gyrs) to $\\pm 0.05$ dex for younger stars.  \\citet{Edvardsson_et.al.1993} derived elemental abundances of O, Na, Mg, Al, Si, Ca, Ti, Fe, Ni, Y, Zr, Ba, and Nd for 189 nearby long-lived disc dwarfs by using high resolution, high S/N, spectroscopic data. Individual ages were derived photometrically from fits in the $\\log T_{\\rm eff} - \\log g$ plane of the isochrones by \\citet{VandenBerg1985}. The uncertainties in the relative ages are about $25 \\%$. Due to metallicity measurements of high precision, \\citet{Edvardsson_et.al.1993} improved greatly the AMR, but ironically the resulting AMR clearly indicated a considerable scatter ($\\sim 0.15$ dex) in the metallicities of disc stars formed at any given time, implying that there is only a very weak correlation between age and metallicity. The scatter seems to be substantially larger than that can be explained by observational errors. If the scatter is real, it would cause a serious difficulty for galactic chemical evolution (GCE) models, because it is easy to fit the average run of the data, but difficult to explain such a large scatter without breaking some of assumptions that GCE models usually make \\citep{Pagel_Tautvaisiene1995}. \\citet{Edvardsson_et.al.1993} suggested that the scatter arises from star formation stimulated from sporadic episodes of gas infall, although it is also possible that a different rate of chemical enrichment, depending on the distance from the Galactic centre, causes a scatter of this kind.    In this article, we derived the ages and orbital parameters for 1658 solar neighbourhood stars, almost ten times more than the previous researches.  Hence we succeed in finding out the new features of the Galaxy using stellar ages, chemical components, and orbits. The article is organised as follows. Section 2 derives the ages and the orbital parameters for sample stars. Section 3 shows the features relevant to thin disc stars while Sect. 4 shows those of thick disc. Section 5 discusses observational error in our data, abundance gradient, the abundance distribution functions, the formation of thick disc, and the scatter along the thin disc AMR. Section 6 concludes the present study. ",
        "conclusions": "\\subsection{Uncertainty in the AMR} \\subsubsection{Inhomogeneity of Spectroscopic [Fe/H] data}  Unfortunately, our AMR contains large errors in metallicity; typically $\\epsilon_{\\rm [Fe/H]}=0.15$ dex due to inhomogeneous data taken from different authors, while $\\epsilon_{\\rm [Fe/H]}=0.1$ dex in \\citet{Edvardsson_et.al.1993} due to their homogeneous [Fe/H] data reduced by the same analysis method. Even if stellar spectra were taken with high S/N and high resolution, [Fe/H] estimates by different authors could result in large scatter in [Fe/H].  Therefore, without examining carefully the details of individual analyses to understand the differences, it would be dangerous to argue too much details of the value of metallicity. We also have to keep in our mind that [Fe/H] determinations are affected by the adopted effective temperatures, gravities and microturbulent velocities, and that a stellar metal abundance can be in error, even if the observations are of excellent quality \\citep{CayrelCatalogue1997}.  However, recent observations are in good agreement for different observers. We have studied the most observed 30 stars in the [Fe/H] catalogue to find the average dispersion in metallicity is 0.13 dex (See. Table \\ref{var_table}). Our sample includes extremely metal-deficient stars, in which the dispersion in metallicity is apt to show the larger values. The dispersion in the metallicity for stars with [Fe/H] $>$ -2.5 is even smaller. Considering that the our sample includes a very small number of extremely metal-deficient stars, we conclude that the dispersion in the metallicity can be estimated to be 0.11 dex."
    },
    "0207/astro-ph0207278_arXiv.txt": {
        "abstract": "s{Kilohertz quasi-periodic oscillations (kHz QPOs) are probably caused by matter in Keplerian orbit at some preferred radius in the accretion disc around a compact star. In a given source, QPO frequencies can drift by a few hundred Hz following changes of the inner disc radius; but the disc cannot move closer to the star than the radius of the innermost stable circular orbit (ISCO) predicted by general realtivity, hence the kHz QPO frequencies must be limited by some maximum frequency.}  Long before kHz QPOs were discovered\\cite{iauc6319,iauc6320}, it had been already proposed$^{3-6}$ that evidence of the ISCO around neutron stars could be observed in flux variability studies of X-ray binaries: For some equations of state, the neutron star lies within the radius of the ISCO. Clumps of matter crossing that radius will no longer be rotationally supported and will fall extremely rapidly onto the neutron star surface; effectively, the accretion disc is terminated at that radius. There is a maximum Keplerian frequency around such a neutron star, corresponding to the minimum possible radius of the inner edge of the disc; variability of the X-ray flux produced in the disc at frequencies larger than $\\nu_{\\rm K} (r_{\\rm ISCO})$ should be strongly suppressed.  KHz QPOs models\\cite{mlp98a} suggest that the radius of the inner disc edge decreases as mass accretion rate, $\\dot M$, increases. Hence, when plotted against a quantity that measures $\\dot M$, QPO frequency should increase, but only until the inner disc edge reaches the ISCO; at that point QPO frequency should remain constant even if the $\\dot M$-related quantity keeps increasing\\cite{mlp98a,kfc97}.  This behavior may have been observed\\cite{zhang98} with the {\\em Rossi X-ray Timing Explorer (RXTE)}. Figure 1a shows, for the X-ray binary 4U 1820-30, the frequencies of both kHz QPOs vs.\\ X-ray intensity, which is commonly assumed to be a good measure of $\\dot M$. Frequencies increase more or less linearly with intensity up to $\\sim$2500 counts s$^{-1}$, and from then on they remain constant, even as intensity increases by $\\sim$30\\,\\%. \\begin{figure}[t] \\centerline{ \\epsfig{file=fig1a.ps, angle=-90, width=4.0cm} % \\epsfig{file=fig1b.ps, angle=-90, width=4.1cm} % \\epsfig{file=fig1c.ps, angle=-90, width=4.1cm} % } \\caption{KHz QPO frequencies vs. X-ray intenisty. {\\em Left panel:} Both kHz QPOs in 4U 1820--30 (data up to June 1998)$^{9}$. {\\em Middle panel:} Lower-frequency kHz QPOs in 4U 1608--52$^{10}$. {\\em Right panel:} Lower-frequency kHz QPOs in 4U 1820--30 (data up to June 2000)$^{11}$} \\end{figure} It is intriguing that such behavior is not observed in other sources with kHz QPOs\\cite{mendez00}. For instance, Figure 1b shows a similar plot for one of the kHz QPOs in 4U 1608--52. Interesting in this plot is the coexsitence of a good frequency-intensity correlation on timescales shorter than $\\sim$1 day (individual segments), with a lack of correlation on longer timescales. These long-term (uncorrelated) changes of frequency and intensity seen in 4U 1608--52 and other sources\\cite{mendez00}, could in principle produce a diagram similar to that shown in Figure 1a. Figure 1c shows QPO frequency vs.\\ X-ray intensity for the lower-frequency kHz QPO in 4U 1820--30, including new RXTE measurements: It is apparent that, as in 4U 1608--52, there are long-term uncorrelated variations of frequency and intensity in 4U 1820--30.  Source intensity may not be a good $\\dot M$ tracer, or QPO frequency may depend upon $\\dot M$ through the disc, with disc accretion not being a fixed fraction of total $\\dot M$\\cite{kaaret98,mendez99}. Because in several sources a one-to-one relation between QPO frequency and spectral properties has been observed\\cite{mendez00}, possible evidence for a QPO frequency saturation vs.\\ spectral related quantities in 4U 1820--30\\cite{kaaret99,bloser00} seemed to argue in favor of the ISCO interpretation for the maximum QPO frequency observed in this source. However, a careful analysis of the same observations shows that the evidence of the saturation is not so compelling, specially when some instrumental corrections, originally not applied, are taken into account\\cite{mendez01}. To fully resolve this issue an X-ray timing mission with $\\sim$10 times the area of RXTE\\cite{barret00} may be needed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207564_arXiv.txt": {
        "abstract": "{ We study the magneto-rotational instability of an incompressible flow which rotates with angular velocity $\\Omega(r)=a+b/r^2$ where $r$ is the radius and $a$ and $b$ are constants. We find that an applied magnetic field destabilises the flow, in agreement with the results of \\cite{rudiger01}. We extend the investigation in the region of parameter space which is Rayleigh stable. We also study the instability at values of magnetic Prandtl number which are much larger and smaller than R\\\"{u}diger \\& Zhang. Large magnetic Prandtl numbers are motivated by their possible relevance in the central region of galaxies (Kulsrud \\& Anderson 1992). In this regime we find that increasing the magnetic Prandtl number greatly enhances the instability; the stability boundary drops below the Rayleigh line and tends toward the solid body rotation line. Very small magnetic Prandtl numbers are motivated by the current MHD dynamo experiments performed using liquid sodium and gallium. Our finding in this regime confirms R\\\"{u}diger \\& Zhang's conjecture that the linear magneto-rotational instability and the nonlinear hydrodynamical instability (Richard \\& Zahn 1999) take place at Reynolds numbers of the same order of magnitude. ",
        "introduction": " ",
        "conclusions": "Our calculations show that many rotation laws of the form $\\Omega(r)=a+b/r^2$ which are hydrodynamically stable (that is to say, they satisfy the Rayleigh criterion) become linearly unstable when a magnetic field is applied. Our results confirm the finding of \\cite{rudiger01} and extend them in the Rayleigh stable region.  We have determined the instability at magnetic Prandtl numbers $\\Pm$ one order of magnitude smaller than R\\\"{u}diger \\& Zhang's, towards the small magnetic Prandtl number limit, which is relevant to possible MHD dynamo experiments with liquid sodium and gallium. Although the power law $\\Rey_{1c}\\propto\\Pm^{-0.5}$ that we find on the Rayleigh line ($\\mu=\\eta^2$) is slightly different from theirs ($\\Rey_{1c}\\propto\\Pm^{-0.65}$ on $\\mu=\\frac1{3}$), it confirms their conjecture that the nonlinear instability found by \\cite{richards99} and the MRI are likely to occur at Reynolds numbers of the same order of magnitude.  We also find that the flow becomes particularly unstable if the magnetic Prandtl number is greater than unity. The instability boundary in the $\\Rey_1$ vs $\\Rey_2$ plane rapidly tends towards the solid body rotation line. This enhanced instability for large $\\Pm$ is consistent with earlier results of \\cite{kurzweg63}. His boundary conditions were selected such as to avoid mathematical difficulties but for small $\\Pm$ agreed well with the % results of \\cite{chandrasekhar61}. The significance of the instability in this case is linked to the possibility (Kulsrud \\& Anderson 1992; Brandenburg 2001) that large values of $\\Pm$ exist in central regions of galaxies."
    },
    "0207/astro-ph0207087_arXiv.txt": {
        "abstract": "We present the results of a VLT observing program carried out in service mode using FORS1 on ANTU in Long Slit mode to determine the optical velocities of nearby low surface brightness galaxies.  As part of our program of service observations we obtained long-slit spectra of several members of the Phoenix dwarf galaxy from which we derive an optical helio-centric radial velocity of $-$13~$\\pm$~9~km/s.  This agrees very well with the velocity of the most promising of the HI clouds seen around Phoenix, which has a helio-centric velocity of $-$23 km/s, but is significantly different to the recently published optical heliocentric velocity of Phoenix of $-$52~$\\pm$~6~ km/s of Gallart {\\it et al.} (2001). ",
        "introduction": "Dynamical measurements of outlying Local Group galaxies are crucial for constraining both the age and the total mass of the Local Group and for probing the nature of the intrinsic dark matter within the galaxies. A necessary part of this process involves investigating the link between possible HI detections, the optical components of the galaxies and the relevance of the HI to the recent star formation history of the system. Optical velocities determine if observations of HI gas in and around these systems are the result of gas associated with these galaxies, a chance superposition with high velocity HI clouds, outlying components of the Magellanic Stream, or just Galactic foreground contamination.  Phoenix is a member of the Local Group, lying about 450~kpc from our Galaxy (Ortolani \\& Gratton 1988) in a fairly isolated location on the opposite side of the Galaxy from the Andromeda sub-system (see Table~1).  It is possibly one of the most distant of the satellites of our Galaxy, or perhaps a free-floating outlying Local Group object. It appears to be a galaxy in transition between a dwarf irregular (dI) and a dwarf spheroidal (dSph) system, having had both a recent burst of star formation and a plausible detection of $\\approx 10^5 \\Msun$ of HI.  The HI gas has been detected close to the position of Phoenix at several different locations, and velocities ($\\vsun =$ 56 km/s, 120 km/s Carignan, Demers \\& C\\^{o}t\\'{e} 1991; $\\vsun = -$23 km/s Oosterloo, Da Costa \\& Staveley-Smith 1996; and from mosaic mapping using the Australia Telescope Compact Array at $\\vsun = -$23, 7, 59, 140 km/s St-G\\'{e}rmain {\\it et al.} 1999 ).  Due to the superior resolution/sensitivity we adopt the latter measurements as defining the possible HI-associated gas throughout the remainder of this paper.  The HI complex at $\\vsun =$ 140~km/s is thought to be an outlying component of the Magellanic Stream, which passes close to the line-of-sight of Phoenix, while the component at $\\vsun = -$7~km/s is undoubtedly Galactic in origin. This leaves two remaining HI components which may be associated with Phoenix. The more compact of the HI clouds at $\\vsun = -$23~km/s is centred about 5 arc-minutes to the southwest of the optical centre of Phoenix and overlaps the optical image.  This is within a distance of about 650 pc if it lies at the same distance as Phoenix, and partially covers the optical image which is about 1 kpc in size.  The component with $\\vsun = $ 59~km/s is much more extended, fragments into four substructures, and is located significantly further to the south of the optical centre.  From the HI morphology and dynamics while it is unclear whether or not the 59~km/s component is, or perhaps was in the past, associated with Phoenix, the evidence for the $-$23~km/s system is more compelling.  The derived HI mass of $\\approx 10^5 \\Msun$ is comparable to that found in the Sculptor dwarf spheroidal and in the dwarf irregular/dwarf spheroidal transition object LGS 3 (Young \\& Lo 1997).   Furthermore, as St. G\\'{e}rmain {\\it et al.} point out, the HI velocity field of the $-$23~km/s component shows clear evidence of a velocity gradient indicating either rotation, or ejection, from Phoenix.  The resolved stellar population of Phoenix has been studied in some detail (Ortolani \\& Gratton 1988; van de Rydt, Demers, \\& Kunkel 1991; Held, Saviane \\& Momany 1999; Mart\\'inez-Delgado, Gallart \\& Aparicio 1999; Holtzmann, Smith \\& Grillmair 2000). Phoenix has obviously experienced recent star formation as can be seen from the sprinkling of blue stars across the field. They are concentrated in ``associations'' near the centre of the galaxy and elongated in the direction of the HI cloud at $-$23km/s. This recent ``episode'' of star formation has been quantified by Held {\\it et al.} (1999) to have started at least 0.6 Gyr ago, but accounts for less than 6\\% of the V-band luminosity of Phoenix and 0.2\\% of the mass. This is thus broadly consistent with the comprehensive modelling of the central region of Phoenix based on HST data by Holtzmann {\\it et al.} (2000). They found that star formation has been roughly continuous over the lifetime of Phoenix, with no obvious evidence for strongly episodic star formation, although a mildly varying star formation rate can fit the data.  Recently in an attempt to probe the optical--HI link further, Gallart {\\it et al.} (2001) published the first optical radial velocity study of a sample of 31 individual stars in Phoenix using VLT/FORS1 in MOS mode. Studying the blue end of the optical spectra in the range $\\approx$ 3500$-$6000 \\AA \\ , which encompasses the Mgb absorption complex, they determined a mean value of $\\vsun = -52\\pm6$ km/s.  There is quite a large offset between this optical measurement and the most likely HI gas velocity component at $-$23~km/s. Gallart {\\it et al.} interpret this difference as either caused by ejection of gas due to supernovae from the most recent burst of star formation in Phoenix about 100 Myr ago, or as a consequence of ram pressure stripping by a hot intergalactic medium within the Local Group.  The putative association of the HI gas with the optical galaxy is quite a crucial measurement, not only because of the direct coupling of the (potentially) recently expelled gas, but also because of the unique position of Phoenix in the Local Group as potentially (by far) the furthest outlying satellite of the Galaxy.  At a Galacto-centric distance of $\\approx$450 kpc, Phoenix could have unprecedented leverage in constraining the mass of the Galactic Halo out to 450~kpc.  As part of a long term VLT service programme to measure the radial velocities of several outlying Local Group satellites we recently acquired some service observations of Phoenix and decided that it was worth re-investigating the optical--HI connection with an additional optical velocity measurement. In this paper we therefore present the results of VLT/FORS1 long slit observations of Phoenix in the spectral region centred on the Ca~II near infra-red triplet. The slit position was aligned with 7 stars bright enough to derive radial velocities from. Although this is a much smaller number of stars than observed by Gallart {\\it et al.}, each spectrum is of sufficiently high S/N ($\\approx$10:1 per continuum \\AA) to derive an accurate velocity measurement. ",
        "conclusions": "Our determination of the optical velocity for Phoenix of $-$13km/s $\\pm$9km/s closely matches the most likely HI velocity for this galaxy of $-$23 km/s derived by St-G\\'{e}rmain {\\it et al.} (1999).  In a recent review of the impact of the VLT on Local Group Galaxies, Held (2001) quotes recent UVES measurements of giant stars in the Phoenix galaxy that agree to within a few km/s of the neutral gas velocity.  It is difficult to directly reconcile these much smaller negative optical radial velocities with the $-$52km/s $\\pm$6 km/s recently determined by Gallart {\\it et al.} (2001).  It is highly unlikely that Galactic foreground contamination could have affected any of the optical results significantly and the HI gas velocity result of $-$23km/s seems convincingly secure as argued by St-G\\'{e}rmain {\\it et al.} (1999).  There is also unlikely to be contamination of our results from the two main tidal streams enveloping our Galaxy: the Magellanic stream and the Sagittarius Dwarf tidal stream.  The Magellanic Stream is predominantly gas with no unambiguous detection of a stellar component yet reported.  Therefore, while the Magellanic Stream can contaminate the gas distribution in the Phoenix direction, it is unlikely to contribute to contamination of the optical velocities.  Likewise, pollution by the Sagittarius Dwarf tidal stream is also highly unlikely since the projection of its current orbit does not pass close to the line-of-sight to Phoenix.  The gradient in the distribution of young stars in Phoenix (e.g., Ortolani \\& Gratton 1988; Mart\\'inez-Delgado {\\it et al.} 1999) suggests that the recent star formation has been moving from east to west across the central component of Phoenix in a manner consistent with self-propagating star formation theories. The youngest stars are thus spatially overlapping the position of the HI cloud at $-$23km/s (as originally pointed out by Young \\& Lo 1997).  The relative position and velocity of this cloud combined with the evidence of recent star formation provides evidence that a burst of star formation can disrupt and potentially blow out gas from the centre of a dwarf galaxy, inhibiting further star formation (e.g., Dekel \\& Silk 1986; Mac Low \\& Ferrara 1999), although it is not the only possible explanation for what is seen (cf. ram pressure stripping scenarios as in Gallart {\\it et al.} 2001).  However, we can conclude that our results in conjunction with those of Held (2001) unequivocally show that modest amounts, ($\\approx 10^5 \\Msun$), of HI gas are associated with the Phoenix dwarf galaxy albeit somewhat offset from the optical centre of the galaxy."
    },
    "0207/astro-ph0207202_arXiv.txt": {
        "abstract": "Density profiles of cosmological virialized systems, or dark halos, have recently attracted much attention. I first present a brief historical review of numerical simulations to quantify the halo density profiles. Then I describe the latest results on the universal density profile and their observational confrontation. Finally I discuss a clustering model of those halos with particular emphasis on the cosmological light-cone effect. ",
        "introduction": "The key assumption underlying the standard picture of structure formation is that the luminous objects form in a gravitational potential of dark matter halos.  Therefore, a detailed description of halo density profiles as well as of their clustering properties is the most basic step toward constructing the formation and evolution of galaxies and clusters.  More specifically, the importance of the detailed studies of density profiles of dark halos is two-fold:  \\begin{description} \\item[(i) theoretical interest;]\\hfill\\par  What is the final (quasi-)equilibrium state of cosmological self-gravitating systems (as long as the energy dissipation is neglected) ?  One may easily think of two quite distinct, but equally plausible, possibilities; \\begin{enumerate} \\item[(A)] the systems reach a certain universal distribution which is independent of the cosmological initial condition. \\item[(B)] the systems somehow keep the memory of the cosmological initial condition even at the highly nonlinear regime. \\end{enumerate} The singular isothermal sphere: \\begin{eqnarray} \\rho(r) = \\frac{\\sigma^2}{2\\pi G} \\frac{1}{r^2} \\end{eqnarray} may be a reasonable possibility along the line of the idea (A).  As a matter of fact, quite often I meet people who argue that the idea (B) is unlikely because of the strong nonlinear nature of the gravitation. In such an occasion, I present an example of the well-known stable clustering solution for the nonlinear two-point correlation function (Davis \\& Peebles 1977): \\begin{eqnarray} P_{\\rm mass}(k) &\\propto& k^n \\quad \\rightarrow\\quad \\xi_{\\rm mass}(r) \\propto r^{-3(n+3)/(n+5)} . \\end{eqnarray} This provides a good specific case that the cosmological initial condition is not erased and imprinted even in the strongly nonlinear behavior as is well confirmed by later numerical simulations (e.g., Suto 1993; Suginohara et al. 2001).  Of course the answer to the final state of cosmological self-gravitating systems may not be unique since it should really depend on the specifics and scales of the systems under consideration. For instance, it is clearly hopeless to extract any meaningful cosmological information from the precise data on the orbit of the earth around the Sun; all the initial memory should have been lost due to the strongly nonlinear and chaotic nature of the gravitation.  Nevertheless the final state of dark matter halos corresponding to galaxy- and cluster-scales is a well-defined problem which may be reliably answered with the current high-resolution numerical simulations independently of the physical intuition.  \\item[(ii) practical application;]\\hfill\\par  Whether or not the density profiles of dark halos keep the cosmological initial memory, the quantitative (empirical) prediction of the profile for a given set of cosmological parameters has several profound astrophysical implications including the rotation curve of spiral galaxies, reconstruction of the mass distribution from the weak-lensing, and X-ray and SZ observations of galaxy clusters. In particular, the confrontation of those testable predictions against the accurate observational data may even challenge the cold dark matter paradigm itself. \\end{description}  In this article, first I will review the summary of the past studies of the density profiles of dark matter halos, and then present some applications of those results. ",
        "conclusions": ""
    },
    "0207/gr-qc0207073_arXiv.txt": {
        "abstract": "The Klein-Gordon-Einstein equations of classical real scalar fields have time-dependent solutions (periodic in time). We show that quantum real scalar fields can form {\\it non-oscillating (static) }solitonic objects, which are quite similar to the solutions describing boson stars formed with classical and quantum complex scalar fields (the latter will be studied in this paper). We numerically analyze the difference between them concerning the mass of boson stars. On the other hand, we suggest an interesting test (a viable process that the boson star may undergo in the early universe) for the formation of boson stars. That is, it is questioned that after a second-order phase transition (a simple toy model will be considered here), what is the fate of the boson star composed of quantum real scalar field. ",
        "introduction": "The presence of dark matter has been established indirectly in a wide range of scale of the universe, from that of individual galaxies to the entire universe itself \\cite{silk}. Though direct measurements of the nature of the dark matter have not yielded any result, speculations on its composition vary from baryonic to non-baryonic matter. Particles like axions and neutralinos are the specific targets for direct observation since the indirect measurements like rotation curves of spiral galaxies and others do not depend on the presence of a particular type of particles. One of the most promising candidates for dark matter is the boson star, which was discovered theoretically over thirty years ago \\cite{kaup}\\cite{ruffini}. Until now, the reality of boson stars has been successfully applied to various plausible physical situations \\cite{mielke}.  The boson star is a self-gravitating compact solitonic object made up of bosonic fields. Non-interacting complex scalar fields \\cite{kaup}\\cite% {ruffini} were originally considered for the constituents composing boson stars. In this case, the resultant configurations are typically `mini'-boson stars, which have small size and mass. This result originates from the following specific feature of boson stars; the boson star is protected from gravitational collapse by the Heisenberg uncertainty principle, instead of the Pauli exclusion principle that applies to fermionic stars, and the characteristic length scale of the former is much smaller than that of the latter. It has been shown that this situation can be dramatically changed by introducing self-interacting complex scalar fields. The self-interaction effectively generates a repulsive force and the maximum mass of stable boson stars can be enhanced up to a size of the order of ordinary fermionic stars % \\cite{colpi}\\footnote{% Analyltic evaluation of the maximum mass and higher order self-interaction effect to boson star configuration can be found in \\cite{jwho}.}.  On the other hand, there is another type of gravitationally bound solitonic object known as oscillating soliton star composed of classical real scalar fields \\cite{seidel}. In this case, the spacetime geometry and real scalar field satisfying the Klein-Gordon-Einstein (KGE) equations are time-dependent (periodic in time). Such objects are of interest to study dark matter, since the most promising candidates for dark matter are described by real scalar fields such as the axion \\cite{kim}. However, even though these solutions are stable with respect to some simple perturbations, it is still unclear whether the stability of the oscillating star can be maintained with general perturbations. Moreover, due to its oscillating feature, real scalar fields may not form primordial solitonic objects in the early universe \\cite{seidel2} (cf. \\cite{kolb}), which are expected to play important roles in galaxy formation, the microwave background, and formation of protostars.  In this paper, we report some very interesting results for the self-gravitating solitonic object formed with real scalar fields: there exist gravitationally bound {\\it non-oscillating (static) }objects composed of {\\it quantum real }scalar fields, instead of {\\it classical real} ones. Our (zero-node) solitonic solution is quite similar to the solutions describing boson stars formed with classical \\cite{colpi} and quantum complex scalar fields (the latter will be studied in this paper). The only difference between them at the equation of motion level is just the effective coefficient of the self-interaction term (within the validity considered in this paper, i.e. zero-node solutions, the semi-classical and Hartree (mean-field) approximations). Therefore, non-self-interacting quantum/classical complex and quantum real scalar fields coupled to gravity can form identical static mini-boson stars.  On the other hand, as many viable cosmological scenarios indicate, quantum fields present in early stage of the universe may experience phase transitions by temperature changes. Thus, if the boson star forms in the early universe, it would undergo phase transitions. It might be an interesting issue to consider the consequence of such a phase transition for the boson star.\\footnote{In \\cite{torres}, a boson star in an evolving cosmological background with time varying gravitational constant (i.e., in the context of scalar-tensor theory of gravity) has been considered.} What is the fate of the scalar field and what are the final products after the phase transition, which is presumably second order? In this paper, as a first step, we consider a simple toy model; without speculating on the detailed procedure of the phase transition, we introduce a plausible form of the Lagrangian for real scalar fields coupled to gravity, which would give the physics after phase transition, and examine that there may still exist boson stars. If they continue to exist, what are the differences between the boson stars before and after phase transition.  In the next section, we start by studying boson stars composed of {\\it % quantum} complex scalar fields, and compare them with the boson stars formed with the {\\it classical }complex scalar fields \\cite{colpi}. The boson stars composed of quantum real scalar fields are studied in Sect.3. In Sect.4, we consider boson stars composed of quantum real scalar fields that undergo a second-order phase transition. Disscussions and some remarks are given in Sect.5. ",
        "conclusions": "In this paper, we have considered self-gravitating solitonic objects made up of quantum complex/real scalar fields. It has been shown that quantum complex/real scalar fields may compose the boson stars that are similar to that formed by classical complex scalar fields. The difference in the theories considered here appears only in the coefficient of the self-interaction term in the Einstein equations, i.e. they are $\\Lambda /2$ for classical complex fields, $5\\Lambda /8$ for quantum complex ones, and $% 3\\Lambda /4$ for quantum real ones. Numerically we have verified, Fig.1 and Fig.2, that the maximum mass of the boson stars increases with the magnitude of the coefficient of the self-interaction term. This result can be understood as follows; after eliminating the time-dependent part the KGE equations for the quantum complex/real scalar fields become effectively for two ``real'' fields as (\\ref{eqmx}) and (\\ref{eeq}). Then, the equations include interactions between the two ``real'' fields as well as self-interactions of each fields. On the other hand, the KGE equations for the classical complex fields considered in Ref.\\cite{colpi} become effectively for one ``real'' field after eliminating the time-dependent part and contain just the self-interaction term of the ``real'' field. Therefore, the coefficients of the resulting interaction terms in the KGE equations for the quantum complex/real scalar fields, which are given by (\\ref{eqm1}), (% \\ref{eqm2}), and (\\ref{eeq1}), (\\ref{eq2}), are greater than that in the case of the classical complex scalar fields. According to the argument that more interaction energy generates effectively a larger maximum mass of the boson star, the maximum masses of the boson stars composed of quantum complex/real scalar fields are greater than that in the case of classical complex scalar fields. This argument, however, has a limitation due to the presence of a crossing between the mass-curves in Fig.1 and Fig.2.  In fact, the existence of such solitonic solutions does not guarantee that the bosonic compact objects can be formed in the universe. (cf. \\cite{lopez}) It has been shown that there is a process that is able to describe the formation of the bosonic compact objects in the early universe \\cite{seidel2}. The process, which is similar to the way of describing the settling of collisionless star systems, starts with collapsing due to a gravitational instability analogous to the Jeans instability. Then, after undergoing the so called gravitational cooling mechanism ejecting part of the scalar field, bosonic compact objects could be formed from the primodial bosonic cloud. In the case of complex scalar fields, this mechanism works to form mini-boson stars. However, in general, the oscillatons made up of classical real scalar fields can be formed only in a short dynamical time scale, but are unstable in such a state. Thus, without introducing additional proviso such as fragmentation of the primodial bosonic cloud, the classical real scalar fields would be ruled out as a candidate for the dark matter.  In this paper, we have shown that quantum real scalar fields can form the {\\it boson star} rather than the {\\it oscillaton}. Especially, without the self-interaction $\\Lambda =0$, the mini-boson star composed of the quantum real scalar fields becomes exactly identical to that made up of the classical (and quantum) complex scalar fields. Thus, considering the quantum effect, the real scalar field can be saved, and can be the most promising candidate for dark matter.  As a corollary, which would be an interesting issue in the study of bosonic objects, the quantum complex/real scalar fields in excited states may be considered. The solitonic solutions, if they exist, are unlikely to be static in both (complex and real) cases. However, it is not obvious if such quantum fields in excited states can be ruled out completely as candidates for the dark matter.  Another interesting test for the boson star composed of the quantum real scalar fields is the second-order phase transition. In this analysis, we have shown that the phase transition makes the fields effectively more massive, $m\\longrightarrow \\sqrt{2}m$, or equivalently less self-interactive, $\\Lambda \\rightarrow \\Lambda /2$, and more importantly the boson stars can exist even after the phase transition.  \\begin{acknowledgement}  We would like to thank B. Teshima for his prominent contribution to the numerical job, and S. P. Kim and S. Sengupta for their helpful discussions and comments. J.H. and C.H.L. also thank to D. Page and V. Frolov for their warm hospitality during staying at Univ. of Alberta. This work was supported in part by Korea Science and Engineering Foundation under Grant no. 1999-2-112-003-5 (J.H. and C.H.L.), the BK21 project at Hanyang University (C.H.L.), and the Natural Science and Engineering Research Council of Canada(J.H., C.H.L. and F.C.K.).  \\end{acknowledgement}"
    },
    "0207/astro-ph0207344_arXiv.txt": {
        "abstract": "Observations of W51 with {\\it Submillimeter Wave Astronomy Satellite} (SWAS) have yielded the first detection of water vapor in a diffuse molecular cloud.  The water vapor lies in a foreground cloud that gives rise to an absorption feature at an LSR velocity of 6 km s$^{-1}$.  The inferred water column density is $2.5 \\times 10^{13} \\, \\rm cm^{-2}$. Observations with the Arecibo radio telescope of hydroxyl molecules at ten positions in W51 imply an OH column density of $8 \\times 10^{13} \\, \\rm cm^{-2}$ in the same diffuse cloud.  The observed H$_2$O/OH ratio of $\\sim 0.3$ is significantly larger than an upper limit derived previously from ultraviolet observations of the similar diffuse molecular cloud lying in front of HD 154368.  The observed variation in H$_2$O/OH likely points to the presence in one or both of these clouds of a warm ($T \\simgt 400$) gas component in which neutral-neutral reactions are important sources of OH and/or H$_2$O. ",
        "introduction": "Since its launch in December 1998, the {\\it Submillimeter Wave Astronomy Satellite} (SWAS; Melnick et al.\\ 2000) has detected water vapor in more than 70 molecular clouds by means of observations the $1_{10} - 1_{01}$ transition of ortho-H$_2$O (e.g. Snell et al.\\ 2000; Ashby et al.\\ 2000; Neufeld et al.\\ 2000a). While emission-line observations form the core of the SWAS program on interstellar water vapor,  absorption-line observations are possible toward a few bright continuum sources; these include Sagittarius B2 (Neufeld et al.\\ 2000b; hereafter N00), Sagittarius A, W49, and W51.  Absorption-line observations typically probe the water vapor abundance in several kinematically-distinct foreground clouds lying along the line-of-sight to the source (e.g. N00). Under typical interstellar conditions, most water molecules are in the lower state of the $1_{10} - 1_{01}$ transition (i.e. in the ground state of ortho-water); thus absorption line observations have the distinct advantage of yielding water vapor column densities that are insensitive both to the physical conditions in the absorbing cloud and to the assumed rate coefficients for collisional excitation of water.  Typically, the H$_2^{16}$O absorption line is very optically-thick, yielding only a lower limit on the water column density, but observations of {\\it optically-thin} absorption by the H$_2^{18}$O isotopologue (less abundant by a factor of 250--500) have led to a quantitative determination of the water column densities in foreground clouds along the Sgr B2 sight-line (N00).  In this {\\it Letter}, we report the results of similar observations carried out towards the star-forming region W51.  The results are  particularly intriguing, because they provide our first detection of a H$_2^{16}$O absorption line which is of only {\\it moderate} optical depth. This feature, observed at an LSR velocity $\\sim 6$~km~s$^{-1}$, originates in a diffuse foreground cloud in the which the water column density is small.  This cloud has previously been detected by means of 21~cm absorption-line observations (Koo 1997): the inferred HI column density is $\\sim 10^{21}\\,\\rm cm^{-2}$.  Spaans et al. (1998; hereafter S98) have argued that measurements of OH and H$_2$O column densities in diffuse clouds provide a valuable probe of the chemistry of interstellar oxygen molecules; in particular, the OH/H$_2$O abundance ratio serves to constrain the branching ratio for the dissociative recombination of the molecular ion H$_3$O$^+$, a crucial parameter in chemical models for both diffuse and dense molecular clouds. Accordingly, we have used the Arecibo Observatory (AO) to carry out OH absorption line observations toward the same source, the results of which are also presented here.  The observations and data reduction are described in \\S 2 below, and the observational results presented in \\S 3.  In \\S 4 we discuss the derived water and OH column densities and the constraints that they place upon the oxygen chemistry in molecular clouds. ",
        "conclusions": "Because of the relative simplicity of the chemical networks involved, diffuse molecular clouds provide a useful laboratory for testing astrochemical models; in particular, the H$_2$O/OH abundance ratio serves as a valuable probe of the chemical network that produces oxygen-bearing molecules (S98). One key uncertainty in that network concerns the dissociative recombination of H$_3$O$^+$ with electrons, and specifically the fraction of such recombinations that produce water, $f_{\\rm H_2O}$, a quantity for which two laboratory groups have obtained highly discrepant results. According to results obtained in the flowing afterglow experiment of Williams et al.\\ (1996), a fraction $f_{\\rm OH}=0.65$ of dissociative recombinations of H$_3$O$^+$ lead to OH, a fraction $f_{\\rm H_2O} = 0.05 $ to H$_2$O, and the remaining fraction $f_{\\rm O}= 1 - f_{\\rm OH} - f_{\\rm H_2O}$ to O. A different experimental technique (Vejby-Christensen et al.\\ 1997), which made use of the ASTRID heavy-ion storage ring in Denmark, yielded significantly different results (Jensen et al. 2000), {\\it viz.} $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O} = 0.74 \\pm 0.02 : 0.25 \\pm 0.01 : 0.013 \\pm 0.005$. Similar results (although with larger error bars) were obtained (Neau et al.\\ 2000) from the CRYRING heavy-ion storage ring facility; they were $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O} = 0.78 \\pm 0.08 : 0.18 \\pm 0.05 : 0.04 \\pm 0.06$,  Taken together with the ground-based observations of OH that we obtained at Arecibo, the SWAS observations of W51 imply an H$_2$O/OH abundance ratio $\\sim 0.3$ in the diffuse cloud that is responsible for the $v_{\\rm LSR} = 6 \\rm \\, km \\, s^{-1}$ feature. Considering the uncertainties in our determination of the H$_2$O and (particularly) the OH column densities, we estimate the H$_2$O/OH ratio to be uncertain by a factor $\\sim 2$. In comparing the observed H$_2$O/OH abundance ratio with theoretical predictions, we have used the steady-state photodissociation region (PDR) model of Kaufman et al.\\ (1999), modified so as to treat the case of a finite slab illuminated from two sides.  Because H$_2$ and CO are photodissociated following line absorption, their photodissociation rates are reduced by self-shielding.  In order to treat correctly the effects of self-shielding for radiation incident upon {\\it both} sides of the slab, the H$_2$ and CO abundances must be obtained by an iterative method.  \\subsection{Standard models of cold diffuse clouds}  In Figure 4, we show the predicted H$_2$O and OH column densities for a variety of {\\it astrophysical} parameters: the total visual extinction through the cloud, $A_V$, in magnitudes; the illuminating UV field, $G_0$, in units of the Habing field; and the cosmic ray ionization rate, $\\zeta_{\\rm cr}$. All results apply to an assumed cloud density, $n_H$, of 100 H nucleons per cm$^{3}$.  The temperature is calculated from considerations of thermal balance and is $\\sim 30$~K at the cloud center. Filled squares apply to models with $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O}=0.75 : 0.25 :  0.0$ (values suggested by the ASTRID storage ring experiment), while filled triangles apply to models with $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O}= 0.65 : 0.05 :  0.30$ (suggested by the flowing afterglow experiment). The different astrophysical parameters for each plotted data point are described in the figure caption.  Black circles represent the column densities observed toward W51 and HD 154368.  A striking feature of Figure 4 is that although the H$_2$O and OH column densities depend strongly on the assumed astrophysical parameters, their {\\it ratio} is determined primarily by the assumed branching ratio $f_{\\rm H_2O}$ and shows almost no dependence upon $A_V$, $G_0$, or $\\zeta_{\\rm cr}$.  This behavior can be understood by means of a simple ``toy\" model, in which we assume OH and H$_2$O to be formed by dissociative recombination of H$_3$O$^+$ and destroyed by photodissociation. The expected H$_2$O/OH ratio is given by \\begin{equation} {n({\\rm H_2O}) \\over n(\\rm OH)} = {\\zeta_{\\rm OH} \\over \\zeta_{\\rm H2O}} \\times {f_{\\rm H_2O} \\over f_{\\rm OH} + f_{\\rm H_2O}} \\sim 0.69 \\times {f_{\\rm H_2O} \\over f_{\\rm OH} + f_{\\rm H_2O}} \\end{equation} where $\\zeta_{\\rm OH} = 3.5 \\times 10^{-10} G_0 \\exp(-1.7\\,A_V) \\,\\rm s^{-1}$ and $\\zeta_{\\rm H2O} = 5.1 \\times 10^{-10} G_0 \\exp(-1.8\\,A_V) \\,\\rm s^{-1}$ are the assumed photodissociation rates for OH and H$_2$O (Roberge et al.\\ 1991).\\footnote{The quantity $f_{\\rm H_2O}$ appears with $f_{\\rm OH}$ in the denominator of the second term on the right-hand-side, because photodestruction of H$_2$O results in the formation of OH.}  Equation (2) yields H$_2$O/OH abundance ratios of 0.172 and 0.049 respectively for the branching ratios assumed for the filled squares and triangles in Figure 4. These ratios are shown by dashed lines in Figure 4, and do indeed yield good agreement with results from the full steady-state PDR model. The assumption of chemical steady-state is justified by the fact that the photodissociation timescale for OH is only $\\sim 100 \\exp(1.7\\,A_V) / G_0$ years.   Given standard models for cold diffuse clouds, the observed H$_2$O/OH abundance ratio of $\\sim 0.3$ in W51 is {\\it consistent} with the case $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O} = 0.75 : 0.25 : 0$ and clearly {\\it inconsistent} with the case $f_{\\rm OH} : f_{\\rm H_2O}: f_{\\rm O} = 0.65 : 0.05: 0.30$. Thus -- if interpreted using standard models for cold diffuse clouds -- the observed H$_2$O/OH abundance ratio argues for the laboratory results obtained in the ASTRID storage ring experiment and against those obtained in the flowing afterglow experiment.  This conclusion, however, is different from that obtained by S98, who used ultraviolet absorption line observations with the Goddard High Resolution Spectrometer of HST to place an upper limit of only 0.06 ($3\\,\\sigma$) on the H$_2$O/OH ratio in an entirely different diffuse cloud, which lies in front of the star HD 154368.  Based upon these observations, S98 argued for a {\\it low} value of $f_{\\rm H_2O}$ that is inconsistent with the ASTRID storage ring experiment.  The puzzling discrepancy between the observed H$_2$O/OH abundance ratio in these two diffuse clouds may point to the importance of additional production mechanisms for OH or H$_2$O that do not involve the dissociative recombination of H$_3$O$^+$.  This possibility is addressed in \\S4.2 below.  \\subsection{Enhanced-temperature models of diffuse molecular clouds}  It has long been recognized (e.g. Elitzur \\& de Jong 1978) that neutral-neutral reactions provide an alternate production route to OH and H$_2$O.  The reactions \\begin{equation} \\rm O + H_2  \\rightarrow OH + H \\\\ \\end{equation} \\begin{equation} \\rm OH + H_2  \\rightarrow H_2O + H \\\\ \\end{equation} possess activation barriers that make them negligibly slow at the low temperatures typical of diffuse interstellar clouds;  at temperatures above $\\sim 300$~K, however, they become important production mechanisms for OH and H$_2$O (Neufeld, Lepp \\& Melnick 1995). If even a small fraction of the gas in the W51 $\\rm 6\\, km\\, s^{-1}$ and/or the HD 154368 cloud were sufficiently warm -- as a result of a weak shock, for example -- then neutral-neutral reactions might perturb the OH and H$_2$O column densities.  To investigate this possibility, we have obtained model predictions for PDRs in which the temperature has been fixed at a variety of temperatures between 100 and 1500~K.  The results are represented by the magenta locus in Figure 4; they were obtained for the astrophysical parameters adopted by S98 for the HD 154368 cloud -- $A_V = 2.65$~mag, $n_H = 325\\,\\rm cm^{-3}$, $G_0 = 3$ -- but for a branching ratio $f_{\\rm H_2O}  = 0.25$ and a cosmic ray ionization rate of $1.8 \\times 10^{-17} \\rm cm^{-3}$.  The OH and H$_2$O column densities are both clearly enhanced by neutral-neutral reactions.  At moderate temperatures, the H$_2$O/OH ratio decreases because the large O/OH ratio makes reaction (3) more important than (4).  At temperatures higher than $\\sim 600$~K, however, the effect on the ratio is reversed, and the H$_2$O/OH ratio is {\\it increased } by neutral-neutral reactions.  The results obtained in enhanced-temperature models suggest a way out of the puzzle posed by the discrepant OH/H$_2$O ratios measured in the W51 $\\rm 6\\, km\\, s^{-1}$ and HD 154368 clouds.  If these clouds possess small (but differing) amounts of warm gas, then the OH/H$_2$O abundance ratios could differ (even though the value of $f_{\\rm H_2O} $ must, of course, be identical in both clouds). We are unaware of any observations that rule out the presence of small amounts of warm gas in these sources; indeed, the presence of such gas in diffuse clouds is predicted by certain models that seek to explain the anomalously-high CH$^+$ abundances observed in many diffuse clouds as resulting from the effects of turbulence or weak shocks\\footnote{Depending upon the geometry, velocity shifts of the OH and H$_2$O lines relative to the lines of other species (e.g. HI) might be an observational signature of a shock production mechanism.  The velocity shifts can be very small, however, if the shock propagates at an oblique angle to the line-of-sight or if multiple shocks are present in the beam.  Thus, the absence of any such signature in the W51 6 km $s^{-1}$ cloud does not argue strongly against the production of OH and H$_2$O in shocks.} (e.g. Joulain et al.\\ 1998, Flower \\& Pineau des Forets\\ 1998, and references therein).  Unfortunately -- as discussed above -- the effect of warm gas upon the H$_2$O/OH ratio depends critically upon the gas temperature -- and even switches sign at $T\\sim 600$~K.  Thus, if warm gas is present, the observed H$_2$O/OH ratio cannot even be used to derive a limit upon $f_{\\rm H_2O} $.  For example, if $\\sim 5\\%$ of the gas in HD 154368 were at $\\sim 500$~K, then the observations of OH and H$_2$O in that source could be reconciled with the large branching ratio $f_{\\rm H_2O}  = 0.25$ derived in \\S4.1 above.  Alternatively, if $\\sim 0.3\\%$ of the gas in the W51 cloud were at $T\\sim 700$~K, then the observed H$_2$O/OH ratio would be consistent with the lower limit on $f_{\\rm H_2O} $ inferred previously by S98.  To summarize, the discrepancy between the H$_2$O/OH ratio reported here for the W51 6 km s$^{-1}$ cloud and that reported previously for the HD 154368 cloud (S98) suggests that a component of warm ($T \\simgt 400$) gas is present in one or both of these sources.  The presence of this warm component makes it difficult to determine observationally the branching ratio for the dissociative recombination of H$_3$O$^+$ with electrons to form OH and H$_2$O. Our new observations of W51 cast doubt upon the conclusion of S98 that the branching ratio to H$_2$O is small, but do not allow the branching ratio to be determined definitively.  This work was supported NASA's SWAS contract NAS5-30702. We gratefully acknowledge the excellent support of the telescope operators at the Arecibo Observatory. The Arecibo Observatory is operated by the National Astronomy and Ionosphere Center under a Cooperative Agreement with the National Science Foundation.  \\clearpage"
    },
    "0207/astro-ph0207458_arXiv.txt": {
        "abstract": "We present a proposal for a gravitational wave detector, based on the excitation of an electromagnetic mode in a resonance cavity. The mode is excited due to the interaction between a large amplitude electromagnetic mode and a quasi-monochromatic gravitational wave. The minimum metric perturbation needed for detection is estimated to the order $7\\times 10^{-23}$ using current data on superconducting niobium cavities. Using this value together with different standard models predicting the occurrence of merging neutron star or black hole binaries, the corresponding detection rate is estimated to 1--20 events per year, with a `table top' cavity of a few meters length. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207172_arXiv.txt": {
        "abstract": "Evidence for a second main-sequence turn-off in a deep colour-magnitude diagram of NGC 1868 is presented. The data were obtained with HST/WFPC2 and reach down to $m_{555} \\simeq 25$. Besides the usual $\\tau \\simeq 0.8~Gyr$ turn-off found in previous analyses, another possible turn-off is seen at $m_{555} \\simeq 21$ ($M_{V} \\simeq 2.5$), which is consistent with an age of $\\tau \\simeq 3~Gyrs$. This CMD feature stands out clearly especially when contaminating field LMC stars are statistically removed. The background subtracted CMD also visibly displays a red subgiant branch extending about 1.5 mag below the younger turn-off and the clump of red giants. The significance of the secondary turn-off in NGC 1868 was confirmed with Monte-Carlo simulations and bootstrapping techniques. Star-counts in selected regions in the cluster CMD indicate a mass ratio of old population/young population in the range $5\\%$ \\ltsima $M_{old} / M_{young}$ \\ltsima $12\\%$, depending on the mass function slope. The existence of such a subpopulation in NGC 1868 is significant even in the presence of uncertainties in background subtraction. The possibility that the secondary turn-off is associated with the field star population was examined by searching for similar features in CMDs of field stars. Statistically significant excesses of stars redwards of the main-sequence were found in all such fields in the range $20$ \\ltsima $m_{555}$ \\ltsima $22$. These however are much broader features that do not resemble the main-sequence termination of a single population. We also discuss other alternative explanations for the feature at $m_{555} \\simeq 21$, such as unresolved binarism, peculiar stars or CMD discontinuities associated with the B\\\"ohm-Vitense gap. ",
        "introduction": "NGC 1868 is a rich Large Magellanic Cloud (LMC) cluster located at $\\alpha = 5^h~14^m$ and $\\delta = -63^o~57'$ (J2000), approximately 6$^\\circ$ away from the LMC's centre. Its age has been estimated both from ground-based photometry and spectroscopy, often yielding discrepant results covering the $3~10^8~yrs$ \\ltsima $\\tau$ \\ltsima $10^9~yrs$ range (Flower et al. 1980, Hodge 1983, Bica \\& Alloin 1986, Chiosi et al. 1986, Olszewski et al. 1991, Corsi et al. 1994 and references therein). Most age estimates come from the interpretation of optical colour-magnitude diagrams (CMDs), usually based on the positions of either the main-sequence turn-off (MSTO) or the clump of red giants (RC) or both. However, none of the CMDs available until now have been deep enough to allow probing the main sequence at magnitudes as faint as $V \\simeq 22$ with small photometric errors (see the web page on www.ast.cam.ac.uk/STELLARPOPS/LMCdatabase for a detailed list of references on NGC 1868 and other rich LMC clusters in the HST Cycle 7 program GO7307).  In this paper we investigate a possible second, fainter and therefore older, MSTO at $V \\simeq 21$, based on a deep CMD of NGC 1868 built from HST/WFPC2 data. This second population of stars may be the result of a strong interaction with another cluster, or of the capture of a lower-mass, older cluster by NGC 1868. In fact, there is evidence for a growing number of clusters in the LMC which may have undergone strong interactions or mergers (Kontizas et al. 1993, Dieball \\& Grebel 1998, Leon et al. 1999, Dieball, Grebel \\& Theis 2000). Confirmation of a merger event in a cluster's history, however, requires detailed photometric and/or spectroscopic data. Sagar et al (1991), based on a CMD of NGC 2214 with two apparent supergiant branches, suggested that it was made up of two distinct populations. However, Lee (1992), Bhatia \\& Piotto (1994) and Banks et al (1995), based on larger and more accurate photometry, found a single population in the CMD of NGC 2214. Given that signatures of mergers or strong interactions remain observable in clusters for at least 1 Gyr (de Oliveira, Bica \\& Dottori 2000), NGC 1868, if confirmed as a merger product, may be a good laboratory for studying the dynamical effects and the final products of such events. Hence, it is essential to investigate closely the possibility that it is made up of two distinct populations.  Other explanations for a feature similar to a MSTO exist. Unresolved binaries, for instance, are usually brighter and redder than the main sequence of single stars. Therefore, an enhanced fraction of unresolved binaries could result in such a feature. Sudden changes in stellar structure, such as the onset of convective envelopes, may lead to CMD features such as the B\\\"ohm-Vitense gap, that may also mimic a MSTO (B\\\"ohm-Vitense 1970). These possibilities are also considered in the analysis of our deep NGC 1868 CMD, as well as of other similar CMDs.  The paper outline is as follows: in \\S 2 we describe the HST/WFPC2 data used, both on NGC 1868 and on control field areas. We also discuss the process of removing contamination by field stars from the on-cluster CMD. In \\S 3 we test the statistical significance of the candidate secondary MSTO by means of bootstrapping realizations. In \\S 4 we use isochrones and star counts in selected CMD areas to extract useful information about this presumed secondary population in NGC 1868. In \\S 5 we explore the possibility that the candidate MSTO is related to field stars, making use of CMDs obtained in control field areas. We also explore the alternative mechanisms that could yield to features similar to a MSTO at that position. Our main conclusions are presented in \\S 6.   \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_cmdcomp.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{ On-cluster (left panel) and off-cluster (right panel) colour-magnitude diagrams. The large trapezium box indicates the locus of evolved stars, whose field star subtraction was carried out separately from that of MS stars. Empirically determined photometric uncertainties are shown on the left panel.} \\end{figure*}  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_cmdclean.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Colour-magnitude diagram resulting from subtracting field stars from the on cluster data. } \\end{figure*}  \\section[]{The data}  The WFPC2 on-cluster data are described in more detail in Santiago et al. (2001). The off-cluster data used for field star subtraction were reduced and analyzed by Castro et al. (2001). In both cases the images were combined and calibrated using the standard pipeline procedure (Holtzman et al. 1995a,b). Photometry was carried out using DAOPHOT tasks and is described in detail by Santiago et al (2001) and Castro et al (2001). Zero-points, CTE and aperture corrections were also applied, again following standard procedures (Santiago et al 2001, de Grijs et al 2002a,b).  Figure~1 shows the resulting CMDs. The left panel shows stars belonging to the on-cluster images. Two such images were taken, one with the Planetary Camera (PC) at the centre of NGC 1868 and the other at its half-light radius. They are the CEN and HALF fields as defined by Santiago et al. (2001). The CMD shown represents the final sample, including stars from both HALF and CEN fields and with no repeats. The total on-cluster solid angle is approximately 7.5 arcmin. As explained in Santiago et al. (2001), the stars in the region common to the HALF and CEN fields had two independent photometric measurements and were thus used to determine uncertainties as a function of $m_{555}$ magnitudes and $m_{555} - m_{814}$ colour. A total of 731 stars were found in the overlap region between the HALF and CEN fields. These empirical error determinations are an essential part of our upcoming analysis. The uncertainties were computed in bins 0.5 mag wide in both $m_{555}$ and $m_{814}$. In each magnitude bin we computed the mean value and standard deviation ($\\sigma$) of the distribution of magnitude differences. The typical uncertainty of a single measurement was then assumed to be $\\sigma / \\sqrt{2}$. The photometric uncertainty analysis is described in more detail in Kerber et al (2002).  The right panel in Figure 1 shows stars in the off-cluster area, located 7.3' away from the centre of NGC 1868 and previously studied by Castro et al (2001). This off-cluster field corresponds to a single WFPC2 field, covering about 5 arcmin. Both CMDs include only objects that were classified as stellar sources in the classification schemes presented in Santiago et al (2001) and Castro et al (2001). No additional cleaning of remaining spurious objects or background galaxies was made, resulting in some objects located well away from the cluster main-sequence (MS), the red giant branch (RGB) and the RC. Most of them should be contaminating unresolved background galaxies or faint stars in our Galaxy.  The on-cluster stars display a clear MS that terminates at $m_{555} \\simeq 19.2$. Previous works find the MS termination at $V = 19.2-19.3$ (Corsi et al 1994, Brocato et al 2001). This is also roughly the magnitude of the RC ($m_{555} - m_{814} \\simeq 0.8$), which contains He burning stars and is largely dominated by the cluster stars. In both panels, a clear RGB stretches downwards from the RC position to the bottom of the subgiant branch (SGB) at $m_{555} \\simeq 22$, $m_{555} - m_{814} \\simeq 0.8$. These stars mostly belong to the old ($\\tau$ \\gtsima $10~Gyrs$) LMC field star population (Holtzman et al. 1997, Castro et al. 2001, Smecker-Hane et al 2002). Besides the main cluster and old field turn-offs, there are two additional features that may be interpreted as MSTOs in the CMD: one has mostly field stars with $m_{555} \\simeq 20.3$ and had already been identified by Castro et al. (2001) as a $\\tau \\simeq 2~Gyrs$ field population; the other is fainter, at $m_{555} \\simeq 21$, $0.4 < m_{555} - m_{814} < 0.5$ and consists entirely of stars in the on-cluster data. It is this latest feature that we concentrate on, as a candidate NGC 1868 subpopulation. The MSTOs described here are indicated in Figure~1.  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_cmdcomp_41-51.ps,height=20cm,width=19.0cm,angle=0}} \\end{center} \\caption{A sample of artificial CMDs built from the NGC 1868 MS fiducial line (white line) and error distribution. The lower-right panel shows the actual NGC 1868 data, for comparison} \\end{figure*}  \\subsection[]{Removing contaminating field LMC stars}  In order to clean the left panel of Figure~1 from background field contamination, we statistically remove the control off-cluster CMD from the on-cluster one. Several procedures for doing so have been considered. One method is to match the off-cluster stars to corresponding stars in the on-cluster image using the estimated probabilities that two stars could be independent photometric measurements of each other. In order to estimate these matching probabilities, we use the empirically determined photometric uncertainties in both images and assume a Gaussian error distribution. Cluster stars are then randomly removed according to their probability of matching any of the off-cluster stars. In this approach a reliable estimate of the photometric uncertainties is very important. As pointed out previously, we used the independent magnitude measurements for the stars in common between the CEN and HALF fields to empirically estimate standard deviations ($\\sigma$) in the error distribution at different magnitude bins (Santiago et al 2001, Kerber et al 2002).  An alternative way of removing contaminating field stars is to bin both on-cluster and off-cluster data in magnitude (or in magnitude and colour) and to subtract the histogram of the latter from that of the former. At each bin, a subset of the on-cluster stars, corresponding to the subtracted histogram, is randomly selected. The two approaches have been tested and yield similar results. As our main concern is to assess the existence of a second, older population superposed to the dominant NGC 1868 population, the branch of evolved stars is critical, since it provides the best opportunity to single out stars from each individual population. Therefore, field subtraction was carried out separately for the MS and for the SGB/RGB locus. This latter is marked on Figure~1.  A more detailed discussion about the issue of field stars removal is presented in an upcoming paper (Kerber et al. 2002).  The cleaned CMD is shown in Figure~2. The CMD is now cut-off at the faint end ($m_{555} > 24$) as this region is not relevant to the current analysis. Below this limit, sampling incompleteness would only further complicate field star subtraction. We also cut it at bright magnitudes ($m_{555} < 19$) to avoid regions where saturation effects start to take place ($m_{555} \\simeq 19$ and $m_555 \\simeq 17.8$ for the HALF and CEN fields, respectively). The statistical removal of field stars has depleted the subgiant branch in the range $21$ \\ltsima $m_{555}$ \\ltsima $22$. A residual number of SGB stars brighter than $m_{555} \\simeq 21$ but much fainter than the cluster RC remains. As this excess of SGB stars reltive to the field is located at brighter magnitudes than the candidate second MSTO, it is a first evidence for the reality of this feature. As for the secondary MSTO itself, at $m_{555} \\simeq 21$, it was left untouched, which reflects the absence of field stars close to it. This again supports the reality of the second population in NGC 1868.  \\section[]{Statistical significance of the candidate secondary MSTO}  In this section we will assess whether the feature we tentatively identify as a secondary MSTO in the CMD of NGC 1868 is statistically significant. More specifically we address the question of whether a similar feature could originate from a single main sequence through a random realization of photometric errors in the data. We explore this possibility by means of Monte-Carlo realizations of the cleaned NGC 1868 CMD. In each realization we collapse the cluster MS onto a fiducial line and redistribute the data, using the measured photometric uncertainties and assuming Gaussian error distributions in magnitude and colour.  The main sequence fiducial line is determined by taking the median $m_{555} - m_{814}$ colour at different magnitude bins. We apply a 3$\\sigma$ clipping to the colour distribution in order to eliminate outliers and iterate until convergence. We initially define this MS line in bins spaced by 0.2 mag. A cubic spline interpolation is then used to fill this line in a much narrower ($\\Delta m_{555} = 0.02$) binning. Each true MS star is then randomly assigned to one of the points along the MS fiducial line, using a two-dimensional Gaussian probability function whose standard deviations are given by the star's measured magnitude and colour uncertainties. We then redistribute the data back onto the CMD plane, the magnitude and colour of each artificial star again being a random realization of the same Gaussian distribution of errors.  One hundred such Monte-Carlo simulations were carried out. A subset with 11 of the resulting CMDs is shown in Figure~3. The actual data are shown in the lower-right panel for comparison. The MS fiducial line is shown as a white line for guidance. The artificial CMDs have similar width as the data but are unable to reproduce the apparent turn-off at $m_{555} \\sim 21$, or the residual of the old field population turn-off ($m_{555} \\sim 22$), whose stars are farther to the red of the MS and much more clumped together than any set of points in any of the simulated CMDs. The simulations also fail to reproduce the main cluster MSTO at $m_{555} \\simeq 19$ and $m_{555} - m_{814} \\simeq 0.3$. This proves that {\\it scatter due to photometric errors alone cannot account for these features seen in the data}. Notice that these experiments do not explicitly account for the effect of unresolved binaries on the CMD. The fact that the distribution of true stars is skewed redwards from the MS fiducial line is likely caused by unresolved binaries. Incorporation of unresolved binaries would help spread out the simulated stars towards redder colours, but is unlikely to reproduce the features mentioned above, especially if we consider that part of the effect is already incorporated in the position of the MS fiducial line. We discuss the effect of binaries in more detail in \\S 5.2.  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_isofit.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Field subtracted cluster colour-magnitude diagram. Some Padova isochrones are superimposed to the data; their [Fe/H] and age (in Myrs) are indicated in the figure, as well as the adopted reddening value.} \\end{figure*}  \\section[]{The candidate NGC 1868 second population}  Assuming that the second MSTO at $m_{555} \\sim 21$ is real and belongs to the cluster, one wishes to determine what are the main characteristics of this population in NGC 1868.  In Figure~4 we again show the cleaned NGC 1868 CMD. Superimposed to the data we show three Padova isochrones (Girardi et al. 2000). In plotting the isochrones, we assume a distance modulus of $m - M = 18.5$ to the LMC (Panagia et al. 1991). The data were extinction corrected assuming E(B-V) =0.04 as indicated. Conversion from E(B-V) to E(555-814) was done as described by Holtzman et al (1995a,b). The $[Fe/H]$ and age (in Myrs) of each isochrone are shown in the figure. The [-0.4,800] isochrone fits the main NGC 1868 population. The other isochrones are attempts to fit the second turn-off and SGB. The younger and more metal-rich isochrone with [Fe/H] = -0.4 and an age of 2500 Myrs provides the best fit to both turn-off and SGB regions.    \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_turnoff_dist.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Dots: projected distribution of stars located at the CMD position of the secondary MSTO. The contours are circles centred on the cluster. The numbers shown vertically are the percentage of all cluster stars interior to each circle.} \\end{figure*}  Using star counts and this best fitting isochrones, we can estimate the mass fraction of this candidate older NGC 1868 population relative to the dominant and younger one. There are 11 stars in the range $2.5$ \\gtsima $M_{555}$ \\gtsima $1.5$ ($21.0$ \\gtsima $m_{555}$ \\gtsima $20.0$) and located along the SGB. Besides these, 13 stars are located within the CMD area whose limits are: $20.85 \\leq m_{555} \\leq 21.3$, $m_{555} - m_{814} \\leq 0.6$ and $m_{555} \\leq 5.333~(m_{555} - m_{814}) + 18.55$. This corresponds to the region where the second MSTO is clearly detached from the MS. Assuming that all such $N_{old} = 24$ stars are in fact associated with an older population belonging to NGC 1868 and using the best fitting isochrone, we infer the masses that correspond to the limiting range in absolute magnitudes: they are $m_{min} \\simeq 1.36~m_{\\odot}$ for the basis of MSTO position and $m_{max} \\simeq 1.49~m_{\\odot}$ for $M_{555} = 1.5$ ($m_{555} = 20$) along the SBG. As for the dominant cluster population we find $N_{young}$ = 732 stars along the cluster MS with $M_{555} < 2$ ($m_{555} < 20.5$) and fainter than the MS termination (we cut it at $m_{555} = 19$). These limits correspond to $1.63$ \\ltsima $m/m_{\\odot}$ \\ltsima $2.10$ for the best fitting isochrone to the NGC 1868 upper MS ([Fe/H],age(Myrs) = -0.4,800). Assuming the cluster's present day mass function (PDMF) to be a power-law with fixed slope $\\alpha$ within the mass ranges considered, the mass ratio will be given by  $${ {M_{old}} \\over {M_{young}} } = { {N_{old}} \\over {[ m_{max,old}^{1-\\alpha} - m_{min,old}^{1-\\alpha} ]} }~{ {[ m_{max,young}^{1-\\alpha} - m_{min,young}^{1-\\alpha} ]} \\over {N_{young}} }$$  Assuming a Salpeter value for the MF ($\\alpha = 2.35$), and the limiting mass values quoted above for each of the subpopulation, we then infer $M_{old} / M_{young} = 0.06$. For $\\alpha = 1.5$ we have $M_{old} / M_{young} = 0.08$ and for $\\alpha = 3.2$, $M_{old} / M_{young} = 0.05$. The mass ratios are also fairly insensitive to adopting the ([Fe/H],age(Myrs)) = (-0.7,3500) Padova isochrone for the secondary population; in this case, the old population mass range would be $1.22 < m/m_{\\odot} < 1.29$ and $M_{old} / M_{young}$ would change by no more than 50\\% for any choice of $\\alpha$. The mass ratio estimates are summarized in Table~1.  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=all_fidfld.ps,height=20cm,width=19cm,angle=0}} \\end{center} \\caption{CMDs of field LMC stars with the 5\\% highest values of $n_{sig}$ as defined in the text. Each panel corresponds to one particular field in the LMC as indicated. The central solid line is the MS fiducial line and the dashed lines correspond to its $\\pm~5\\sigma$ deviations. } \\end{figure*}  In addition to the PDMF slope and the model mass-luminosity relation, another source of uncertainty in the mass ratio quoted above is the residual contamination by field stars to the $N_{old}$ counts in the SGB region. Previous experiments with the field star subtraction (see \\S 2.1) methods have led to variations of up to $35\\%$ in the resulting value of $N_{old}$. If we assume this to be the uncertainty caused by field contamination, a similar relative error is expected to propagate into $M_{old} / M_{young}$. Notice that this is a conservative reasoning, since residual contamination by field LMC stars should also affect the younger population. We should also point out that this $35\\%$ uncertainty is larger than the Poisson fluctuation expected for the $N_{old}$ counts.  {\\bf \\begin{table} \\caption{Mass ratios for the two subpopulations of NGC 1868 stars for different choices of PDMF slope and assumptions for the age, metallicity and mass range of evolved stars in the secondary population} \\begin{tabular}{c c c c c c c c} \\hline isochrone & SGB mass range (solar) & $\\alpha$ & $M_{old} / M_{young}$ \\\\ \\hline $[-0.4,2500]$ & 1.36/1.49 & 2.35 & 0.06 \\\\ $[-0.4,2500]$ & 1.36/1.49 & 1.50 & 0.08 \\\\ $[-0.4,2500]$ & 1.36/1.49 & 3.20 & 0.05 \\\\ $[-0.7,3500]$ & 1.22/1.29 & 2.35 & 0.09 \\\\ $[-0.7,3500]$ & 1.22/1.29 & 1.50 & 0.12 \\\\ $[-0.7,3500]$ & 1.22/1.29 & 3.20 & 0.06 \\\\ \\hline \\end{tabular} \\end{table} }  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=all_histredfld.ps,height=20cm,width=19cm,angle=0}} \\end{center} \\caption{Frequency distribution of the stars shown in the previous figure as a function of $m_{555}$ magnitude. Again each panel corresponds to one of the LMC fields, as indicated. The histogram shown as a dotted line corresponds to the frequency distribution of all MS stars} \\end{figure*}  An important question is how the stars in this candidate secondary population are distributed throughout the cluster. Differences in spatial distribution between the two populations could indicate a recent merger, in which the two systems have not yet attained dynamical equilibrium. Figure~5 shows the projected distribution of the stars located at the second turn-off; they are the same 13 stars used in estimating the $M_{old}/M_{young}$ ratio earlier in this section. The 11 more evolved stars along the SBG were left out of Figure~5 because they represent just a statistical excess relative to field counts; hence, they are not necessarily true members of the older population. The contours represent circles around the cluster centre and the attached labels show the percentage of all cluster stars located inside them. We note that, within the fluctuations expected by the small numbers, the projected distribution of these stars is consistent with that of the entire cluster.  \\section {Alternative interpretations}  We have so far been interpreting the CMD feature at $m_{555} \\simeq 21$ and $m_{555} - m_{814} \\simeq 0.4$ as a second MSTO associated with NGC 1868. In this section we discuss other possible interpretations.  \\subsection {Is it a population associated with the LMC field?}  One obvious alternative is that the feature in the NGC 1868 CMD is in fact a MSTO, but associated with stars belonging to the general LMC field population rather than to the cluster itself. Notice however that this is unlikely, since most contaminating field stars have been statistically removed from the on-cluster CMD. Furthermore, Figure 5 shows that the stars located at the position of the second turn-off are concentrated towards the cluster centre, a strong evidence in favour of a cluster origin for them. On the other hand, the star formation history in the LMC is known to be complex, especially in the last few Gyrs, when the bar was formed (Gallagher et al 1996, Elson et al 1997, Smecker-Hane et al 2002). Thus, small scale variations in the field CMDs are possible, making field subtraction more uncertain.  One way to investigate possible MSTOs in field populations is to search for similar features in the 7 CMDs of field stars studied by Castro et al (2001). These are fields located 7.3' from the target clusters of the GO7307 HST project. In fact, Castro et al have visually identified possible turn-offs, with ages in the range $2$ \\ltsima $\\tau$ \\ltsima $4$ Gyrs, in several of these fields. In addition to those, turn-offs associated with an old ($> 10$ Gyrs) population were found in all field CMDs. In some cases, visual inspection of the CMDs revealed a broadening in the MS for $m_{555}$ \\ltsima $21.5$, indicative of continuous star formation in the LMC.  In order to quantify these visual impressions we tried to identify features in the field CMDs containing stars that are highly detached from the MS line, similar to the presumed NGC 1868 MSTO. One obvious such locus is the SGB/RGB region, which is always present in the CMDs of field LMC stars. It thus has to be eliminated from the analysis {\\it a priori}. We must also avoid contamination from background galaxies and faint stars belonging to the Galaxy. Unresolved galaxies should be limited to faint magnitudes, $m_{555}$ \\gtsima $22.5$, as brighter ones are visibly extended in our WFPC2 images. From counts of faint compact galaxies, one expects $\\sim 50$ contaminating galaxies within the $22.5 \\leq m_{555} \\leq 25$ range in our on-cluster field (Abraham et al 1996). As for foreground stars, a similar number is expected in the entire observed CMD (Santiago et al 1996). Together both types of contaminating sources contribute with a few percent of the total sample and will be spread out in the CMD.  We therefore proceeded as follows. For each field a fiducial line representing the CMD MS was defined in the same way as described in \\S 3. A low-order polynomial was fit to the MS fiducial line of most fields in order to smooth out the wiggles caused by noise in the median $m_{555} - m_{814}$ value at each $m_{555}$ bin. For NGC 1805 and NGC 1818, which have much larger numbers of MS stars than the other fields, especially at the bright end ($20.5 < m_{555} < 19$), the raw MS line was used. Evolved stars, possible background galaxies and faint stars belonging to the Galaxy were then eliminated by cutting out all CMD objects beyond $\\pm 5~\\sigma$ from the MS line, where $\\sigma$ is the empirically determined standard deviation in the $m_{555} - m_{814}$ colour distribution at each magnitude. For the remaining stars, we computed the number of standard deviations by which each star is detached from the MS fiducial line, $n_{sig} = \\Delta (m_{555} - m_{814}) / \\sigma$.  In Figure~6, we show the stars whose $n_{sig}$ values fall at the 95\\% position or beyond in each LMC field. The number of stars in each panel varies from 30 to 130. These stars would be $> 2~\\sigma$ (and $< 5~\\sigma$) events of a Gaussian error distribution in colour. The MS fiducial line and the $\\pm 5~\\sigma$ lines are also shown for guidance. Assuming that the photometric uncertainties have been adequately measured over the entire $m_{555}$ range and that these stars just reflect the high tail of the error distribution, we would expect them to follow the distribution of MS stars. Therefore, the shape of the distribution of such stars may reveal features that are not accounted for by errors alone. As an example, a clumped distribution of such stars at some magnitude range, as compared to the smooth distribution of MS stars, would indicate the existence of features such as a MSTO.  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_fidclus.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{CMDs of the NGC 1868 stars with the 5\\% highest values of $n_{sig}$ as defined in the text. The central solid line is the MS fiducial line and the dashed lines correspond to its $\\pm~5\\sigma$ deviations. } \\end{figure*}  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=n1868_histredclus.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Solid line: Frequency distribution of the stars shown in the previous figure as a function of $m_{555}$ magnitude. The dotted line corresponds to the frequency distribution of all MS stars in NGC 1868} \\end{figure*}  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=4clusters_coldist.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Colour distributions for NGC 1805, NGC 1818 (upper panel), NGC 1831 and NGC 1868 (lower panel). } \\end{figure*}  \\begin{figure*} \\begin{center} \\centerline{\\psfig{file=4clusters_gapdet.ps,height=10cm,width=9.5cm,angle=0}} \\end{center} \\caption{Panel {\\it a}: Gap center plotted against gap width for all high confidence gap candidates found in NGC 1805, according to the $\\chi^2$ statistics defined by Rachford \\& Canterna (2000). Panel {\\it b}: the same as in panel {\\it a}, but for NGC 1818. Panel {\\it c}: the same as in panel {\\it a}, but for NGC 1831. Panel {\\it d}: the same as in panel {\\it a}, but for NGC 1868.} \\end{figure*}  There is a common pattern in the distribution of these 5\\% reddest stars in the 7 LMC fields. In all cases, the majority of them are within the range $20 < m_{555} < 22$. This can be also seen by the frequency distribution of these stars as a function of $m_{555}$, which is shown in Figure~7. In all fields, large $n_{sig}$ stars show a distinct peak in their magnitude histogram when compared to the distribution of normal MS stars. This peak is usually broad and covers the $m_{555}$ range mentioned above. Examples are the fields close to NGC 1805, NGC 1818 and NGC 1868. NGC 1805 and NGC 1818 are close to star formation regions and their neighbouring field includes stars of different magnitudes (masses) leaving the MS towards the RGB. As for the field close to NGC 1868, a 2 Gyr MSTO is seen in its CMD (Castro et al 2001), contributing with stars brighter than $m_{555} = 20.5$.  In some cases, on the other hand, high $n_{sig}$ field stars show a narrow and very tall peak. This is the case of the field stars in the neighbourhood of NGC 1831, which display a distinct peak at $m_{555} \\simeq 21$; this is consistent with the position of the $\\simeq 4~Gyrs$ turn-off identified in this field by Castro et al (2001). Other examples of high peaks in the distribution of detached field stars include the fields close to Hodge 14 and NGC 2209. These peaks are at least $0.5$ mag fainter than that of NGC 1831 and are close to the basis of the SGB of the old LMC field populations. Inspection of the stars contributing to these peaks in fact confirms that they are dominated by SGB stars which were not successfully removed using the $5~\\sigma$ colour cut-off.  Figures 8 and 9 show, respectively, the CMD and $m_{555}$ histogram for the stars in the on-cluster, background cleaned NGC 1868 data whose $n_{sig}$ values are at the 95\\% position or beyond. These figures follow the same conventions as Figures 6 and 7. A total of 201 stars contribute to Figures 8 and 9. The larger number of stars allowed a magnitude binning in Figure 9 narrower than in Figure 7. Notice that the on-cluster histogram of high $n_{sig}$ stars as a function of magnitude now shows a large and narrow peak at a bright magnitude, $m_{555} \\simeq 19.5$, which is close to the dominant NGC 1868 MSTO. A second peak is seen in the range $20.5 < m_{555} < 22$, similar to those found in the field CMDs. This latter peak is well centered where the candidate secondary MSTO was found. Apart from the bright peak, however, the distribution of high $n_{sig}$ stars as a function of $m_{555}$ in the cluster data is not markedly different from those of field stars.  Thus, we conclude that a field origin for the second turn-off in NGC 1868 may not be ruled out. However, the differences in shape between this feature and those known to be associated with field LMC stars, plus the fact that the NGC 1868 CMD had previously been removed of field stars, argue against this interpretation. Furthermore, the suspected secondary turn-off in NGC 1868 is detached enough from the MS that the two reddest stars contributing to it actually are located beyond the $+5~\\sigma$ line shown in Figure 8 (see Figure~2), and therefore do not contribute to the histogram in Figure 9.  \\subsection {Selective unresolved binarism and Am stars}  As mentioned in \\S 3, there are other causes, besides a MS turn-off, that could lead to a set of stars being significantly detached from the MS towards redder colours.  Unresolved binarism is one such possibility. The presence of a secondary star, whose flux is added to the primary, will shift the position of the system towards brighter magnitudes in a CMD, relative to the position of the primary. If both are MS stars, the secondary will be redder, thus also affecting the colour of the system. Unresolved binaries are certainly present in any CMD and should be distributed all over the MS and among evolved stars. Their main effect is to make the MS broader with a red tail in the colour distribution. The feature seen in Figures 1 and 2, at $m_{555} \\simeq 21$, could be explained only if there is a specific increase in the fraction of cluster binaries at an absolute magnitude $M_{555} \\simeq 2.5$.  The apparent second MSTO occurs at the luminosity typical of peculiar A stars (Ap, Am, etc). Due to their chemical anomalies, Am stars in particular tend to be redder than normal A stars by about 0.05 mag as a result of increased line blanketing. This effect alone is too small to account for the feature in question but certainly goes in the right direction. As many Am stars are also known to be binaries (Carquillat et al 2001, Debernardi et al 2001), a combination of blanketing and enhanced binarism might accommodate the observed feature in the CMD of NGC 1868. However, more stringent constraints on the frequency of Am stars and on their binary fraction must be placed in order to test this possibility. Another problem with this interpretation is that global binary fraction estimates tend to be smaller in rich clusters than in the general field (Mayor et al 1996, Cot\\'e et al 1996, Elson et al 1998).  \\subsection {The B\\\"ohm-Vitense gap}  The onset of a convective envelope is known to occur for MS stars of spectral type A or later. This change in energy transport is expected to occur abruptly for stars with $T_{eff} \\simeq 7500~K$, or $(B-V) \\simeq 0.22$ (B\\\"ohm-Vitense 1958). As a result, stars cooler than this critical effective temperature would become redder than the slightly hotter ones with fully radiative envelopes (B\\\"ohm-Vitense 1970). The resulting discontinuity in the CMD is the so-called B\\\"ohm-Vitense gap and is thought to have an amplitude of $\\Delta (B-V) \\simeq 0.10$. Recent surveys combining precise photometry and distance measurements from astrometry have shown that a gap does exist and is located not too far from where expected by theory. But the evidence is often weak or subject to selection effects. Newberg \\& Yanny (1998) studied bright and nearby field stars with available photometry plus distances from Hipparcos and found a gap at $0.2 < (B-V) < 0.3$ in the CMD of field stars of luminosity class V only. In other words, the gap is masked out by evolved stars contaminating the apparent MS in their CMD. Rachford \\& Caterna (2000) detected a gap at $(B-V) \\simeq 0.35$ in most nearby open clusters they studied, but they argue that this gap may be unrelated to the onset of envelope convection.  We have carried out the same gap detection analysis as Rachford \\& Caterna (2000). We computed $\\chi^2$ values as defined by those authors for gap candidates within the $0.15 < m_{555} - m_{814} < 0.60$ colour range and with widths varying from 0.02 to 0.10 mag. Figure 10 shows the $m_{555} - m_{814}$ frequency distribution, not only for NGC 1868, but also for 3 other rich LMC clusters in our sample, each one with more than 5000 stars in their CMD (Santiago et al 2001). The clusters are paired according to age, NGC 1805 and NGC 1818 being younger ($\\tau < 10^{8}$ yrs, Johnson et al 2001) than NGC 1831 and NGC 1868. The colour distributions for the older clusters is slightly flatter than those for the younger ones. Apparent gaps are seen at $m_{555} - m_{814} \\simeq 0.40$ (in all cases) and at $m_{555} - m_{814} \\simeq 0.52$ (except NGC 1831). A broader gap is also present at $m_{555} - m_{814} \\simeq 0.25-0.30$ in NGC 1805 and NGC 1818 and possibly in NGC 1831 and NGC 1868.  Figure 11 shows the gap center plotted against gap width for all our gap candidates with $> 90\\%$ confidence ($\\chi^2 > 2.8$). Each panel corresponds to one of the 4 rich clusters discussed. The gaps at $m_{555} - m_{814} \\simeq 0.4$ are significant in all clusters except NGC 1818. At $m_{555} - m_{814} \\simeq 0.52$, only narrow and fairly unconspicuous gaps are confirmed. The bluer gap, at $m_{555} - m_{814} \\simeq 0.25$, takes different widths and positions, being more significant in NGC 1805 and NGC 1831. Hence, the main feature, clearly present in 3 rich clusters, is the gap at $0.35 < m_{555} - m_{814} < 0.40$. Notice that this colour coincides closely with that of the candidate MSTO, adding support to the statistical reality of this CMD feature. These gap limits correspond approximately to $0.23 < B-V < 0.27$, which is close to the expected position of the B\\\"ohm-Vitense gap and somewhat displaced from the gap position favoured by Rachford \\& Caterna (2000).  The presence of colour gaps in most of our rich LMC clusters calls for a more common mechanism than merger events. Our gap detection results, therefore, favour the possibility that the B\\\"ohm-Vitense gap was detected for the first time in LMC clusters. On the other hand, our gap colour is inconsistent with the one found by de Bruijne et al (2000, 2001) in their very accurate Hyades data. In fact, the very narrow CMDs by these authors provide the strongest limits on the gap position and associated colour discontinuity, which seems to be of $\\Delta (B-V) \\simeq 0.05$. This is significantly smaller than the colour difference between MS stars and the secondary MSTO found in our NGC 1868 data. ",
        "conclusions": "In this paper we have shown possible evidence for a second main sequence turn-off in a deep WFPC2/HST colour-magnitude diagram of NGC 1868. This feature is clearly visible in the cluster CMD, at $m_{555} \\simeq 21$ and $m_{555} - m_{814} \\simeq 0.4$, especially after contaminating field stars are removed. The presence of these stars, as well as of a residual SGB/RGB brighter than $m_{555} = 21$, is evidence for a secondary cluster population. In fact, previous ground-based data also show a hint of this feature: the distribution of stars seen redwards of the main-sequence in the NGC 1868 CMD by Corsi et al. (1994), despite the large scatter, has an excess of stars at the same position.  Assuming that the feature is really a MSTO associated with NGC 1868, isochrone fits yield an age of $2.5-3.5~Gyrs$  and a metallicity of $-0.7 < [Fe/H] < -0.4$ for this subpopulation. Using star counts in CMD regions where we expect only one cluster subpopulation to be present, we estimate the mass ratio of the older subpopulation relative to the younger to be in the range $0.05$ \\ltsima $M_{old} / M_{young}$ \\ltsima $0.12$. This range of values incorporates uncertainties in mass function slope, in age and metallicity of the secondary population, as well as uncertainties in the statistical removal of contaminating field LMC stars.  Even though the candidate second turn-off remains untouched after field stars are statistically subtracted from the CMD of NGC 1868, the possibility that this feature is associated with a field star population was investigated in detail and cannot be completely ruled out. CMDs of field stars from several different positions in the LMC show an excess of stars located in loci close to that of the NGC 1868 candidate turn-off. This excess is measured relative to the expected number of stars scattered to these loci due to photometric errors. Most of the features in the field CMDs, however, are broad and visually less conspicuous than the apparent secondary turn-off in NGC 1868, thus being more consistent with periods of enhanced field star formation lasting for longer than $1~Gyr$. Another argument against a field origin for the candidate second turn-off is that its stars are concentrated towards the cluster centre.  Other possible explanations for the CMD feature have been explored. They include CMD spread due to unresolved binaries, which could mimic a MS turn-off if binarism is enhanced within a narrow range of primary star masses. It is interesting to notice that the mass range in question is close to that of Am stars, which are redder than normal A stars and for which a larger than usual binary fraction may exist. Estimates of binary fraction within globular clusters, however, yield smaller values than in the field. Besides, the binary fraction among Am stars is not yet well constrained. It is thus unlikely that the CMD feature discussed here is due to Am stars or unresolved binaries.  A final and exciting alternative would be that the B\\\"ohm-Vitense gap was for the first time detected in a stellar population outside the Galaxy. The fact that gaps were found in the colour distribution of two other LMC clusters besides NGC 1868, all of them at $m_{555} - m_{814} \\simeq 0.4$ ($(B-V) \\simeq 0.25$), calls for a more common mechanism than mergers. In fact, a merger of the type suggested by the presence of a second and much fainter MSTO as discussed in this paper should be rare (Vallenari et al 1998, Dieball \\& Grebel 2000). Most model predictions favour merging units that are nearly coeval and usually formed through encounters within the same giant molecular cloud, either before or after they are fully formed clusters (Fujimoto \\& Kumai 1997, Efremov \\& Elmegreen 1998, Leon et al 1999).  On the other hand, it is unclear to what extent the B\\\"ohm-Vitense gap may mimic a main sequence turn-off. The recent and precise CMD data on the Hyades by de Bruijne et al (2000, 2001) do not reveal strong turn-off like features associated with the B\\\"ohm-Vitense gap candidates found by those authors; the colour discontinuity in their CMD is of smaller amplitude than necessary to account for the NGC 1868 feature studied here. Besides, the gap associated with our NGC 1868 CMD feature is bluer than the gaps found by de Bruijne et al.  We should point out that the nature of the stars that make up the candidate MSTO in NGC 1868, which currently is uncertain and accountable for by at least two astrophysically interesting alternatives (merging and B\\\"ohm-Vitense gap), may be established by spectroscopic classification, a task likely to demand large, 8m class, telescopes."
    },
    "0207/astro-ph0207491_arXiv.txt": {
        "abstract": "{\\small We analyzed 139 $\\chi$-state observations of \\grs\\ with {\\it RXTE} from 1997 to 2000 and found i) that the observations fall into two groups with different Comptonization behavior, ii) that the slope of the hard X-ray component correlates with the radio flux, thus revealing the interaction of jet and corona, and iii) a 590\\,days long term periodicity in the hard X-ray and radio components.} ",
        "introduction": "\\vspace{-0.1cm} The prototypical microquasar \\grs\\ was discovered by {\\it Granat} \\cite{cbl92} as a transient X-ray source. {\\it RXTE} observations revealed astonishing X-ray variability \\cite{gmr96} which has been classified by Belloni \\etal \\cite{bmk97} into twelve different X-ray states.  The analysis presented here is based on the data reduction described in full detail in Rau \\& Greiner \\cite{rg02}. We analyzed 139 observations from 89 days from the RXTE public archive of \\grs\\ from November 1996 to September 2000, when the source was in the $\\chi$-state \\cite{bmk97}. The absence of large amplitude variations and structured variability in these states suggests quasi-stable geometry and parameters during each observation and allows to fit the spectrum of an entire observation at once. $\\chi$-states are connected with radio emission of varying strength and the most common states observed.  From each observation we used data from PCU0 and HEXTE cluster 0 and fitted these with a model consisting of cold absorption (WABS), a disk blackbody (DISKBB) and a power-law spectrum reflected from an ionized relativistic accretion disk (REFSCH \\cite{frs89},\\cite{mz95}) in XSPEC \\cite{a96}. Energies between 4 and 8.5\\,keV were ignored during the fitting process because of the known response problems of the PCA at these energies. The hydrogen column density was fixed at $N_H$=5$\\cdot$10$^{22}$\\,cm$^{-2}$.  Our model implies obvious simplifications, e.g. compared to the theoretical Comptonization models the power law overestimates the flux at low energies. Despite the simplicity of the three component model, we achieved surprisingly good fit results for our sample of $\\chi$-state observations (reduced $\\chi^2$$<$2 for all and $\\chi^2$$<$1.1 for more than 60\\% of the observations).   \\vspace{-0.1cm} ",
        "conclusions": ""
    },
    "0207/astro-ph0207485_arXiv.txt": {
        "abstract": "The nearby Mira-like variable \\lpup{} is shown to be undergoing an unprecedented dimming episode.  The stability of the period rules out intrinsic changes to the star, leaving dust formation along the line of sight as the most likely explanation.  Episodic dust obscuration events are fairly common in carbon stars but have not been seen in oxygen-rich stars. We also present a 10-$\\mu$m spectrum, taken with the Japanese IRTS satellite, showing strong silicate emission which can be fitted with a detached, thin dust shell, containing silicates and corundum. ",
        "introduction": "\\lpuppis{} (HR 2748; HIP 34922) is a bright nearby red giant with a pulsation period of about 140\\,d.  Its spectral type of M5eIII and luminosity of 1500 L$_\\odot$ indicate that it is evolving towards the tip of the Asymptotic Giant Branch (AGB).  Evidence for mass loss at a rate of $3 \\times 10^{-7}\\rm \\, M_\\odot \\, yr^{-1}$ \\citep{JCP2002} supports this. \\lpup{} is possibly the nearest star in this evolutionary phase, at a Hipparcos distance of $61 \\pm 5$\\,pc.  Among known long-period AGB stars, only R~Doradus has a similar distance.  At 12 microns, \\lpup{} is among the 15 brightest sources in the IRAS point source catalogue.  \\lpup{} is unusual in several respects.  Firstly, it shows a high degree of optical polarization, with a variable wavelength dependence that implies a long timescale for the growth and dissipation of dust grains (of the order of a decade; \\citealt{MCLG86}).  Secondly, CO measurements by \\citet{K+O99} indicate a very low expansion velocity (about 2.5\\,km\\,s$^{-1}$), which led them to label \\lpup{} as an extreme case, with one of the smallest expansion velocities ever measured for an AGB star.  The slow wind from \\lpup{} led \\citet{WLBJ2000,WLBN2002} to suggest that this star could represent their B-model, in which mass loss is driven entirely by pulsations, without any significant input from radiation pressure on dust grains.  This has been further discussed by \\citet{JCP2002}, who modelled the mass loss and suggested that the pulsations may be non-radial.  Thirdly, as we report here, this star has shown a remarkable change in mean visual magnitude over the past century, and is currently undergoing a dramatic dimming.  We present visual and infrared photometry which characterizes this behaviour, and argue that the most likely cause is the formation of dust along the line of sight.  We also present the first 10-micron spectrum of \\lpup, obtained with the Japanese IRTS satellite, which shows strong silicate emission. ",
        "conclusions": "Visual photometry of \\lpup{} shows an unprecedented dimming over the past 5 years.  The long-term light curve shows stable periodicity, and we argue that \\lpup{} is Mira-like and should be classified as SRa.  The period stability implies a constant stellar radius, which rules out temperature and/or luminosity variations as the cause of the dimming.  Rather, the dimming seems to arise from an episode of dust formation close to the extended atmosphere.  Episodic dust obscuration events are fairly common in carbon stars but have not been seen in (non-symbiotic) oxygen-rich stars. We suggest that dust forms continuously but anisotropically, with the current dimming event being due to to dust formation along the line of sight.   The red colours indicate reddening from dust, but the extinction curve is greyer than found for ISM dust.  This could reflect a higher fraction of oxides.  The change of colour during the dimming indicates that already before the dimming, the $V$-band magnitude was significantly affected by circumstellar or atmospheric extinction.  \\lpup{} was one of the few stars located below the Mira $P$--$L$ relation: the derived $K$-band extinction puts the star in closer agreement with this relation.  We present a 10-$\\mu$m spectrum showing strong silicate emission.  These observations were carried out in 1995, just prior to the recent dimming. The silicate feature can be fitted with a detached, thin dust shell, with inner radius $7 \\times 10^{14}\\,\\rm cm$ and outer radius $\\approx 1.5 \\times 10^{15}\\,\\rm cm$.  Longer wavelength photometry shows no evidence for more distant, colder dust. We derive a mass-loss rate of $\\dot M_{\\rm g} \\approx 5 \\times 10^{-7}\\,\\rm M_\\odot \\,yr^{-1}$, but this value depends on the assumed expansion velocity and metallicity --- if the dust velocity is high, the actual mass-loss rate could be higher."
    },
    "0207/astro-ph0207166_arXiv.txt": {
        "abstract": "{ We present the results of the X-ray spectral analysis of the first deep X-ray survey with the XMM-Newton observatory during Performance Verification.  The X-ray data of the Lockman Hole field and the derived cumulative source counts were reported by Hasinger et al. (2001).  We restrict the analysis to the sample of 98 sources with more than 70 net counts (flux limit in the [0.5-7] keV band of $1.6 \\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$) of which 61 have redshift identification. We find no correlation between the spectral index $\\Gamma$ and the intrinsic absorption column density N$_{\\rm H}$ and, for both the Type-1 and Type-2 AGN populations, we obtain $\\langle\\Gamma\\rangle \\simeq 2$.  The progressive hardening of the mean X-ray source spectrum with decreasing flux is essentially due to an increase in intrinsic absorption. The marked separation between the two AGN populations in several diagnostics diagrams, involving X-ray colour, X-ray flux, optical/near IR colour and optical brightness, is also a consequence of different absorption column densities and enables the classification of optically faint obscured AGN. The Type-2 and obscured AGN have weaker soft X-ray and optical fluxes and redder R$-$K$^\\prime$ colours. They follow the evolutionary tracks of their host galaxies in a color-redshift diagram.  About 27$\\%$ of the subsample with R$-$K$^\\prime$ colour are EROs (R$-$K$^\\prime \\geq 5$) and most of these 18 X-ray selected EROs contain an obscured AGN as revealed by their high X-ray-to-optical/near IR flux ratios.  There are six sources in our sample with ${\\rm L_X}$[0.5-10]$>10^{44}$ erg s$^{-1}$ and ${\\rm log(N_H)}>10^{22}$ cm$^{-2}$: which are likely Type-2 QSOs and we thus derive a density of $\\sim 69$ objects of this class per square degree. ",
        "introduction": "The deep ROSAT survey of the Lockman Hole showed that about 80$\\%$ of the soft (0.5-2 keV) X-ray background (XRB) is resolved into discrete sources (Hasinger et al. \\cite{gunther98}). These findings have recently been confirmed and strengthened using the two deep Chandra surveys of 1 Msec each (Brandt et al. \\cite{brandt01}; Rosati et al. \\cite{piero02}). An important population of X-ray sources with hard spectra, most probably obscured active galactic nuclei (AGN), is present in the Chandra (Barger et al. \\cite{barger01}; Hornschemeier et al. \\cite{hornschemeier01}; Rosati et al. \\cite{piero02}) and XMM-Newton (Hasinger et al. \\cite{paper1}, hereafter Paper I) deep surveys; a few objects of this class had already been detected in ROSAT deep and shallower surveys (Lehmann et al. \\cite{ingo01a}; Mittaz et al. \\cite{mittaz99}). In the hard band (2-10 keV), the X-ray source density derived from the number counts in the two Chandra deep surveys is about 4000 deg$^{-2}$ (Brandt et al. \\cite{brandt01}; Rosati et al. \\cite{piero02}) resolving $\\sim 85-90 \\%$ of the 2-10 keV XRB. This population of X-ray sources show a progressive hardening of the average X-ray spectrum towards fainter fluxes (Tozzi et al. \\cite{paolo01}; Mittaz et al. \\cite{mittaz99}).   The XMM-Newton deep survey ($\\simeq$ 100 ksec of good quality data) of the Lockman Hole was obtained during Performance Verification. The X-ray data reduction and analysis (restricted to sources within a 10 arcmin radius) was reported in Paper I where it was demonstrated that the different populations of X-ray sources are well separated in X-ray spectral diagnostics based on hardness ratios. The extensive optical follow-up programs of this field (Lehmann et al. \\cite{ingo01a}, and references therein) provide an understanding of the physical nature of the X-ray sources. The point sources detected in the soft band by ROSAT are predominantly unobscured (in both optical and X-ray bands) AGN spanning a wide redshift range.  In the XMM-Newton sample, there is a significant fraction of sources with hard spectra. This new population is most probably dominated by intrinsically absorbed AGN. This assumption can be tested using the available optical spectra and, more efficently, by X-ray spectral study.  To this aim, we have performed an X-ray spectral analysis of the sources in the Lockman Hole to understand their physical nature combining the X-ray data with the optical/near IR information. We also use the subsample with redshift identification to check the validity of our conclusions concerning the specific properties of the obscured AGN population. Preliminary results of this work were reported by Mainieri et al. (\\cite{mainieri02}).\\\\ In the following we will refer to Type-1 (broad and narrow emission lines) and Type-2 AGN (high ionization narrow emission lines) using the optical spectroscopic classification.  The observations are presented in Sect. 2. The results of the spectral analysis are described in Sect. 3, in particular the range of the X-ray spectral index, the observed ${\\rm N_H}$ distribution and colour-colour diagnostics diagrams. The optical/near IR properties are discussed in Sect. 4 together with a comparison with QSO and galaxy evolutionary tracks. The search for relations between X-ray and optical/near IR fluxes is presented in Sect. 5. The effect of the absorbing column density on the X-ray luminosity and the Type-2 QSO candidates are discussed in Sect. 6.  Representative spectra of the different classes of X-ray sources are given in Sect. 7. Finally, our conclusions are outlined in Sect. 8. ",
        "conclusions": "We have discussed the X-ray spectral properties of a sample of 98 sources found in the 100 ksec XMM-Newton observation of the Lockman Hole, using data from the EPIC-pn detector. The large throughput and the unprecedented sensitivity at high energies of the X-ray telescope and detectors allow us, for the first time, to measure separetely the intrinsic absorption and the slope of the power law emission spectrum for the faint source population. We have derived the spectral index ($\\Gamma$) and the column density (${\\rm N_H}$) for sources with more than 70 counts in the [0.5-7] keV band. We find that the value of $\\Gamma$ is independent of the absorption level with $<\\Gamma> \\approx 2$. Thus, we infer that the progressive hardening of the X-ray spectra of faint sources observed in Chandra deep fields (Giacconi et al. \\cite{giacconi01}; Tozzi et al. \\cite{paolo01}; Brandt et al. \\cite{brandt01}) is mainly due to the increasing level of intrinsic absorption rather than intrinsically flat spectra.  We confirm that the ${\\rm R-K}^\\prime$ colours of X-ray counterparts get redder towards fainter R magnitudes. Such a trend is not present between ${\\rm R-K}^\\prime$ and the K$^\\prime$ magnitude; this is likely due to a combination of a less pronounced absorption effect in this band, a different K-correction for AGN-type spectra (small) and star-like galaxy spectra (large), as well as an increased contribution of the host galaxy light in the K$^\\prime$ band relative to that of the AGN.  Comparing the ${\\rm R-K}^\\prime$ colours of the X-ray sources with evolutionary tracks of various galaxy-types as a function of redshift, we find that Type-2 AGN have colours dominated by the host galaxy and are also significantly absorbed (log ${\\rm N_H}>21.5$). On the other hand, for Type-1 AGN, the large majority of which are unabsorbed, the nuclear component is significantly contributing to their optical colours.  In addition, there is a strong correlation between the ${\\rm R-K}^\\prime$ colour and the amount of intrinsic X-ray absorption.  We have also defined an X-ray selected sample of 18 EROs (${\\rm R-K}^\\prime \\geq 5$) and found that it mainly comprises X-ray absorbed objects with a strong correlation between colour and intrinsic column density.  We have derived the unabsorbed rest-frame luminosities of the sources with strong intrinsic absorption. There are six absorbed, bright X-ray objects in our sample with ${\\rm L_X}[0.5-10]>10^{44}$ erg s$^{-1}$ and ${\\rm log(N_H)}>10^{22}$ cm$^{-2}$: one is an optically classified Type-1 QSO (source $\\#96$ see Sect. \\ref{sec:qso1}), two are Type-2 AGN and the remaining three have a photometric redshift and due to their X-ray absorption and optical/near-IR colours likely Type-2 AGN. Four of them are also EROs (${\\rm R-K}^\\prime \\geq 5$). These are likely to be Type-2 QSO candidates and we derive a density of $\\sim 69$ objects of this class per square degree at a flux limit in the [0.5-7] keV band of $1.6 \\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$.  Our analysis of the unidentified sources (mostly newly detected XMM sources) shows that the majority of these sources have absorbed X-ray spectra and are consequently located in the harder part of the diagnostic X-ray colour-colour diagrams. They are also optically fainter ($\\sim80\\%$ of them have R$>24$) and their optical-to-near-IR colours are redder ($\\sim90\\%$ have R$-$K$^\\prime \\geq 4$) than already identified sources. Their X-ray-to-optical flux ratios are $\\log (\\frac{f_X[2-10]}{f_R})>1$. From these properties, we argue that the majority of these sources are Type-2 AGN. This is confirmed by our on-going optical spectroscopic survey which is showing that the bulk of these sources is at $z<1$. Two X-ray bright optically ``normal'' galaxies are present in our sample. Their X-ray spectra are clearly absorbed suggesting the presence of an obscured AGN. We expect this class of objects to increase from the optical identification of the newly detected XMM-Newton sources."
    },
    "0207/astro-ph0207399_arXiv.txt": {
        "abstract": "In this contribution, I touch on a subset of our recent efforts in spectral and opacity modeling aimed at improving our understanding of brown dwarfs, L dwarfs, and T dwarfs.  I discuss theoretical calculations of the alkali line profiles, newly generated CrH opacities, new evidence for refractory rainout in T dwarfs from optical spectral measurements, and the distinction between brown dwarfs and planets. ",
        "introduction": "The subject of brown dwarfs (and substellar-mass objects in general) is entering a new phase of rapid expansion and discovery.  More than 200 L dwarfs are now known, and they are joined by $\\sim$30 T dwarfs.  This development requires a corresponding expansion in theoretical effort, involving at its core evolutionary, spectral, and compositional modeling.  Crucial to spectral modeling are molecular and atomic opacities, many of which have not been addressed before with the degree of seriousness and completeness that standard stellar atmospheres studies have long enjoyed.  However, the pace of relevant spectral and opacity calculations is accelerating and in this spirit I summarize in this paper a few such topics of recent interest.  In \\S\\ref{rain}, I show new proof of the ``rainout\" and settling of refractory elements in T dwarfs.  Rainout leaves as one of its consequences the lower-temperature reaches of T dwarf atmospheres enhanced in sodium and potassium atoms.  In \\S\\ref{profile}, I touch on some new calculations of the alkali line wings to many 1000's of \\AA\\ detunings. The red wing of the K I line centered at 0.77 \\mic in particular defines much of the T dwarf continuum between 0.77 \\mic and 1.0 \\mic. This is followed in \\S\\ref{crh} by a short discussion of the new CrH opacities recently generated by our group.  Finally, and a bit tentatively, in \\S\\ref{name} I finish with a short discussion on nomenclature, a subject that continues to exercise the brown dwarf and extrasolar planet communities. ",
        "conclusions": ""
    },
    "0207/astro-ph0207350_arXiv.txt": {
        "abstract": "The dust produced in the Kuiper Belt (KB) spreads throughout the Solar System forming a dust disk. We numerically model the orbital evolution of KB dust and estimate its equilibrium spatial distribution and its brightness and spectral energy distributions (SED), assuming greybody absorption and emission by the dust grains. We show that the planets modify the KB disk SED, so potentially we can infer the presence of planets in spatially unresolved debris disks by studying the shape of their SEDs. We point out that there are inherent uncertainties in the prediction of structure in the dust disk, owing to the chaotic dynamics of dust orbital evolution imposed by resonant gravitational perturbations of the planets. ",
        "introduction": "Main sequence stars are commonly surrounded by cold far-IR-emitting material. The fact that this infrared excess is not restricted to young stars, and that the dust grain removal processes, Poynting-Robertson (P-R) and solar wind drag, act on timescales much smaller than the age of the system, indicate that: (1) a reservoir of undetected dust-producing planetesimals exists; and (2) to induce frequent mutual collisions, their orbits must be dynamically perturbed by massive planetary bodies. The Solar System is also filled with interplanetary dust. In the inner Solar System, this dust, which gives rise to the zodiacal light, has been observed by Pioneer 10 (out to 3.3 AU) and by the infrared telescopes IRAS and COBE. The dominant sources of the zodiacal cloud are debris from Jupiter family short period comets and asteroids (Liou et al.,~\\citeyear{liou95}; Dermott et al.,~\\citeyear{derm92}). The discovery of a debris disk around $\\beta$-Pictoris, extending to 100s of AU, together with the confirmation of the existence of the theoretically predicted Kuiper Belt objects (KBOs) (Jewitt \\& Luu,~\\citeyear{jewi95}), suggest that significant dust production may also occur in the outer Solar System due to mutual collisions of KBOs (Backman \\& Paresce,~\\citeyear{back93}; Backman, Dasgupta \\& Stencel,~\\citeyear{back95}; Stern,~\\citeyear{ster96}) and collisions with interstellar grains (Yamamoto \\& Mukai,~\\citeyear{yama98}).  Dust particles are small enough to experience the effect of radiation and stellar wind forces. Radiation pressure makes their orbital elements and specific orbital energy change immediately upon release from parent bodies.  If their orbital energy becomes positive, the dust particles escape on hyperbolic orbits. In the Solar System, these particles are known as $\\beta$-meteoroids (Zook \\& Berg,~\\citeyear{zook75}). If their orbital energy remains negative, the dust particles stay on bound orbits. P-R and solar wind drag tends to circularize and decrease the semimajor axis of these orbits, forcing these particles to slowly drift in towards the central star (Burns, Lamy \\& Soter,~\\citeyear{burn79}). Assuming that the dust particles are constantly being produced, this drifting in creates a dust disk of wide radial extent, that we refer to as a $\\it{debris~disk}$. Debris disks are systems that satisfy the following conditions: (1) their age is longer than the P-R and collisional lifetimes; (2) they are optically thin to stellar radiation, even along the mid plane; and (3) they have little or no gas, so that the dust dynamics is controlled by gravitation and radiation forces only (Backman, \\citeyear{back02}).  When planets are present, the journey of the dust particle towards the central star is temporarily interrupted by the trapping of the particle in Mean Motion Resonances (MMRs). MMRs occur when the orbital period of the particle is in a ratio of small integers to that of the perturbing planet. [The p:q MMR means that the orbital period of the particle is p/q times that of the planet.] In an MMR, the drifting in is halted because the energy loss due to P-R drag is balanced by the resonant interaction with the planet's gravity field. This trapping can potentially create structure in debris disks, as the particles accumulate at certain semimajor axes. Sufficiently massive planets may also scatter and eject dust particles out of a planetary system, creating dust free or depleted zones. This structure, if observed, can be used to infer the presence of planets. Liou \\& Zook (\\citeyear{liou99a}, hereafter LZ99) found that the presence of the Giant Planets has an important effect on the structure of the debris disk that is presumably generated in the KB: Neptune creates a ring-like structure between 35 and 50 AU, due to the trapping of particles in exterior MMRs, and Jupiter and Saturn are responsible for the ejection of about 80\\% of particles from the Solar System (Liou, Zook \\& Dermott,~\\citeyear{liou96}, hereafter LZD96). The latter creates a clearing in the inner 10 AU that resembles the inner gap in the $\\beta$-Pictoris disk. If observed from afar, the KB disk would be the brightest extended feature in the Solar System, and its structure, if spatially resolved, could be recognized as harboring at least two giant planets: an inner planet (Jupiter plus Saturn) and outer planet (Neptune) (LZD96). In anticipation of future observations of debris disks, whose structure is likely to be spatially unresolved, in this paper we are interested in studying how the structure affects the shape of the disk SED and consequently if the SED can be used to infer the presence of planets.  In this paper we are going to follow numerically, from source to sink, the evolution of several hundred dust particles from the KB in the size range from 1 to 40 $\\mu$m (for $\\rho$=2.7 g/cm$^{3}$), or from 3 to 120 $\\mu$m (for $\\rho$=1 g/cm$^{3}$), under the combined effects of solar gravity, solar radiation pressure, P-R and solar wind drag and the gravitational forces of 7 planets (excluding Mercury and Pluto). The sinks of dust included in our numerical simulations are: (1) ejection into unbound orbits; (2) accretion onto the planets; and (3) orbital decay to less than 0.5 AU heliocentric distance. The equations of motion are integrated using a modification of the multiple time step symplectic method SyMBA (DLL98). In $\\S$2 we describe our numerical integration method and the tests performed to check the suitability of the code. $\\S$3 describes our methods for deriving the equilibrium spatial distribution of the dust disk. $\\S$4 explains the distribution of parent bodies and the orbital evolution of dust. In $\\S$5 we discuss the formation of structure in the KB debris disk and its observational signatures. Dust destruction processes are discussed in $\\S$6, and $\\S$7 summarizes our results.  \\label{intro} ",
        "conclusions": "\\label{concl} (1) We have followed, from source to sink, the orbital evolution of dust particles from the Kuiper Belt. To integrate the equations of motion efficiently, we have introduced radiation and solar wind forces in the multiple time step symplectic integrator of DLL98. We have established the suitability of our code by comparison between numerical results and analytical solutions to 2-body and restricted three-body cases, as well as comparison with other numerical results in the literature (LZD96, LZ99).  (2) We have carried out numerical simulations for single size particle disks in the presence and in the absence of planets in order to estimate the uncertainties inherent in the prediction of structure in the outer solar system debris disk, owing to the chaotic dynamics of dust orbital evolution.  We simulate dust particle initial conditions according to the wider distribution of parent bodies indicated by the recent observed distribution of KBOs, and our simulations extend to larger particle sizes than previous studies.  (3) We find that the distribution of KB dust particle lifetimes in the Solar system are described as a sum of a gaussian and a nearly uniform distribution; the latter represents only a small fraction of all particles but extends to very long lifetimes, while the gaussian represents the dominant fraction of particles. The mean and dispersion of the gaussian component increases systematically with particle size, and is in the range of [few million years] for [1--100 $\\mu$m] particle sizes. We do not find any correlations between the initial orbital elements and dynamical lifetimes of dust particles.  (4) We have examined carefully the method used by LZ99 to estimate the equilibrium spatial distribution of KB dust in the Solar System. This method is based on the ergodic assumption, so the dust structures obtained are determined to a large extent by the longest lived particles, which represent only a very small fraction of the dust population. The ergodic assumption is generally not applicable in chaotic dynamical systems. Nevertheless, we have established that in practice this method gives reliable results for several aspects of dust dynamical studies for three reasons: (i) the distribution of dust particle lifetimes is described as a sum of a gaussian plus a nearly uniform distribution, i.e. the longest-lived particles are not anomalous, they are statistically representative of the long tail population; (ii) the dust spatial structure is created quickly; (iii) the radial profile of the equilibrium number density distribution does not strongly depend on the longest-lived particles (although the azimuthal structure does).  (5) Overall, the number density of the KB dust disk shows a depletion of dust in the inner 10 AU, due to gravitational scattering by Jupiter and Saturn, and an enhanced dust density in a ring between 35 and 50 AU, due to trapping of particles in MMRs with Neptune. The structure is more pronounced for larger particle sizes. The brightness distribution shows a bright ring between 10 and 15 AU with a sharp inner edge (particles ejected by Saturn and Jupiter), and a steep increase in brightness in the inner few AU (a combination of the decreasing density and increasing grain temperature).  (6) We find that the azimuthal structure of the dust disk is not predictable in detail, except for a `gap' near the outermost planet Neptune.  This is because the azimuthal structure depends sensitively on the long lived particles trapped in mean motion resonances with Neptune, and the times of residence in the various resonances are highly variable and unpredictable.  (7) We have calculated disk brightness density and spectral energy distributions (SED), assuming greybody absorption and emission from the dust grains. We find that the presence of planets modifies the shape of the SED. The Solar System debris disk SED is particularly affected by the clearing of dust from the inner 10 AU due to gravitational scattering by Jupiter and Saturn.  (8) Grain physical lifetimes are limited by collisions and sublimation. The comparison of the dynamical lifetime of particles, the timescale for structure formation and the collisional time between KB and interstellar grains indicates that, if the current estimates for the flux and the size distribution of interstellar grains are correct, collisional destruction is important for grains larger than about 6 $\\mu$m. For smaller particles, debris disk structure will be able to survive, although the smaller particles have less prominent structure associated with the outer planets. Depending on their composition, sublimation of particles may or may not play an important role in the destruction of structure. If KB dust has water ice composition, and assuming a sublimating temperature of 100 K, it is likely that even large 120 $\\mu$m particles will sublimate before reaching the inner 4 AU of the Solar System. We conclude that grain destruction processes need to be examined more carefully in future applications of our studies to infer the presence of planets from structure in debris disks.  This work is part of the SIRTF FEPS Legacy project\\footnote{http://feps.as.arizona.edu} (P.I. M. Meyer), with the goal ``to establish the diversity of planetary architectures from SEDs capable of diagnosing the radial distribution of dust and the dynamical imprints of embedded giant planets''. The modeling of a particular system is very complex, because it involves a large number of free parameters. We have therefore chosen a forward modeling approach: a grid of models will be created for different planetary masses and orbital radii, parent bodies' masses and orbital distribution, total mass in dust particles, etc. We will produce dust spatial distributions like the ones presented here which will be used as input for a radiative transfer calculation to generate SEDs containing all the important spectroscopic features.  This will be more detailed than the simple greybody approximation used in the present work. This ``library'', that as part of our Legacy will be available to the community, will contain the templates to which we will compare the dust SEDs derived from the SIRTF observations for their interpretation in terms of planetary architectures.   \\begin{center} {\\it Acknowledgments} \\end{center} We thank Hal Levison for providing the SKEEL computer code. AMM is supported by NASA contract 1224768 administered by JPL. RM is supported by NASA grants NAG5-10343 and NAG5-11661."
    },
    "0207/astro-ph0207216_arXiv.txt": {
        "abstract": "The remarkably filamentary spatial distribution of young stars in the Taurus molecular cloud has significant implications for understanding low-mass star formation in relatively quiescent conditions.  The large scale and regular spacing of the filaments suggests that small-scale turbulence is of limited importance, which could be consistent with driving on large scales by flows which produced the cloud.  The small spatial dispersion of stars from gaseous filaments indicates that the low-mass stars are generally born with small velocity dispersions relative to their natal gas, of order the sound speed or less.  The spatial distribution of the stars exhibits a mean separation of about 0.25 pc, comparable to the estimated Jeans length in the densest gaseous filaments, and is consistent with roughly uniform density along the filaments. The efficiency of star formation in filaments is much higher than elsewhere, with an associated higher frequency of protostars and accreting T Tauri stars. The protostellar cores generally are aligned with the filaments, suggesting that they are produced by gravitational fragmentation, resulting in initially quasi-prolate cores.  Given the absence of massive stars which could strongly dominate cloud dynamics, Taurus provides important tests of theories of dispersed low-mass star formation and numerical simulations of molecular cloud structure and evolution. ",
        "introduction": "The Taurus-Auriga molecular cloud long has been a touchstone for studies of star formation.  Although Taurus is not typical of most (massive) star-forming regions, its proximity and low extinction mean that its stellar population is the best determined of any star-forming cloud.  Perhaps more importantly, Taurus is relatively quiescent, with low turbulent velocities and a lack of massive stars to dissociate, ionize, and otherwise disrupt the cloud.  If the simple, static models of the original paradigm of low-mass star formation (e.g., Shu, Adams, \\& Lizano 1987) can be applied anywhere, they should work in Taurus.  The spatial distributions with which stars are formed can provide important clues to the processes of star formation. Even though Taurus does not contain populous clusters, many of its stars fall into loose groups (Jones \\& Herbig 1979; Gomez \\etal 1993).  So far, the most detailed studies of the stellar spatial distribution in Taurus have considered the two-point correlation function (Gomez \\etal 1993) and the related mean surface density of companions (MSDC; Larson 1995).  Larson (1995) found that the MSDC exhibited roughly two distinct power-law distributions at small and large scales, the inner region corresponding to close binaries and multiple systems, and the outer region representing the clustering properties of the stars.  A break between these two distributions was identified at about $\\sim 0.04$~pc, which Larson suggested was the Jeans length in Taurus (see also Simon 1997 and Bate, Clarke, \\& McCaughrean 1998).  The molecular gas in Taurus has long been recognized to be filamentary in nature (e.g., Schneider \\& Elmegreen 1979; Scalo 1990; Mizuno \\etal 1995).  It is also well-recognized that the young stars are strongly correlated with gas and dust; i.e., that the stellar distribution must exhibit filamentary structure as well.  However, the extent of the stellar filamentary distribution, and its physical significance, has not been given sufficient attention, nor has it been viewed in the context of recent numerical simulations of molecular clouds. The spatial distribution of young objects in Taurus is reexamined from this point of view, with emphasis on its relationship to our previous suggestion that the Taurus cloud (like other nearby molecular clouds) was formed by large-scale flows in the interstellar medium (Ballesteros-Paredes, Hartmann, \\& Vazquez-Semadeni 1999, BHV; Hartmann, Ballesteros-Paredes, \\& Bergin 2001, HBB). ",
        "conclusions": "The spatial distribution of young stars in Taurus is remarkably filamentary and well-organized.  The large-scale coherence of the spatial distribution suggests that small-scale turbulence is not dominating the cloud structure, consistent with driving by external flows which could have formed the cloud. The spatial distribution of young stars is consistent with a roughly uniform distribution along bands and filaments, with a spacing comparable to the local Jeans length. The efficiency of star formation in filaments is much higher than elsewhere; the higher frequency of protostars and accreting T Tauri stars indicates that the population in these filaments is relatively young. The protostellar cores from which the stars formed are often elongated along the filaments in which they reside, consistent with formation by gravitational fragmentation. Fragmentation of filaments naturally produces elongated, roughly prolate cloud cores whose structure is well-suited to producing binaries after collapse.  Numerical simulations of turbulent clouds with higher spatial resolution, and turbulent driving on large scales, are urgently needed to compare with the observations of this nearby region which constitutes the main example of ``quiescent'' star formation.  The fragmentation of filaments into stars in Taurus has an echo on larger scales in regions of high mass and clustered star formation, a topic addressed in the second paper in this series.  The discussion of this paper was strongly influenced by my collaboration with Javier Ballesteros, especially the many discussions we have had concerning turbulent support.  The paper also benefited greatly from the suggestions, comments, and encouragement of Phil Myers, who also suggested the importance of the roughly periodic banded/filamentary structure of Taurus, and its implications in terms of fragmentation from a sheet-like structure. This work was supported in part by NASA grant NAG5-9670."
    },
    "0207/astro-ph0207093_arXiv.txt": {
        "abstract": "Interferometric observations of stars in late stages of stellar evolution and the impact of VLTI observations are discussed. Special attention is paid to the spectral information that can be derived from these observations and on the corresponding astrophysical interpretation of the data by radiative transfer modelling. It is emphasized that for the robust and non-ambiguous construction of dust-shell models it is essential to take diverse and independent observational constraints into account. Apart from matching the spectral energy distribution, the use of spatially resolved information plays a crucial role for obtaining reliable models. The combination of long-baseline interferometry data with high-resolution single-dish data (short baselines), as obtained, for example, by bispectrum speckle interferometry, provide complementary information and will improve modelling and interpretation. ",
        "introduction": "The Very Large Telescope Interferometer (VLTI; see Glindemann, this volume) of the European Southern Observatory with its four 8.2\\,m unit telescopes (UTs) and three 1.8\\,m auxiliary telescopes (ATs) will certainly establish a new era of studying the late stages of stellar evolution within the next few years. With a maximum baseline of up to more than 200\\,m, the VLTI will allow observations with unprecedented resolution opening up new vistas to a better understanding of the physics of evolved stars and thus of stellar evolution.  Stars in late stages of stellar evolution form therefore an important group among the VLTI key targets. During the Red Giant phase, strong  winds erode the stellar surfaces leading to the formation of circumstellar shells which absorb an increasing fraction of the visible light and re-emit it in the infrared regime. Accordingly, most of these evolved stars are bright infrared objects. The heavy mass loss leads to the chemical enrichment of the interstellar medium and therefore plays a crucial role for the understanding of the galactic chemodynamical evolution.  The vast majority of all stars, which have left their main sequence phase and become Red Giants, are of low and intermediate mass and finally evolve along the Asymptotic Giant Branch (AGB). These luminous, frequently pulsating and heavily mass-losing AGB stars form an important stellar population which contributes considerably to light, chemistry and dynamics of galaxies. The envelopes of AGB stars are the major factories of cosmic dust. Accordingly, AGB stars are often heavily enshrouded by dust exposing high fluxes in the infrared and are ideal laboratories to investigate the interplay between various physical and chemical processes. Most dust shells around AGB stars are known to be spherically symmetric on larger scales, whereas most objects in the immediate successive stage of proto-planetary nebulae appear in axisymmetric geometry. Evidence is growing that this break of symmetry takes place already at the very end of the AGB evolution.  Mass loss is also one of the dominant effects during the evolution of massive stars, virtually leading to an almost complete peeling of the star. Circumstellar dust shells found around evolved massive supergiants often show features of non-spherical outflows. Observing and modelling the circumstellar shells surrounding these stars, unveil details of evolution as, for instance, mass-loss rates. The presence of fossil shells even gives clues for the evolutionary history.  Dust formation around evolved stars can even continue beyond the Red Giant stage, as, e.g., in R\\,CrB stars or late-type Wolf-Rayet stars. The production of dust in such hostile environments is still challenging to theory. In the instance of Wolf-Rayet stars colliding winds due to binarity is one of the favored scenarios.  High-resolution interferometric observations reveal details of  disks and dust shells of evolved stars and thus improve our knowledge % of, for example, the mass-loss process and its evolution. In the following sections, we discuss high spatial resolution observations and their interpretation by radiative transfer calculations for some prominent evolved stars. ",
        "conclusions": ""
    },
    "0207/astro-ph0207570_arXiv.txt": {
        "abstract": "We report sensitive ATCA radio continuum observations toward IRAS 15596$-$5301 and 16272$-$4837, two luminous objects (${\\cal L}>2\\times10^4$\\Lsun) thought to represent massive star forming regions in early stages of evolution (due to previously undetected radio emission at the 1$\\sigma$ level of 2 mJy per beam). Also reported are 1.2 millimeter continuum and a series of molecular line observations made with the SEST telescope.  The radio continuum observations toward IRAS 15596$-$5301 reveal the presence of three distinct compact sources, with angular sizes of 2.7\\arcsec\\ to 8.8\\arcsec\\ (FWHM), all located within a region of 30\\arcsec\\ in diameter. Assuming that these are regions of ionized gas, we find that they have diameters of 0.06-0.2 pc, electron densities of $8\\times10^2 - 2\\times10^3$ cm$^{-3}$, and that they are excited by early B type stars. The 1.2-mm observations show that the dust emission arises from a region of $42\\arcsec\\times25\\arcsec$ (FWHM) with a total flux of 5.8 Jy, implying a mass of $1.4\\times10^3$ \\Msun. The line observations indicate that IRAS 15596$-$5301 is associated with a molecular cloud with a FWHM angular size of 37\\arcsec\\ ($\\sim0.4$ pc radius at the distance of 4.6 kpc), a molecular hydrogen density of $\\sim4\\times10^5$ cm$^{-3}$ and a rotational temperature of $\\sim27$ K. We suggest that the massive dense core associated with IRAS 15596$-$5301 contains a cluster of B stars which are exciting compact \\hii\\ regions that are in pressure equilibrium with the dense molecular surroundings.  No radio continuum emission was detected from IRAS 16272$-$4837 up to a $3\\sigma$ limit of 0.2 mJy. However, the 1.2-mm observations show strong dust emission arising from a region of $41\\arcsec\\times25\\arcsec$ (FWHM) with a total flux of 13.8 Jy, implying a mass of $2.0\\times10^3$ \\Msun. The line observations indicate the presence of an elongated molecular cloud with FWHM major and minor axes of 61\\arcsec\\ and 42\\arcsec ($0.50\\times0.35$ pc in radius at the distance of 3.4 kpc), a molecular hydrogen density of $\\sim2\\times10^5$ cm$^{-3}$ and a rotational temperature of $\\sim27$ K. The high luminosity ($2.4\\times10^4$\\Lsun) and lack of radio emission from this massive core suggest that it hosts an embedded young massive protostar that is still undergoing an intense accretion phase. This scenario is supported by the observed characteristics of the line profiles and the presence of a bipolar outflow detected from observations of the SiO emission. We suggest that IRAS 16272$-$4837 is a bona-fide massive star forming region in a very early evolutionary stage, being the precursor of an ultra compact \\hii\\ region. ",
        "introduction": "The earliest phase of high-mass star formation is possibly one of the least understood stage of evolution of massive stars. Massive stars (M$>8$ \\Msun) are known to be formed in dense molecular cores, however the sequence of processes leading to their formation is not yet well established. In particular, the role of coalescence (Stahler et al. 2000) and accretion (Osorio, Lizano, \\& D'Alessio 2000) processes in the assembling of a massive star is still under debate. The determination of the physical conditions of the gas during the formation and early evolution of a massive star is difficult because of their scarcity and rapid evolution. In addition, massive stars are usually born in clusters or groups hence their individual studies are usually afflicted by confusion, particularly because they are found located, on the average, at larger distances from the Sun than sites of low-mass star formation. Massive objects in early evolutionary stages, namely in the process of assembling the bulk of their final mass, should be identified by having high bolometric luminosities ($>10^4 $\\Lsun), strong dust emission, and very weak or no detectable free-free emission at cm radio wavelengths. The bolometric luminosity has contributions from the accretion of infalling material and nuclear burning. Up to date only a few systematic searches for high mass protostellar objects have been carried out (Molinari et al. 1996, 1998, 2000; Sridharan et al. 2002).  We have recently started a multi-wavelength study of a sample of 18 luminous IRAS sources in the southern hemisphere thought to be representative of young massive star forming regions (Mardones, Garay, \\& Bronfman 2002). The goal is to understand the physical and chemical differences between different stages of early evolution. The objects were taken from the Galaxy-wide survey of CS(2$\\rightarrow$1) emission towards IRAS sources with IR colors typical of compact \\hii\\ regions (Bronfman, Nyman, \\& May 1996). We selected sources based primarily on the observed CS(2$\\rightarrow$1) line profiles; looking for self-absorbed lines consistent with inward or outward motions (e.g., Mardones 1998), and/or with extended line wings, possibly indicating the presence of bipolar outflows. In addition, the sources were required to have IRAS 100$\\mu$m fluxes greater than $10^3$ Jy and to be in the southern hemisphere ($\\delta < -20\\arcdeg$). The luminosity of the IRAS sources, computed using the IRAS energy distribution and the distances derived by Bronfman (2002) are in the range $2\\times10^4 - 4\\times10^5$ \\Lsun, implying that they contain at least an embedded massive star.  Most of the selected objects are expected to be associated with ultra compact (UC) \\hii\\ regions which are thought to be manifestations of newly formed massive stars that are still embedded in their natal molecular cloud. This expectation is confirmed by the radio continuum observations of Walsh et al. (1998) which show that 9 of the 12 sources in both samples have detectable radio continuum emission (above a $3\\sigma$ limit of 6 mJy/beam at 8.64 GHz with angular resolution of $\\sim$1.5\\arcsec). The objects that were not detected at radio wavelengths are suitable candidates for massive stars in very early stages of evolution in which dense material is still falling toward a massive OB-type protostar. In this accretion phase, the high-mass accretion rate of the infalling material quenches the development of an UC \\hii\\ region (Yorke 1984; Walmsley 1995), and the free-free emission from the ionized material is undetectable at centimeter wavelengths. The mass accretion rate might also be large enough that the ram pressure of the infalling gas could provide the force to prevent the expansion of an \\hii\\ region. We note, however, that due to the limited sensitivity of the Walsh et al. (1998) survey, low-density \\hii\\ regions with emission measures smaller than $4\\times10^5$ pc cm$^{-6}$ were not detectable. Hence, the lack of detection at the above limit does not rule out the presence of an optically thin compact \\hii\\ region.  In this paper we report sensitive ATCA radio continuum observations toward two sources in our sample, IRAS 15596$-$5301 and 16272$-$4837, without previously detected radio continuum emission to place stringent limits in their radio flux density. The $1\\sigma$ sensitivity level of 70 $\\mu$Jy at 4.8 GHz achieved in our observations is thirty times smaller than in previous studies and is sufficient to detect the emission measure corresponding to any ionizing OB star within the Galaxy. The main goal was to establish whether or not these objects correspond to very young massive objects; that is massive protostars which have not yet ionized significant amounts of the surrounding gas. Also reported in this paper are millimeter continuum and molecular line observations of IRAS 15596$-$5301 and 16272$-$4837 made with the Swedish-ESO submillimetre telescope. The latter observations are part of a molecular line survey toward several high-mass star forming regions made in order to determine their physical characteristics and investigate possible chemical differences. ",
        "conclusions": "\\subsection{Spectral energy distribution}  Figure~\\ref{fig-sed} shows the spectral energy distribution (SED) of IRAS 15596$-$5301 and 16272$-$4837 from 12 $\\mu$m to 1.2 mm, which is mainly due to thermal dust emission. We fitted the SED with modified blackbody functions of the form $ B_{\\nu}(T_d)\\left[1-\\exp(-\\tau_{\\nu})\\right]\\Omega_s~, $ where $\\tau_{\\nu}$ is the dust optical depth, $B_{\\nu}(T_d)$ is the Planck function at the dust temperature $T_d$, and $\\Omega_s$ is the solid angle subtended by the dust emitting region. The opacity was assumed to vary with frequency as $\\nu^{\\beta}$, i.e. $\\tau_{\\nu}= \\left(\\nu/\\nu_o\\right)^{\\beta}$, where $\\nu_o$ is the frequency at which the optical depth is unity. Due to the limited number of spectral points we have set the value of $\\beta$ equal to 2.0, consistent with tabulated opacities (Ossenkopf \\& Henning 1994) and derived values for high mass star forming regions (Molinari et al. 2000). A single temperature model produced poor fits, underestimating the emission observed at wavelengths smaller than 25$\\mu$m, and therefore we used a model with two temperature components.  From the fits (long-dashed lines) we derive that the colder dust component (short-dashed lines) toward 15596$-$5301 and 16272$-$4837 have, respectively, temperatures of 27 and 25 K, angular sizes (assuming a Gaussian flux distribution) of 30\\arcsec (FWHM), and wavelengths at which the opacity is unity of $\\sim$90 $\\mu$m and 140 $\\mu$m. The temperature of the hot dust component is 100 K for 15596$-$5301 and 115 K for 16272$-$4837. The thermal dust emission at 1.2 mm is therefore optically thin ($\\tau_{1.2mm}\\sim 5\\times10^{-3}$), and thus the observed flux density at 1.2 mm allows to obtain an additional mass estimate of the dense cores. In general, for an isothermal dust source the total gas mass, $M_{g}$, is given in terms of the observed flux density, $S_{\\nu}$, at an optically thin frequency, $\\nu$, by (e.g. Chini, Krugel, \\& Wargau 1987) $$ M_{g} = {{S_{\\nu} D^2}\\over{R_{dg} \\kappa_{\\nu} B_{\\nu}(T_d)}} ~~, $$ where $\\kappa_{\\nu}$ is the mass absorption coefficient of dust, $R_{dg}$ is the dust-to-gas mass ratio (assuming 10\\% He), and $B_{\\nu}(T_d)$ is the Planck function at the dust temperature $T_d$. The main source of uncertainty in the conversion of the observed flux density into gas mass is the $R_{dg}\\kappa_{\\nu}$ factor, or total mass opacity, which is a poorly known quantity (e.g. Gordon 1995). Using a dust opacity at 1.2 mm of 1 cm$^2$ g$^{-1}$, as computed by Ossenkopf \\& Henning (1994) for dense and cold protostellar cores, $R_{dg}=0.01$, the fitted dust temperatures, and the observed flux densities, we derive masses of $1.4\\times10^3$ \\Msun\\ for IRAS 15596$-$5301 and $2.0\\times10^3$ \\Msun\\ for IRAS 16272$-$4837. These masses derived from the dust emission are in good agreement with those derived from the molecular line intensities and from the virial assumption.  \\subsection{Evolutionary stages }  \\subsubsection{IRAS 15596$-$5301 (G329.40-0.46)}  The radio continuum observations toward IRAS 15596$-$5301 indicate the presence, in a region of $\\sim$0.3 pc in radius, of three distinct \\hii\\ regions with diameters ranging from 0.06 to 0.2 pc. The multiple structure of the ionized gas is typical of galactic \\hii\\ regions, and is most likely due to the presence of a cluster of exciting stars. If components A, B, and C are excited by individual ZAMS stars, the rate of UV photons needed to ionize them (see Table 3) imply exciting stars with spectral types of B0, B0.5, and B1, respectively. The total luminosity emitted by this cluster of B stars, as inferred from the radio observations, is 4.1$\\times10^4$ \\Lsun. On the other hand, the total far-infrared luminosity computed using the IRAS fluxes (see Casoli et al. 1986) is $\\sim6.5\\times10^4$ \\Lsun\\ (assuming a distance of 4.6 kpc; Bronfman 2002). The difference between the radio derived luminosity and the IRAS luminosity could be explained by the presence of dust within the \\hii\\ regions. Garay et al. (1993) found that the fraction of Lyman continuum photons absorbed by dust within \\hii\\ regions is typically~55\\%. Alternatively, it could be explained by the presence, in addition to the B stars, of several less massive stars that will contribute to the FIR luminosity but that are not hot enough to contribute to ionization. It is not easy, however, to disentangle which of these effects play the predominant role.  The compact \\hii\\ regions are found projected toward the peak of the CS($5\\rightarrow4$) emission map (see Figure~\\ref{fig-cs54maps}), suggesting that they are deeply embedded within the dense molecular core. From the observed sizes, and assuming a sound speed in the ionized gas of 11.4 \\kms, we estimate that the \\hii\\ regions have dynamical ages between 3$\\times 10^3$ to 8$\\times 10^3$ yrs. If these correspond to the actual ages of the compact \\hii\\ regions, then we should conclude that they are very young objects. The dynamical time-scales, however, may not provide a realistic estimate of the actual age of \\hii\\ regions. The large number of UC \\hii\\ regions and their short dynamical ages poses the well known problem that the rate of massive star formation appears to be much greater than other indicators suggest (Wood \\& Churchwell 1989, Churchwell 1990). Due to the high density of the molecular gas in which they are embedded, we suggest instead that the \\hii\\ regions within the IRAS 15596$-$5301 massive core might be in pressure equilibrium with the surrounding dense ambient medium, and are currently stalled at their equilibrium radius. (e.g. De Pree, Rodr\\'\\i guez, \\& Goss 1995). The molecular density of the ambient gas needed to stall an \\hii\\ region at radius, $R_f$, is given by (e.g. Garay \\& Lizano 1999) $$ \\left({{n_{H_2}}\\over{10^{5}~\\rm cm^{-3}}}\\right) =  1.9 \\left({{N_{\\rm u}}\\over{10^{49}~\\rm s^{-1}}}\\right)^{1/2} \\left({{T_e}\\over{10^4~\\rm K}}\\right) \\left({{30~\\rm K}\\over{T_o}}\\right) \\left({{\\rm pc}\\over{R_f}}\\right)^{3/2}~~, $$ where $N_u$ is the rate of ionizing UV photons emitted by the exciting star, $T_e$ is the electron temperature of the ionized gas, and $T_o$ is the temperature of the ambient gas. Using the observed radius and the derived ionizing rate of UV photons of the \\hii\\ regions within IRAS 15596$-$5301, we find that molecular densities of $\\sim5\\times10^5$ cm$^{-3}$ are needed for them to be pressure confined by the dense environment. These densities are similar to those derived from the molecular observations. The time needed for the \\hii\\ regions to achieve pressure equilibrium are between $1.0\\times10^5$ yrs to $2.5\\times10^5$ yrs, implying that massive star formation started within this core more than $2.5\\times10^5$ yrs ago. We conclude that the dense massive core is in an advanced stage of early evolution, in which multiple OB star formation have already taken place near its central region.  \\subsubsection{IRAS 16272$-$4837 (G335.58-0.28)}  The total far-infrared luminosity of IRAS 16272$-$4837 computed using the IRAS fluxes (see Casoli et al. 1986) is $\\sim2.4\\times10^4$ \\Lsun\\ (assuming a distance of 3.4 kpc; Bronfman 2002). The luminosity obtained integrating under the fitted curve in Figure~\\ref{fig-sed} is similar to the IRAS luminosity. The high luminosity suggests that the IRAS 16272$-$4837 massive core hosts a young massive protostar inside; and the lack of radio emission suggests that it is still undergoing an intense accretion phase. Models of massive envelopes accreting onto a young massive central B type star (e.g. Osorio et al. 2000) require accretion rates in the envelopes $\\dot M >5\\times10^{-4} M_{\\odot}$ yr$^{-1}$ in order to fit the observed SEDs. The high-mass accretion rate of the infalling material quenches the development of an UC \\hii\\ region (Yorke 1984), and the free-free emission from the ionized material is undetectable at centimeter wavelengths.  This hypothesis is supported by the characteristics of the line profiles observed toward IRAS 16272$-$4837, which suggest that the molecular gas is undergoing infalling motions. From the spectra of the optically thick HCO$^+(1\\rightarrow0)$ line we measure a velocity difference between the red and blue peaks of 2.7 \\kms, and brightness temperature of the blue peak, red peak, and dip of 5.4, 3.6, and 2.9 K, respectively. From the spectra of the optically thin H$^{13}$CO$^+(1\\rightarrow0)$ line we measure a FWHM line width of 3.18 \\kms. From these values, using the simple model of contracting clouds of Myers et al. (1996), we derive a characteristic inward speed of 0.5 \\kms. We note that this value is considerably smaller than the free-fall velocity expected for a cloud with a total mass of $\\sim3\\times10^3$\\Msun\\ at its outer envelope radius of 0.4 pc, suggesting that the collapse is not dynamical. Using the derived values of the infall speed, molecular density, and core size, we obtain a mass infall rate, $\\dot{M}_{in}$, of $1\\times10^{-2} M_{\\odot}$ yr$^{-1}$, large enough to prevent the development of an UC \\hii\\ region. The high value of the mass infall rate rises the question as to which is the fraction of the total luminosity due to accretion.  The accretion luminosity, $L_{acc}$, is $$ L_{acc} = {{G~ f~ \\dot{M}_{in}~ M_p}\\over{R_p}}~~,$$ where $f$ is the fraction of the large scale mass infall rate that goes into accretion onto the protostar, $M_p$ is the mass of the protostar, and $R_p$ the radius where the shock occurs. None of these three parameters are known for IRAS 16272$-$4837. Assuming $M_p\\sim10$\\Msun, $R_p\\sim3\\times10^{12}$ cm, and $f\\sim0.05$ (e.g, Norberg \\& Maeder 2000) we obtain $L_{acc}\\sim3.6\\times10^3$ \\Lsun, about 15\\% of the total luminosity. We emphasize that this value of the accretion luminosity corresponds only to a rough estimate, particularly because the value of $f$ is highly uncertain.  Additional evidence for IRAS 16272$-$4837 to be in a collapsing stage is provided by the presence of bipolar outflowing gas, phenomenon which is thought to be closely related to accretion processes. The 22$\\mu$m object lies at the center of symmetry of the SiO outflow, suggesting it is intimately associated with the energy source of the outflow. The early evolutionary stage of this region is also sustained by the presence of 6.67 GHz methanol masers, which are thought to be signposts of young regions of massive star formation (Walsh et al. 1997, 1998). Notice that the 22$\\mu$m source is elongated, with the 6.67 GHz methanol masers being aligned along its major axis.  Finally, we mention that the high value of the mass to luminosity ratio of IRAS 16272$-$4837, $M/L = 0.083$, about 4 times higher than that of IRAS 15596$-$5301, is another indicator of its youth, as argued by Sridharan et al (2002)."
    },
    "0207/astro-ph0207436_arXiv.txt": {
        "abstract": "We examine the propagation of two-dimensional relativistic jets through the stellar progenitor in the collapsar model for gamma-ray bursts (GRBs). Each jet is parameterized by a radius where it is introduced, and by its initial Lorentz factor, opening angle, power, and internal energy. In agreement with previous studies, we find that relativistic jets are collimated by their passage through the stellar mantle. Starting with an initial half-angle of up to 20 degrees, they emerge with half-angles that, though variable with time, are around 5 degrees. Interaction of these jets with the star and their own cocoons also causes mixing that sporadically decelerates the flow.  We speculate that this mixing instability is chiefly responsible for the variable Lorentz factor needed in the internal shock model and for the complex light curves seen in many gamma-ray bursts. In all cases studied, the jet is shocked deep inside the star following a brief period of adiabatic expansion. This shock converts most of the jet's kinetic energy into internal energy so that even initially ``cold'' jets become hot after going a short distance.  The jet that finally emerges from the star thus has a moderate Lorentz factor, modulated by mixing, and a very large internal energy.  In a second series of calculations, we follow the escape of that sort of jet. Conversion of the remaining internal energy gives terminal Lorentz factors along the axis of approximately 150 for the initial conditions chosen. Because of the large ratio of internal to kinetic energy in both the jet ($\\geq 80\\%$) and its cocoon, the opening angle of the final jet is significantly greater than at breakout. A small amount of material emerges at large angles, but with a Lorentz factor still sufficiently large to make a weak GRB. This leads us to propose a ``unified model'' in which a variety of high energy transients, ranging from x-ray flashes to ``classic'' GRBs, may be seen depending upon the angle at which a standard collapsar is observed.  We also speculate that the breakout of a relativistic jet and its collision with the stellar wind will produce a brief transient with properties similar to the class of ``short-hard'' GRBs.  Implications of our calculations for GRB light curves, the luminosity-variability relation, and the GRB-supernova association are also discussed. ",
        "introduction": "Growing evidence connects GRBs to the death of massive stars. Analysis by \\citet{fra01} and others of the radio afterglows of ``long-soft'' GRBs suggests that, despite great diversity in apparent brightnesses, these events have a total kinetic energy in relativistic matter tightly clustered around $3 \\times 10^{51}\\,\\erg$ \\citep{fra01} - a supernova-like energy.  ``Bumps'' resembling the light curves of Type I supernovae have also been seen in the optical afterglows of at least four GRBs (GRB 980326, Bloom et al. 1999; GRB 970228, Reichart 1999, Galama et al. 2000; GRB 011121, Bloom et al. 2002, Garnavich et al. 2002; and GRB 020405, Price et al. 2002) and GRB 980425 has been associated with an optical supernova, SN 1998bw (e.g., Galama et al. 1998; Iwamoto et al. 1998; Woosley, Eastman, \\& Schmidt 1999). The observational evidence that (long-soft) GRBs are associated with regions of star formation has become overwhelming \\citep{blo02b}. Given the evidence for beaming and relativistic motion, it is probable that at least one major subclass of GRBs is a consequence of massive stars that, in their explosive deaths, produce relativistic jets.  Here we examine the passage of relativistic jets through a collapsing massive star and their breakout. We begin with a rotating star that has already experienced 10 seconds of collapse. Ten seconds is the nominal time for a disk to form around a black hole and the polar region to become sufficiently evacuated for a polar jet to propagate \\citep{mac99}. The initial progenitor is a helium core of $15\\,\\Msun$ evolved by \\citet{heg02} to iron core collapse with approximations to all (non-magnetic) forms of angular momentum transport included. The missing 10 seconds is followed using a non-relativistic two-dimensional code as in \\citet{mac99}.  We pick up the calculation when the jet, which presumably began in a region $\\sim 30\\,\\km$ in size, has already reached a radius of $2000\\,\\km$ and do not consider what has gone on inside. For present purposes, details of whether the jet was made by black hole angular momentum, MHD processes in the disk, or neutrinos do not concern us. The jet is initiated in a parametric way based upon its power, opening angle, Lorentz factor, and internal energy. Its propagation to the stellar surface at $800,000\\,\\km$, and its interaction with the stellar mantle is then followed (\\S~\\ref{jin}). Additional calculations (\\S~\\ref{jout}) also examine what happens to the jet immediately after it escapes the star and converts its residual internal energy into additional relativistic motion.  We find, in agreement with \\citet{alo00}, that the passage of the jet through the star leads to its additional collimation.  We also find that instabilities along the beam's surface lead to mixing with the nearly stationary stellar material and cocoon. The mixing produces variations in the mass loading, and therefore the Lorentz factor of the jet. The opening angle of the jet also varies with time, gradually, if irregularly, growing as the star is blown aside. These results, discussed in \\S~\\ref{obs}, have important implications for the observed light curves and energies of GRBs and imply that what is seen may vary greatly with viewing angle. In particular, we predict the existence of a large number of low energy GRBs with mild Lorentz factors (\\S~\\ref{unified}) that may be related to GRB 980425/SN 1998bw and to the recently discovered ``hard x-ray flashes'' \\citep{hei01}.  Finally, we consider the breakout of the jet. As the shock breaks out a small amount of material being pushed ahead by the jet head is accelerated to relativistic speeds. Interaction of this material with the stellar wind of the progenitor will produce a transient of some sort \\citep{woo99a, tan01}.  We speculate that this is the origin of a hard precursor to GRBs that, at least in some cases, might be characterized as a ``short-hard GRB'' in isolation(\\S~\\ref{shb}). ",
        "conclusions": ""
    },
    "0207/astro-ph0207600_arXiv.txt": {
        "abstract": "We present a detailed investigation of issues related to the measurement of peculiar velocities and temperatures using Sunyaev-Zel'dovich (SZ) effects. We estimate the accuracy to which peculiar velocities and gas temperatures of distant galaxy clusters could be measured. With $\\mu K$ sensitivity on arcminute scales at several frequencies it will be possible to measure peculiar velocities to an accuracy of $\\sim$130 km\\,s$^{-1}$ and gas temperatures to better than 1 keV. The limiting factor for the accuracy of $v_{pec}$ is the presence of bulk motions within the galaxy cluster, even for apparently relaxed clusters.  The accuracy of the temperature is mainly limited by noise. These results are independent of redshift. Such constraints can best be achieved with only three frequencies: one in the Rayleigh-Jeans region ($\\nu<40$ GHz), one near 150 GHz, and the third at 300 GHz or higher. Measurements at the null of the thermal SZ effect are of marginal utility, other than as a foreground/background monitor. ",
        "introduction": "\\label{sec:intro}  Observations of the Sunyaev-Zel'dovich (SZ) effect \\citep{sunyaev72} are currently at a mature stage. Highly significant detections are now routine and the next generation of instruments is about to exploit the SZ effect to provide deep surveys of galaxy clusters (for a recent review see Carlstrom {\\em et~al.} 2002 \\nocite{carlstrom02}). All current and near-future instruments are devoted to studies of the {\\em thermal} SZ effect. There are a host of yet more subtle distortions of the cosmic microwave background (CMB) spectrum that contain a wealth of information related to the cluster's peculiar velocity ($v_{pec}$) and the gas temperature ($T$) of the intracluster medium.  In this paper we investigate the most significant of the subtler distortions, namely those due to the line-of-sight motion (the kinetic SZ effect) and the distortion of the spectrum due to relativistic effects (the relativistic thermal SZ effect). These effects are discussed in several reviews \\citep{rephaeli95,birkinshaw99} and are presented in great detail by several authors \\citep{sunyaev80,rephaeli95,challinor98,sazonov98,itoh98,nozawa98b,molnar99, dolgov01}.  At sensitivities below $\\sim$1 $\\mu K$ some higher order effects could be important such as multiple scatterings and transverse velocities. At such low levels, CMB anisotropies will be also be a very difficult contaminant. For simplicity, we restrict ourselves to the most significant effects.  The importance of using the SZ effect to its fullest potential comes from its redshift independence. As a spectral distortion of the CMB, it redshifts along with the CMB and the amplitude of the distortion does not suffer from cosmological dimming. If there are clusters at redshift $z \\sim 2$ (as there should be in standard models of structure formation) they will be very faint in X-rays, and perhaps undetectable if the gas has been preheated by an early burst of star formation. An independent probe of the gas temperature would be invaluable, and preliminary steps along these lines are being made \\citep{hansen02}. Measurements of peculiar velocities would allow reconstruction of the large scale density field on scales comparable to the horizon \\citep{dore02}. A firm understanding of the large scale density inhomogeneities would allow new tests of galaxy formation and provide a view of the evolution of structure.  In the next section we lay out the physical effects which allow such powerful tests. In \\S\\ref{sec:freqs} we investigate the observing frequencies which are best suited for such an investigation. The details of the simulations and map-making methods are outlined in \\S\\ref{sec:sims} and results are presented in \\S\\ref{sec:results}. We close with a discussion of practical issues and a summary. ",
        "conclusions": "\\label{sec:disc}  Measurements of peculiar velocities of galaxy clusters at microwave frequencies will soon be possible. We have shown that multi-frequency, sensitive observations could measure peculiar velocities to an accuracy of roughly 100 km/s. Measurements of gas temperatures will be useful, with uncertainties possibly smaller than could be achieved with X-ray spectroscopy. The redshift independence of the SZ effect will make this an extremely powerful tool for studies of distant clusters. Exploring differences between X-ray emission weighted and SZ emission weighted temperature maps will no doubt be interesting.  Foregrounds and backgrounds will be important barriers to such precise studies of peculiar velocities. The primary anisotropies of the CMB will become problematic on scales larger than a few arcminutes, and point source removal  will be very difficult. It has not yet been demonstrated that the atmosphere will not be a problem for ground-based experiments, but there is no {\\em a priori} evidence that there will be a problem. Inter-frequency calibration to the requisite precision will be a significant technical challenge.  Point sources come in (at least) two varieties. At frequencies below $\\sim$ 90 GHz radio point sources (primarily extragalactic AGN but also star-forming galaxies) have historically been a problem for SZ measurements and they are unlikely to go away. The best solution seems to be simultaneous monitoring at extremely high (a few arcseconds) resolution. At higher frequencies, dusty star-forming galaxies are ubiquitous. If no source subtraction is done, confusion could easily be at the level of 10 $\\mu K$, comparable to the kinetic SZ signal \\citep{blain98}. Currently very little is known about these sources, making spectral subtraction (measuring at a higher frequency where the SZ signal is negligible) difficult; the ultimate solution  may require something like ALMA to remove the point sources at mm wavelengths. There is no evidence for variability in these sources, so it will not be necessary to do the subtraction simultaneously, as is required for the occasionally variable radio point sources.  The best frequencies for observation turn out to not include the null of the thermal SZ effect. The best strategy is to have a Rayleigh- Jeans band (below 90 GHz), a high frequency band (above 300 GHz) and a band near 150 GHz.  Much has been made of the null of the thermal SZ effect, and it will be important as a check for systematic errors, but it is not a good frequency for cluster studies.  Bulk velocities within the cluster, combined with contamination from the anisotropies of the CMB, lead to a limit of roughly 100 km/s on the possible accuracy of kinetic SZ velocity measurements. This is much higher than what would be expected from considerations of the background noise levels and the few tens of km/s that arise from the difficulty in choosing the appropriate definition of velocity. More work on simulations could shed significant light on optimal strategies for estimating the true peculiar velocity as well as provide a much better estimate of the distribution of errors that could be expected from an ensemble of galaxy cluster peculiar velocity measurements.  Measurements at cm and mm wavelengths are opening a new window on cosmology. It will soon be possible to measure gas temperatures and peculiar velocities to good accuracy out to $z\\sim 2$, allowing unprecedented tests of structure formation as well as an excellent understanding of the topography of much of the observable universe."
    },
    "0207/astro-ph0207526_arXiv.txt": {
        "abstract": "We extend the analysis of Gabor transforms on a Cosmic Microwave Background (CMB) temperature map \\cite{gabortrans} to polarisation. We study the temperature and polarisation power spectra on the cut sky, the so-called pseudo power spectra. The transformation kernels relating the full-sky polarisation power spectra and the polarisation pseudo power spectra are found to be similar to the kernel for the temperature power spectrum. This fact is used to construct a fast power spectrum estimation algorithm using the pseudo power spectrum of temperature and polarisation as data vectors in a maximum likelihood approach. Using the pseudo power spectra as input to the likelihood analysis solves the problem of having to invert huge matrices which makes the standard likelihood approach infeasible. ",
        "introduction": "Most theories of the early universe predict the temperature and polarisation fluctuations of the CMB to be Gaussian distributed. In such models, the angular temperature and polarisation power spectra contain all the information about the cosmological parameters which one can determine from observations of the CMB sky. As several combinations of the cosmological parameters can give rise to similar temperature power spectra, estimating the polarisation power spectra will break the degeneracy and will be of great importance for accurate estimation of cosmological parameters. Also the error bars on these parameters can be reduced by exploiting the extra information present in the CMB polarisation power spectra.\\\\  Much effort has been made recently in order to find methods to analyse the CMB temperature power spectrum \\cite{gabortrans,OhSpergelHinshaw,pseudo,ringtorus1,ringtorus2,bond,BJK,bartlett,tegmark,dore,szapudi,master,amad}. There has been very few publications confronting the even harder task of estimating the polarisation power spectra. The framework for analysing the polarisation power spectra has been set in \\cite{pol1,pol2} but these papers only describe the full likelihood method which is far too time consuming also when only considering the temperature power spectrum. In \\cite{tegmark0} a quadratic polarisation power spectrum estimation method was introduced, similar to the one presented in \\cite{tegmark} for temperature only.\\\\  In this paper we will extend the method of using the pseudo power spectrum as input to a likelihood estimation procedure of the power spectrum as described in \\cite{gabortrans} (from now on called HGH). We will include the $E$ and $C$ mode polarisation pseudo power spectra in the data vector and use techniques similar to those described in HGH to estimate the power spectra. This can be done because the kernels that connect the full sky polarisation power spectra with the polarisation pseudo power spectra on an apodised sky are similar to the kernel for the temperature power spectrum. In the first part of this paper we will derive the formulae for these kernels and for the polarisation pseudo power spectra and discuss their shapes. Then in the second part this will be used for likelihood estimation.\\\\  In this paper the $B$ component polarisation will mostly be neglected. The $B$ polarisation power spectrum is expected to be very small and will hardly be detectable by the upcoming $MAP$ or $Planck$ satellite experiments \\cite{jaffe0}. Also the $E$ and $B$ components of polarisation mix on the cut sky as will be discussed in this paper, making the $B$ polarisation pseudo spectrum to be dominated by the $E$ component \\cite{sepeb1,sepeb2,tegmark0,sepeb3}. ",
        "conclusions": "We have presented a maximum likelihood method to simultaneously estimate the temperature and polarisation power spectra from high resolution CMB data in the presence of non-uniform noise and a symmetric Gabor window. An extension of the power spectrum estimation method developed in HGH has been made in order to estimate for the polarisation power spectra in addition to the temperature power spectrum. In most standard theories for the early universe, the $B$ component polarisation will be too small to be observed by the MAP and Planck experiments. For this reason, the method has been tested here under the assumption that the $B$ mode polarisation is negligible. In this case the method appears to give unbiased estimates of the polarisation power spectra also in the presence of non-uniform noise and a Gabor window.\\\\  The kernels connecting the full sky polarisation power spectra with the cut sky polarisation pseudo power spectra were studied and found to be very similar to the kernel for the temperature power spectrum. For this reason the effect of a cut sky and a Gabor window on the polarisation power spectra is similar to the effect on the temperature power spectrum. This explains that the method of estimating the power spectrum from the pseudo power spectrum for polarisation was as successful as it was for temperature.\\\\  One issue which has not been studied fully here is the inclusion of the $B$ mode polarisation. We demonstrated that the $E$ and $B$ mode polarisation power spectra are mixing on the cut sky making detections of the much weaker $B$ component difficult. Further work needs to be done in order to include the $B$ component in the likelihood analysis.\\\\  In HGH it was discussed how one can find the noise correlation matrix for temperature using Monte Carlo. This might be faster than the analytical approach presented here when the size of the dataset is very huge. With a sufficient number of Monte Carlo simulations this was shown to give similar error bars as the analytic treatment. The results for the temperature noise matrix is expected to be valid also for the polarisation noise matrices and can be used when the dataset is so big that the Monte-Carlo approach is significantly faster, or when correlated noise is present.\\\\  Another extension which was discussed in HGH was the simultaneous analysis of several patches on the CMB sky. This extension was shown to work for the temperature power spectrum and we expect that this could work also for polarisation, allowing data from several different experiments to be analysed together. In HGH it was shown that with this power spectrum estimation method, huge datasets like the ones to be expected from MAP or Planck can be analysed in a reasonable amount of time. The most time-consuming step in the method is the construction of the correlation matrix, which in the case of polarisation is 9 times longer. This makes also the joint temperature and polarisation power spectrum estimation feasible for huge datasets.   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=mergepclpolnoavg.ps,bbllx=0pt,bblly=0pt,bburx=696pt,bbury=842pt,height=18cm,width=22cm} \\caption{The result of a joint likelihood estimation of the temperature power spectrum (upper plot) and the $E$ (middle plot) and $C$ (lower plot) polarisation power spectra. The dotted line shows the full sky average spectrum. The histogram shows the binned input pseudo spectrum without noise. The shaded areas around the binned average full sky power spectrum (not shown) show the expected deviations from the average using the approximate formula for uniform noise. The bright shaded area shows the cosmic and sample variance only whereas the dark shaded area also shows expected variance due to noise. The dots show the estimate with $1\\sigma$ error bars taken from the inverse Fisher matrix. In the analysis a $15$ degree FWHM Gaussian Gabor window with a $\\theta_C=3\\sigma$ cutoff was used.} \\label{fig:polpcl} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=mergepclpol.ps,bbllx=0pt,bblly=0pt,bburx=696pt,bbury=842pt,height=18cm,width=22cm} \\caption{Same as figure (\\ref{fig:polpcl}) but the dots here are the average of 60 estimates from Monte Carlo simulations. The error bars are the average deviations taken from the simulations. The dotted line shows the average full sky spectrum. The shaded areas which are plotted around the binned full sky power spectrum (not shown) show the variance taken from the approximate variance formula for uniform noise. The dashed lines show the expected variance taken from the inverse Fisher matrix.} \\label{fig:polpclnoavg} \\end{center} \\end{figure}"
    },
    "0207/astro-ph0207460_arXiv.txt": {
        "abstract": "{The evidence for refractive interstellar scintillation (RISS) being the main cause for rapid intraday variations (Intraday Variability, IDV) in Quasars and BL Lacs has recently become stronger. If IDV is still a complex composition of extrinsic and source intrinsic effects, the intrinsic part of the IDV pattern should show up in the millimeter and sub-millimeter regime due to the frequency dependence of RISS. Hence, observations at higher frequencies are essential in order to exclude RISS as the sole cause of IDV. Here we report on our new attempt to search for rapid variations at much higher frequencies. In addition, the possibility of a direct detection of the scattering screen in front of IDV sources will be discussed. Our recent line observations towards a few IDV sources lead to the first detection of a high latitude molecular cloud in front of an intraday variable radio core. } \\authorrunning{Fuhrmann et al.} ",
        "introduction": "There is still a hot debate about the origin of short term radio variations in total as well as polarized flux of flat spectrum Quasars and BL Lacs. Such variations on time scales of hours to days (IntraDay Variability, IDV) are common in this type of objects and reveal the tiny dimensions of the emitting region. The cause of the variations seen in these sources is currently controversial with claims being made for either: 1) a source-intrinsic or 2) extrinsic origin due to scattering in the interstellar medium (ISM) (e.g. Wagner \\& Witzel 1995 and ref. therein). Observational findings suggest, that IDV is a complicated mixture of both effects (Krichbaum et al. 2002). Due to the involved small source sizes, refractive interstellar scintillation (RISS) must play an important role in the cm-radio regime, while intrinsic variations require extreme high Doppler boosting or special source geometries.\\\\ In this paper we will present our new observational approaches: In the first part we will report on the attempt to search for IDV at much higher frequencies than presently done (up to 345 GHz). In the second part we will introduce the possibility of a direct search for the required scattering screen and show first results of line observations towards IDVs. ",
        "conclusions": "During the first half of 2002 we started two new observational approaches: a search for high frequency IDV (32 GHz and 345 GHz) and the attempt to detect directly the possible foreground screen via spectral line observations. Since in both cases the data reduction is still in progress, preliminary results show: 1) if confirmed, at least the BL Lac object 0716+714 seems to show sub-mm IDV. This can not be explained by RISS and must be due to intrinsic mechanism. 2) a CO cloud in the direction of 0954+658. Such high latitude clouds could serve as the origin of the scattering material producing interstellar scintillation. This has to be confirmed by additional investigations of ionized material, like ionized HCO or the detection of RRLs. A further option could be faraday rotation measures on the line of sight to the source compared with those measured a few degrees away."
    },
    "0207/astro-ph0207183_arXiv.txt": {
        "abstract": "We present a simulated cluster of galaxies, modeled with a pre-heated intracluster medium, that exhibits X-ray features similar to the `cold fronts' seen in {\\sl Chandra\\/} observations.  Mock observations at a particular epoch show factor two discontinuities in X-ray temperature and factor four in surface brightness on a spatial scale $\\lta\\! 100$~kpc. Analysis of the cluster's dynamical history reveals that the front is a transient contact discontinuity created by an ongoing merger of two roughly equal mass subgroups.  The cold front feature in this realization is amplified by the adiabatic expansion of one of the subgroups following its ablation from the center of its local dark matter potential.  The presence of cold front features in a cluster modeled without radiative cooling or magnetic fields implies that such relatively complex physics is not a necessary element of the phenomenon and suggests that the prevalence of such features in high resolution \\xray images of clusters may simply reflect the high frequency of ongoing mergers driven by gravity and comparatively simple hydrodynamics. ",
        "introduction": "The improved spatial imaging of the hot intracluster medium (ICM) in galaxy clusters by the {\\sl Chandra\\/} Observatory has resulted in the discovery of apparent contact discontinuities, termed ``cold fronts'', in many clusters (Markevitch \\etal 2000;  Vikhlinin, Markevitch \\& Murray 2001; Mazzotta \\etal 2001;  Sun \\etal 2002). These cold fronts occur at the boundary of a local peak in the \\xray surface brightness and are typically located within a few hundred kpc of the core of the cluster.  The fronts exhibit a drop in temperature and corresponding rise in density upon entering the emission peak.  The term ``cold front'' is applied to contrast the phenomenon with a shock front where the temperature would be elevated in the direction of increasing density.  The origin of these cold fronts is thought to be related to cluster mergers (Markevitch \\etal 2000). From an extended Press--Schechter treatment, Fujita \\etal (2002) predict that up to one-third of present clusters will contain large \\xray subhalos. In this {\\sl letter}, we present evidence supporting the merger origin hypothesis from a simulated cluster evolved under a `preheated' assumption for ICM evolution (Kaiser 1991; Evrard \\& Henry 1991). In this treatment, the proto-ICM gas at high redshift is assumed to lie a fixed, elevated adiabat that results from heating due to star formation and/or AGN activity (Bower 1997;  Cavaliere, Menci \\& Tozzi 1998; Balogh, Babul \\& Patton 1999; Wu, Fabian \\& Nulsen 2000; Tozzi \\& Norman 2001; Voit \\& Bryan 2001).  The gas subsequently evolves under gravitationally-driven shock-heating, with magnetic fields and radiative cooling ignored.  As we prepared this work for publication, Nagai \\& Kratsov (2002), using completely independent techniques, present cold front phenomena in a pair of simulations of non-preheated cluster models.  Their simulations have higher spatial resolution compared to the one we present, but our Lagrangian simulations have an advantage in allowing the history of gas parcels to be tracked over time.  We use this ability to show that material once at the core of a merger progenitor is directly responsible for the cold emission seen at later stages of the merger.    In section~\\ref{sect:sim}, we present the simulation and compare its cold front properties to observations. The merger history and thermodynamics responsible for the cold front are examined in section~\\ref{ssect:phys}. Note that all scales quoted in this paper assume a Hubble constant H$_0 \\!=\\! 70 \\kmsmpc$. ",
        "conclusions": "We show that clusters modeled under the assumption of a preheated intracluster medium can exhibit features similar to the cold fronts observed in high resolution spectroscopic imaging of \\xray emission from clusters.   We present a particular realization, displaying a temperature dip coincident with a peak in \\xray surface brightness, where the cold front is a transient feature created by ablated core material of a merging subgroup.   Freed from its confining local dark matter potential, adiabatic expansion cools the core while its density remains sufficiently high to create a strong feature in emission.  This result supports the merger origin assumption for cold fronts and demonstrates that their observed characteristics can be reproduced by a gas dynamic treatment that ignores radiative cooling and magnetohydrodynamics.  A similar conclusion using independent methods is reached by Nagai \\& Kratsov (2002).  Although these studies provide an existence proof that cold fronts can result from mergers, we do not yet know if {\\sl all} observed cold fronts are consistent with this formation mechanism. It remains to be seen what fraction of mergers result in cold fronts and what combination of parameters --- mass ratio, impact parameter, angular momentum, viewing angle --- favors such an outcome. Although the cold front in this study occurs when the subcluster's gas strays from its local potential minimum and expands adiabatically, we do not know if this is a necessary condition.  Improved understanding of cold fronts will require more extensive searches within well-defined samples of observed and simulated clusters. The relatively small numbers of observed and simulated cold fronts must be increased to enable secure statistical studies of this phenomenon in the cluster population."
    },
    "0207/astro-ph0207656_arXiv.txt": {
        "abstract": "The time- and ensemble-averaged mechanical energy outputs of radio galaxies may be large enough to offset much of the cooling inferred from X-ray observations of galaxy clusters.  But does this heating actually counterbalance the cooling, diminishing cooling flows or quenching them altogether?  I will argue that energy injection by radio galaxies may be important even in clusters where no active source is present, due to the likely intermittency of the jets.  If the energy injected by radio galaxies percolates through the intracluster medium without excessive mixing, it could stabilize the atomic cooling responsible for X-ray emission. ",
        "introduction": "The impact of radio galaxies on their surroundings is probably far out of proportion to their numbers.  Although only a small minority of AGNs seems to produce powerful jets, virtually all of the mechnical energy output is transferred to the ambient medium.  Relatively little of this energy can be radiated away (Scheuer 1974), and therefore the majority must go into heat and motion.  In contrast, the large, easily detected supplies of radiant energy pouring out of most AGNs probably have little effect on the surroundings. Cool, dense clouds ($\\la 10^4$ K) in the interstellar medium of the host galaxy can easily reradiate this energy, while the hotter, diffuse gas responds thermally through the Compton effect, with a low efficiency $\\sim \\tau (h\\nu/m_e c^2)$, where $\\tau\\ll 1$ is the electron scattering optical depth.  Moreover, with Compton temperatures $\\la 10^7$ K for typical AGN spectra, irradiation is more likely to {\\it cool} the hot phase of the ISM/ICM through the inverse Compton effect, rather than heat it.  Nevertheless, in a time-averaged or ensemble-averaged sense, there seems to be plenty of energy associated with radio jets alone.  Peres et al.~(1998) compared the radio and X-ray luminosities of 58 clusters in a {\\it ROSAT} flux-limited survey.  The ratio of 1.4 GHz power to bolometric X-ray power from within the ``cooling radius\" of those sources identified as cooling flows was $\\sim 0.008$, just shy of 1\\%.  Now, the maximum possible synchrotron emissivity of radio lobes in internal pressure equilibrium is typically smaller than a few percent of the kinetic power, assuming equipartition and unit filling factor of relativistic electrons in the lobes.  This fraction evolves toward smaller values as the source expands (Begelman 1996, 1999), and is also decreased by any deviation from equipartition. (Note that it can be increased somewhat if the emission comes mainly from small regions at high pressure, e.g., strong shocks within the lobes, but this seems unlikely to produce a large correction.)  Thus it is very likely that the ensemble-averaged jet  power is at least of the same order as the observed X-ray power, and may be larger. But one must address at least two issues before leaping to the conclusion that energy injection by radio jets offsets X-ray cooling. First, most of the radio flux from the {\\it ROSAT} sample is contributed by a small number of clusters with very powerful radio sources.  Typically, only about 10\\% of the cooling flow clusters show strong {\\it current} activity. Second, it is notoriously difficult for simple heating mechanisms to balance X-ray cooling in a stable way.  I will attempt to address these issues below. ",
        "conclusions": "In a time- and ensemble-averaged sense, it appears that radio galaxies can supply enough energy to offset the observed cooling of intracluster gas.  But the relative rarity of active radio sources would require that they be intermittent, and that the heating effects due to mechanical energy injection persist long after the radio lobes have faded. There is good circumstantial evidence that the former is true, but whether the latter occurs is an open question.  I have speculated on a possible mechanism by which mechanical heating can offset cooling in a stable way. In order for effervescent heating to work, the buoyant fluid injected by the radio galaxy would have to spead its energy evenly through the cluster atmosphere, without mixing into the background at a microscopic level.  Intracluster ``weather\" and perhaps a certain amount of thermal conduction could help with this."
    },
    "0207/astro-ph0207130_arXiv.txt": {
        "abstract": "Over the last few years, numerical models of the behavior of solar magnetic flux tubes have gone from using methods that were essentially one-dimensional (i.e.\\ the thin flux tube approximation), over more or less idealized two-dimensional simulations, to becoming ever more realistic three-dimensional case studies. Along the way a lot of new knowledge has been picked up as to the e.g.\\ the likely topology of the flux tubes, and the instabilities that they are subjected to etc. Within the context of what one could call the ``flux tube solar dynamo paradigm,'' I will discuss recent results of efforts to study buoyant magnetic flux tubes ascending from deep below the photosphere, before they emerge in active regions and interact with the field in the overlying atmosphere (cf.\\ the contributions by Boris Gudiksen and {\\AA}ke Nordlund): i.e.\\ I am not addressing the flux tubes associated with magnetic bright points, which possibly are generated by a small-scale dynamo operating in the solar photosphere (cf.\\ the contribution by Bob Stein). The presented efforts are numerical MHD simulations of twisted flux ropes and loops, interacting with rotation and convection. Ultimately the magnetic surface signatures of these simulations, when compared to observations, constraints the dynamo processes that are responsible for the generation of the flux ropes in the first place. Along with these new results several questions pop up (both old and new ones), regarding the nature of flux tubes and consequently of the solar dynamo. ",
        "introduction": "Buoyant magnetic flux tubes are an essential part of the framework of the current theories of dynamo action in both the Sun and solar-like stars: it is believed that when formed near the bottom of the convection zone (CZ), by a combination of rotation and turbulent convection, toroidal flux tubes buoyantly ascend in the form of tubular $\\Omega$-shaped loops. Rising under the influence of rotational forces, they finally emerge after a few months as slightly asymmetric and tilted bipolar magnetic regions at the surface. Many models of buoyant magnetic flux tubes are based on the thin flux tube approximation \\citep{Spruit1981} that treats the tubes as strings, much thinner than e.g.\\ the local pressure scale height, moving subjected the Coriolis and drag forces. This essentially 1-d approximation is consistent with the observations when used to study the latitudes of emergence, tilt angles, and the tilt-scatter of bipolar magnetic regions on the Sun, but only if the initial field strength of the tubes are of the order of 10 times the convection equipartition value near the bottom of the CZ \\citep{DSilva+Choudhuri93,Fan+ea94,Caligari+ea95}. A major problem in dynamo theory is to understand how the field strength can become so high, corresponding to an energy density 100 times larger than that of the available kinetic energy.   However, as buoyant flux tubes rise they expand and the assumption that they are thin breaks down some 20 Mm below the solar surface. Traditionally, the latter fact is taken as the main reason for ``going into higher dimensions'', i.e.\\ for submitting to 2-d (actually 2.5-d) and 3-d models. So far, a lot of questions still remain unanswered, e.g.\\ whether the quasi-steady state topology that the flux ropes reach in the later phase of their rise in 2-d simulations (e.g.\\ Emonet \\& Moreno-Insertis 1998 and Dorch et al.\\ 1999) is stable towards perturbations from the surroundings, and whether the results found for 3-d flux ropes moving in a 1-d average static stratification, at all are valid in the more realistic case. Hence when attempting to review the status of buoyant magnetic flux tube models, there are several questions (Q's) that one may ask. Below are a few examples: ~\\\\ {{\\bf Q1:} In general the tubes may be twisted, but how much twist is needed and warranted?}\\\\ {{\\bf Q2:} How does the tube's twist evolve as they ascend: do they e.g.\\ kink due to an increasing degree of twist?}\\\\ {{\\bf Q3:} Are there other instabilities besides the magnetic Rayleigh-Taylor (R-T) and kink instabilities?}\\\\ {{\\bf Q4:} How do the tubes interact with the flows within the CZ and at the surface?}\\\\ {{\\bf Q5:} What happens to the less buoyant magnetic subsurface structure as the tubes rise (the wake)?}\\\\ {{\\bf Q6:} What happens when the ropes become thick, i.e.\\ large comparable to the local pressure scale height?}\\\\ {{\\bf Q7:} What happens at emergence? Does the field topology change (how)?}\\\\ {{\\bf Q8:} How does the twist arise? What is the appropriate initial condition?}  In the following I will try to answer some of the above Q's by reviewing some of the main results of 2-d and 3-d models of buoyant magnetic flux tubes with emphasis on the most recent results from the turn of this millennium. ",
        "conclusions": "In this review I have tried to shed light on some of the ``big'' questions in the theory of buoyant magnetic flux tubes. Below I summarize some of the answers (A's) to the questions posed in the Introduction.  ~\\\\ {{\\bf A1:} The critical pitch angle is about $30^{\\rm o}$ (for thick ropes), in the general 3-d case including (large-scale) convective flows, but without rotational effects. The result of \\cite{Abbett+ea00} that the inclusion of rotation lowers the critical pitch angle relative to the 2-d result, comes from non-convective simulations; i.e.\\ to finally answer this Q we need a simulation including both rotation and realistic convection.}\\\\ {{\\bf A2:} I found no evidence that the ropes kink as they rise, at least if their initial twist is low enough: one may speculate that an initially R-T unstable rope may stabilize due to the increase of the twist (which increases only slightly in the less symmetric 3-d models).}\\\\ {{\\bf A3:} There may be other instabilities associated with the motion through the uppermost strongly stratified super-adiabatic layers, but it is not know from the models presented here.}\\\\ {{\\bf A4:} Besides the shuffling of the flux ropes by the convective flows, the primary interaction is the cause of a flux-loss due to advective erosion.}\\\\ {{\\bf A5:} The weak magnetic field and the flux lost from the rising flux tubes are transported downwards by the flux pumping-effect.}  A note on A4: The numerical simulations show that the interaction of a buoyant twisted flux rope with stratified convection leads to a magnetic flux-loss from the core of the rope. During the simulation, the flux rope rises 96 Mm, and loses about 25\\% of its original flux content. This, ceteris paribus, leads to a small increase in the amount of toroidal flux that must be stored at the bottom of the CZ during the course of the solar cycle: Solar toroidal flux ropes rise about 200 Mm before emerging as bipolar active regions. One may thus expect them to lose even more of their initial flux, which would then be pumped back towards the bottom of the CZ. Moreover, the relative slip does not remain constant throughout the rope's rise \\citep{Dorch+ea01}.  Of course a lot of Q's still remain to be answered (three from my list): the most fundamental problems remaining are those of the origin of the twist, and the question of how it arises, Q8. This is not addressed by any of the models discussed here, but in my view one likely process is the generation of twisted field lines in large-scale flux bundles located near the bottom of the convection zone, connecting across the solar equator: such flux bundles would experience a rotating motion since their lower parts are located in a region rotating slower than their uppermost parts. This rotation would transmit a twist to the parts of the flux bundle a slightly higher latitudes, thereby possibly giving rise to a twisted toroidal flux system."
    },
    "0207/astro-ph0207306_arXiv.txt": {
        "abstract": "\\noindent We present the model for determination of pulsar distances or average electron distribution using a method similar to the widely used dependence of A$_V$ on distances in different directions. To have reliable pulsar distances, we have used several natural requirements and distances of pulsar-calibrators. We have constructed  dispersion measure-distance relations for pulsars in 48 different directions.   KEY WORDS: PULSAR DISTANCES, ELECTRON DISTRIBUTION, GALAXY ",
        "introduction": "In this paper, our goal is to give dependence between the dispersion measure (DM) and distances (d) for radio pulsars (PSRs) in each different direction, similar to the dependence between A$_V$ and distance. Naturally, degree of irregularity in electron distribution in the Galaxy is very small when compared with dust distribution. Most of the times, as a rule, a mathematical model for electron distribution is used to find distances of PSRs in the Galaxy. Using such models, PSR distances can be found with much error. Moreover, some of the PSR distances do not agree with the models. For them, individual distances are given. This is natural and this will not change even if perfectly complete models are constructed and used. We determined the distances to PSRs by using natural requirements (which will be explained) and distances which are well known by independent methods. This provided us with the knowledge to construct DM-distance relation for PSRs and the Galactic electron distribution very easily and with much less error. ",
        "conclusions": "To compare the distances of PSRs (new distances) that is obtained by our model and the distances obtained by the model of Taylor \\& Cordes (1993), in Figure 5 we have plotted out our new distances vs. Taylor \\& Cordes distances of PSRs. For most of these PSRs (from Table 1) almost the same independent distances are agreed on, but for some of them we have taken considerably different independent distances from that of Taylor \\& Cordes (1993) model gives. The names of these PSRs, as shown in Figure 5, are as follows; PSR J1302-6350, J1748-2021, J1804-0735 and J1910-59. As seen in Figure 5, the new distances and the distances given by Taylor \\& Cordes (1993) model differ up to more than twice. Most of these PSRs are located in south semi-sphere and were discovered during the 1400 MHz survey (with l between 300$^o$-360$^o$). Among the PSRs which are close to the Sun, the greatest difference in the distance given by our model and the model of Taylor \\& Cordes is owned by PSR J1939+2134 (see the distance value calculated from Taylor \\& Cordes model in the book by Lyne \\& Graham-Smith (1998).\\\\  To decide which model give the true distances, we need to test. In the previous section, in finding the PSR distance we have several criteria. One of these criteria is that the PSRs at the same age should be at almost the same distances from the Galactic plane. As we said before, the surveys done at 1400 MHz scanned the Galactic plane with $|b|<$5$\\degr$. Hence, they have discovered a lot of young PSRs. It is necessary to mention that the ratio of flux at 1400 MHz to flux at 400 MHz of young PSRs is several times more than this ratio for old PSRs (Guseinov et al. 2002a). This makes the discovery of young PSRs easier. In Figure 6, latitude (b)-longitude (l) distribution of PSRs whose l is between 300$^o$-360$^o$ and characteristic ages below 10$^6$ years, is given. As seen from this figure, dominant number of these PSRs have $|b|<$2$\\degr$. In Figures 7 and 8, the distance from the Galactic plane (z) vs. the distance from the Sun (d) for PSRs with age$<$10$^6$ years and in the longitude interval l=300$^o$-360$^o$ is given according to the new model and old model of electron distribution, respectively. The cone corresponds to 2 degrees since most of the young PSRs have $|b|<$2$\\degr$ as can be seen from Figure 6. It gives the limits of z at each distance d. As seen from these figures there is not an important difference neither in z nor in d. For only four of these PSRs Taylor \\& Cordes (1993) model has given very large distance values. This indicates that for small b values ($|b|<$2$\\degr$) there is no significant difference in electron distributions of the two models for l between 300$^o$-360$^o$. This is in general also true for other longitude directions. \\\\  For PSRs whose characteristic ages are between 10$^{6}-$10$^7$ years, b-l distribution with l between 300$\\degr$ and 360$\\degr$, is shown in Figure 9. As seen from the figure dominant number of these PSRs have $|b|<$5$\\degr$. In Figures 10 and 11, the distance from Galactic plane z vs. the distance from the Sun d is given according to the new model and old model electron distribution, respectively. PSRs with high b values are nearer PSRs. As seen from Figures 10 and 11, for PSRs farther than $\\sim$3 kpc, the new distances (d,z) are considerably smaller than the distances (d,z) given by the old model. It is seen from Figure 10 that for all distances number density of PSRs are higher near the Galactic plane, and for PSRs farther than 6 kpc the distances from the Galactic plane is in average the same. The distance distribution according to the old model does not agree with this criterion. Thus according to the Taylor \\& Cordes model, as the distances of PSRs increase, average distances from Galactic plane increases. It indicates that the space velocities of farther PSRs are larger, however there is no reason for this. \\\\  $Acknowledgments$ We thank T\\\"{U}B\\.{I}TAK, the Scientific and Technical Research Council of Turkey, for support through TBAG-\\c{C}G4. \\\\ \\\\ \\clearpage \\newpage"
    },
    "0207/astro-ph0207076_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "We present the results of a detailed kinematical and dynamical analysis performed in two highly interesting gas--rich Irregular galaxies (Irrs): IC 1613 and NGC 4449. The analysis has been accomplished by means of optical Fabry--Perot interferometry mapping the H$\\alpha$ and [SII] lines. Our interest was centered on global and local scales in both galaxies. The results focused on the global radial velocity field in both galaxies show that they are different from each other. IC 1613 displays a highly perturbed velocity field due to local motions caused by superbubbles. NGC 4449 shows a velocity field which is relatively well behaved with a weak decreasing range in radial velocities ($\\sim$60 km s$^{-1}$) from NE to SW. On local scales, we found that the HII region population for both galaxies displays a velocity dispersion mean value of $\\sim$20 km s$^{-1}$, which does not reflect the environment of the host galaxy in each case. The HII region population in these galaxies display standard values according to their diameter distribution and luminosity functions (D$_0$= 30--65 pc; $\\alpha$= 1.5--1.9). The superbubble population in IC 1613 displays expansion velocities (20--40 km s$^{-1}$), sizes (100--320 pc), dynamical ages (0.7--2 Myr) and mechanical energies (0.6--8$\\times10^{50}$ ergs) which can be explained --in the framework of the standard model-- in terms of the activity derived from the massive stars (stellar winds and supernova explosions) at those locations. Based on our results (Valdez--Guti\\'errez et al. 2001, 2002, Rosado et al. 2001), we conclude that the impressive distribution and \\hbox {kinematical/dynamical} behavior of the ionized gas in these two gas--rich Irrs has been moulded --on global and local scales-- by the activity of their massive stars; i.e. such activity has ``etched\" globally and locally the galaxies as they look today. Due to the nature of the present work --and as far as we know--, this is the first time that the kinematics and dynamics (on global and local scales) of the ionized gas in nearby gas--rich Irrs is unveiled in such detail."
    },
    "0207/astro-ph0207595_arXiv.txt": {
        "abstract": "Observed novae abundances and explosion energies estimated from observations indicate that there must be significant mixing of the heavier material of the white dwarf (C+O) into the lighter accreted material (H+He). Accordingly, nova models must incorporate a mechanism that will dredge up the heavier white dwarf material, and fluid motions from an early convection phase is one proposed mechanism.  We present results from two-dimensional simulations of classical nova precursor models that demonstrate the beginning of a convective phase during the `simmering' of a Nova precursor.   We use a new hydrostatic equilibrium hydrodynamics module recently developed for the adaptive-mesh code FLASH.  The two-dimensional models are based on the one-dimensional models of Ami Glasner\\cite{Glasner1997}, and were evolved with FLASH from a pre-convective state to the onset of convection. ",
        "introduction": "As a classical nova precursor accretes material from its neighbor, it heats up; by the time its peak temperature becomes roughly $4 \\times 10^7 {\\mathrm ~K}$ -- well before the final stages of runaway -- the accreted atmosphere becomes convectively unstable.  The resulting convective motions may be important for the process of dredging up white dwarf material into the accreted atmosphere.  In this paper, we examine the turn-on of convective motions in a white dwarf atmosphere based on one-dimensional early-time models provided to us by Ami Glasner.  This initial model is the same used in other multidimensional studies \\cite{Glasner1997,Kercek1998}, but taken at an earlier time -- at the last timestep before the onset of convection in the 1-d model code.  We map this model into the multidimensional FLASH code \\cite{flash} using techniques developed in \\cite{hse}, and perturb the models to investigate the onset of convective motions. ",
        "conclusions": ""
    },
    "0207/astro-ph0207240_arXiv.txt": {
        "abstract": "A variety of observations suggest that magnetic fields are present in all galaxies and galaxy clusters.  These fields are characterized by a modest strength $(10^{-7}-10^{-5}\\,{\\rm G})$ and huge spatial scale $(\\la 1\\,{\\rm Mpc})$.  It is generally assumed that magnetic fields in spiral galaxies arise from the combined action of differential rotation and helical turbulence, a process known as the $\\aod$-dynamo. However fundamental questions concerning the nature of the dynamo as well as the origin of the seed fields necessary to prime it remain unclear.  Moreover, the standard $\\aod$-dynamo does not explain the existence of magnetic fields in elliptical galaxies and clusters.  The author summarizes what is known observationally about magnetic fields in galaxies, clusters, superclusters, and beyond.  He then reviews the standard dynamo paradigm, the challenges that have been leveled against it, and several alternative scenarios.  He concludes with a discussion of astrophysical and early Universe candidates for seed fields. ",
        "introduction": "\\label{sec:introduc}  The origin of galactic and extragalactic magnetic fields is one of the most fascinating and challenging problems in modern astrophysics. Magnetic fields are detected in galaxies of all types and in galaxy clusters whenever the appropriate observations are made.  In addition there is mounting evidence that they exist in galaxies at cosmological redshifts.  It is generally assumed that the large-scale magnetic fields observed in disk galaxies are amplified and maintained by an $\\aod$-dynamo wherein new field is regenerated continuously by the combined action of differential rotation and helical turbulence.  By contrast, the magnetic fields in non-rotating or slowly rotating systems such as elliptical galaxies and clusters appear to have a characteristic coherence scale much smaller than the size of the system itself.  These fields may be generated by a local, turbulent dynamo where, in the absence of rapid rotation, the field does not organize on large scales.  In and of itself, the dynamo paradigm must be considered incomplete since it does not explain the origin of the initial fields that act as seeds for subsequent dynamo action.  Moreover, the timescale for field amplification in the standard $\\aod$-dynamo may be too long to explain the fields observed in very young galaxies.  It is doubtful that magnetic fields have ever played a primary role in shaping the large-scale properties of galaxies and clusters.  In present-day spirals, for example, the energy in the magnetic field is small as compared to the rotation energy in the disk.  To be sure, magnetic fields are an important component of the interstellar medium (ISM) having an energy density that is comparable to the energy density in cosmic rays and in the turbulent motions of the interstellar gas.  In addition, magnetic fields can remove angular momentum from protostellar clouds allowing star formation to proceed. Thus, magnetic fields can play a supporting role in the formation and evolution of galaxies and clusters but are probably not essential to our understanding of large-scale structure in the Universe.  The converse is not true: An understanding of structure formation is paramount to the problem of galactic and extragalactic magnetic fields.  Magnetic fields can be created in active galactic nuclei (AGN), in the first generation of stars, in the shocks that arise during the collapse of protogalaxies, and in the early Universe.  In each case, one must understand how the fields evolve during the epoch of structure formation to see if they are suitable as seeds for dynamo action.  For example, magnetic fields will be amplified during structure formation by the stretching and compression of field lines that occur during the gravitational collapse of protogalactic gas clouds.  In spiral galaxies, for example, these processes occur prior to disk formation and can amplify a primordial seed field by several orders of magnitude.  In principle, one should be able to follow the evolution of magnetic fields from their creation as seed fields through to the dynamo phase characteristic of mature galaxies.  Until recently, theories of structure formation did not possess the sophistication necessary for such a program.  Rather, it had been common practice to treat dynamo action and the creation of seed fields as distinct aspects of a single problem.  Recent advances in observational and theoretical cosmology have greatly improved our understanding of structure formation. Ultra-deep observations, for example, have provided snapshots of disk galaxies in an embryonic state while numerical simulations have enabled researchers to follow an individual galaxy from linear perturbation to a fully-evolved disk-halo system.  With these advances, a more complete understanding of astrophysical magnetic fields may soon be possible.  This review brings together observational and theoretical results from the study of galactic and extragalactic magnetic fields, the pieces of a puzzle, if you like, which, once fully assembled, will provide a coherent picture of cosmic magnetic fields.  An outline of the review is as follows: In Section II we summarize useful results from magnetohydrodynamics and cosmology.  Observations of galactic and extragalactic magnetic fields are described in Section III.  We begin with a review of four common methods used to detect magnetic fields; syncrotron emission, Faraday rotation, Zeeman splitting, and optical polarization of starlight (Section III.A).  The magnetic fields in spiral galaxies, ellipticals, and galaxy clusters are reviewed in Sections III.B-III.D while observations of magnetic fields in objects at cosmological redshifts are described in Section III.E.  The latter are essential to our understanding of the origin of galactic fields since they constrain the time available for dynamo action.  Section III concludes with a discussion of observational limits on the properties of cosmological magnetic fields.  Magnetic dynamos are discussed in Section IV.  We first review the primordial field hypothesis wherein large scale magnetic fields, created in an epoch prior to galaxy formation, are swept up by the material that forms the galaxy and amplified by differential rotation. The model has serious flaws but is nevertheless instructive for the discussion that follows.  Mean-field dynamo theory is reviewed in Section IV.B.  The equations for a disk dynamo are presented in Section IV.C and a simple estimate for the amplification rate in galaxies is given in Section IV.D.  The standard mean-field treatment fails to take into account backreaction of small-scale magnetic fields on the turbulent motions of the fluid.  Backreaction is a potentially fatal problem for the dynamo hypothesis for if magnetic fields inhibit turbulence, the dynamo will shut down.  These issues are discussed in Section IV.E.  Galactic magnetic fields, like galaxies themselves, display a remarkable variety of structure and thus an understanding of galactic dynamos has required full three-dimensional simulations.  Techniques for performing numerical simulations are reviewed in Section IV.F and their application to the problem of diversity in galactic magnetic fields is discussed in Section IV.G.  In Section IV.H we turn to alternatives to the $\\aod$-dynamo.  These models were constructed to address various difficulties with the standard scenario.  Section IV ends with a brief discussion of the generation of magnetic fields in elliptical galaxies and galaxy clusters.  The question of seed fields has prompted a diverse and imaginative array of proposals.  The requirements for seed fields are derived in Section V.A.  Section V.B describes astrophysical candidates for seed fields while more speculative mechanisms that operate in the exotic environment of the early Universe are discussed in Section V.C.  The literature on galactic and extragalactic magnetic fields is extensive.  Reviews include the excellent text by Ruzmaikin, Sokoloff, \\& Shukurov (1988a) as well as articles by Rees (1987), Kronberg (1994), and Zweibel \\& Heiles (1997).  The reader interested in magnetohydrodynamics and dynamo theory is referred to the classic texts by Moffatt (1978), Parker (1979), and Krause \\& R\\\"{a}dler (1980) as well as ``The Almighty Chance'' by Zel'dovich, Ruzmaikin, \\& Sokoloff (1990).  A survey of observational results from the Galaxy to cosmological scales can be found in Vall\\'{e}e (1997).  The structure of galactic magnetic fields and galactic dynamo models are discussed in Sofue, Fujimoto \\& Wielebinski (1986), Krause \\& Wielebinski (1991), Beck et al.\\,(1996), and Beck (2000) as well as the review articles and texts cited above. ",
        "conclusions": "It was the late 1940's when the Galactic magnetic field was independently proposed by theorists and detected by observers.  Since then, galactic and extragalactic magnetic fields have been the subject of intense and fruitful research.  Nevertheless, fundamental questions concerning their origin, evolution, and nature remain unanswered.  Magnetic fields have been detected in over one hundred spiral galaxies, in numerous elliptical and irregular galaxies, in galaxy clusters, and in the Coma supercluster complex.  New instruments such as the planned square kilometer array radio telescope will no doubt reveal new magnetic structures.  It is of interest to note that at present, there is no example of a meaningful null detection of magnetic field in a collapsing or virialized system.  Conversely, only upper limits exist on the strength of truly cosmological magnetic fields.  The fact that these limits are several orders of magnitude lower than the strength of galactic and cluster fields suggests that magnetic fields are amplified, if not created, during structure formation and evolution.  The magnetic fields found in spiral galaxies are unusual in that the strength of the large-scale component is comparable to that of the tangled component.  By contrast, the fields in ellipticals are random on $\\la 100\\,{\\rm pc}$ scales.  Likewise, cluster fields are tangled on the scale of the cluster itself.  The distinction no doubt reflects a key difference between spiral galaxies on the one hand and ellipticals and clusters on the other.  Namely, the stellar and gaseous disks of spiral galaxies are dynamically ``cold'', rotationally supported systems while ellipticals and clusters are dynamically ``hot''.  Evidently, the scale of the largest component of the magnetic field in any system is comparable to the scale of the largest coherent bulk flows in that system.  The $\\aod$-dynamo is the most widely accepted paradigm for the amplification and maintenance of magnetic fields in spiral galaxies. The hypothesis that magnetic fields are continuously regenerated by the combined action of differential rotation and helical turbulence is compelling especially in light of the observation that the magnetic structures in disk galaxies are in general spiral.  One may think of these structures as the MHD analogue of material spiral arms.  Spiral structure is believed to be a wavelike phenomenon where the crests of the waves are characterized by enhanced star-formation activity which, in turn, is triggered by an increase in the local density.  Likewise, magnetic spiral arms may reflect low-order eigenmodes in a disklike magnetized system.  A more direct connection between material and magnetic spiral structure is evident in certain galaxies where strong magnetic fields appear in the regions between the material arms.  The FIR-radio continuum correlation provides further evidence in support of a connection between star formation and large-scale galactic magnetic fields.  The magnetic fields in ellipticals and clusters require a different explanation.  Mergers may play the central role in the establishment of these fields since they are likely to be present in merger remnants, typically tidally shredded spiral galaxies.  The magnetic debris from merger events can act as seeds for subsequent dynamo action.  In addition, the energy released during a merger event can drive turbulence in the interstellar or intercluster medium.  It is unlikely that either ellipticals or clusters will support an $\\aod$-dynamo since differential rotation in these systems is relatively weak.  However, they may support fluctuation dynamos in which turbulence amplifies magnetic fields on scales up to the size of the largest eddies in the systems.  From a theoretical perspective, the greatest challenge in the study of galactic magnetic fields comes from the tremendous dynamic range involved.  The scale radius and height of a typical galactic disk are of order $10\\,{\\rm kpc}$ and $1\\,{\\rm kpc}$, respectively, while turbulent eddies in the ISM extend in size from subparsec to $100\\,{\\rm pc}$ scales.  Naive arguments suggest that even if initially, magnetic energy is concentrated at large scales, in a turbulent medium, there is a rapid cascade of energy to small scales. A mean-field approximation, where velocity and magnetic fields are decomposed into large-scale and small-scale components, bypasses this problem.  The equation for the large-scale magnetic field, known as the dynamo equation, incorporates the effects of the small-scale fields through the $\\alpha$ and $\\beta$ tensors which, in turn, attempt to capture the gross properties of the turbulence (e.g., helicity, spatial anisotropy).  The dynamo equation for an axisymmetric thin disk can be solved by means of a quasi-separation of variables which leads to eigenvalue equations in $t$, $\\phi$, $R$, and $z$.  The $t$-eigenvalue gives the growth rate of the magnetic field while the $\\phi$-eigenvalue characterizes the symmetry of the field under rotations about the spin-axis of the disk.  The $R$ equation is similar, in form, to the Schr\\\"{o}dinger equation and its eigenvalue feeds back into into the value of the growth rate.  The separation-of-variables analysis has yielded a number of encouraging results.  Chief among these is a demonstration of principle, namely, that disklike systems with a rotation curve similar to that of a spiral galaxy, can support a magnetic dynamo.  Numerical simulations provide the means to study more realistic models.  In particular, the effects of a finite disk thickness and deviations from axisymmetry can be explored.  These investigations suggest ways in which bisymmetric and/or odd party magnetic fields can be excited.  Unless one is willing to accept magnetic fields as a property of the Big Bang, their existence today implies a violation of the MHD approximation at some stage during the history of the Universe.  While MHD processes can stretch, twist, and amplify magnetic field, by definition, they cannot generate new field where none already exist. Proposals for the origin of the first magnetic fields are as varied as they are imaginative.  For example, interest in the exotic environment of the very early Universe, and in particular cosmological phase transitions, has spawned numerous ideas for the creation of seed magnetic fields.  A perhaps more appealing set of proposals relies on the ordinary astrophysical phenomena that occur during structure formation.  Magnetic fields will develop in AGN, stars, and the shocks that arise during gravitational collapse.  Indeed, rough estimates suggest that AGN and/or an early generation of stars will yield fields of strength $10^{-11}\\,{\\rm G}$ on galactic scales.  Dynamo action can amplify a field of this strength to microgauss levels by a redshift $z\\simeq 2$, a result consistent with observations of magnetic fields in high-redshift radio galaxies.  The astrophysical mechanisms mentioned above were proposed at a time when our understanding of structure formation was relatively crude. It is in part for this reason that the creation of seed fields and the dynamo have been treated as separate and distinct processes. Indeed, most studies of disk dynamos do not make specific references to particular models for seed field production.  Likewise, few papers on seed fields follow the resultant fields into the dynamo regime.  Today, semi-analytic models and numerical simulations enable us to study galaxy formation in detail, taking into account hierarchical clustering, tidal torques from nearby objects, gasdynamics, and feedback from star formation.  Moreover, observations of high-redshift supernovae, the CMB angular anisotropy spectrum, and large scale structure have pinned down key cosmological parameters such as the densities of baryons and dark matter and the Hubble constant.  In light of these developments, it may be possible to achieve a more complete description of the origin of galactic magnetic fields, one that begins with the production of seed fields and follows smoothly into the dynamo regime."
    },
    "0207/astro-ph0207254_arXiv.txt": {
        "abstract": "{\\small Mean orbital light curves of SS433 in different precessional phases are analysed for active and passive states separately. In passive states the mean brightness depends strongly on the disk orientation, the star is fainter by a factor $\\approx 2.2$ in the disk edge-on positions. In active states the brightness does not depend significantly on the precessional phase. We suggest that in active states hot gas cocoons surrounding the inner jets grow and can not be shielded by the disk rim in the edge-on phases. Brightest optical flares are clear separated in two groups in orbital phases, it is considered as indication of orbital eccentricity. Bright flares prefer specific precession and nodding phases, it favours the slaved disk model and the flares as disk perturbations by a torque applied to outer parts of the accretion disk.} ",
        "introduction": "Active states of SS\\,433 were isolated using the GBI radio monitoring program data (http://www.gb.nrao.edu/fgdoss/gbi/gbint.html) and direct inspection \\vspace*{-0.5cm} \\begin{figure}[h] \\centering \\epsfig{file=fabr_f1.ps,width=11.5cm} \\vspace*{-0.9cm} \\caption{ Radio and optical data in 1985 when SS\\,433 was mainly in quiet state and in 1980 when it was mainly active. The radio flux (GBI radio monitoring data) in Jy, optical flux in V magnitudes} \\label{fig1} \\end{figure} of the optical data. Fig.\\,1 shows radio and optical data in two observational seasons (1980 and 1985). Active states are clearly seen in radio. In visible region flares destroy the regular orbital and precessional variabilities. Orbital variability is seen \\begin{figure}[htb] \\vspace*{-0.4cm} \\centering \\epsfig{file=fabr_f2.ps,width=11cm} \\caption{ Mean orbital light curves for passive (circles, down curves) and active (crosses, upper curves) states in different phases of precession. Obvious flares were excluded. Relative intensity I\\,$=0$ corresponds to V\\,=\\,14\\magdot0. The accretion disk eclipses are phased at $\\phi_{orb}=0$} \\label{fig2} \\end{figure} as deep primary eclipses of the accretion disk by the donor--star (Min\\,I, $\\phi_{orb}=0$, $P_{orb}=13\\daydot08$). Precessional variability is a brightening when the accretion disk is the most open to observer ($T_3$ moment, $\\psi_{pr}=0$, $P_{pr}=162\\daydot4$) and weakening when the disk is in edge--on positions ($T_{1,2}$ moments, $\\psi_{pr}=0.34, 0.66$).  We analyse the orbital light curves of SS433 in different precessional phases for active and passive states separately. Both original and all published data of optical V--band photometry for 1979--1996 were used. The data--base consists of 2200 individual observations collected in Sternberg Institute. We used 1491 observations in passive states of SS\\,433 and 584 observations in active states, where obvious flares were excluded. We find that the light curve in active state is about the same as that of in passive with primary and secondary minima (Fig.\\,2, Fig.\\,3). However, it is very important that in active states the brightness does not depend significantly on precessional phase and the primary minima are not so deep as they are in passive states. We suggest a geometry of the inner disk parts as two hot gas cocoons surrounding the two inner jets. In active states the cocoons grow and they can not be shielded by the disk rim when the disk is edge--on. In Fig.\\,4 we show a sketch of the disk and cocoons in active and passive periods.  \\begin{figure}[htb] \\vspace*{-0.5cm} \\centering \\epsfig{file=fabr_f3.ps,height=8cm,width=12.5cm} \\caption{ Mean precessional light curves for passive (circles, down curves) and active (crosses, upper curves) states in a middle of the primary minimum (left) and in elongations (right) in phase intervals $\\Delta \\phi = \\pm 0.05$} \\label{fig3} \\end{figure}  When the disk is the most open to observer the mean brightness in elongations is the same in active and passive states. Probably the cocoon surrounding the approaching jet is not shaded up to its base by the disk rim in these precession phases and luminosity of the cocoon does not depend notably on its size. This may be in a case if the cocoon scatters ($\\tau_T \\sim 1$) inner radiation coming from the accretion disk funnel. The cocoons can be identified with a source of the UV radiation of SS\\,433, where Dolan et al. \\cite{Dea97} have detected the strong linear polarization directed along jets. The cocoons can be also identified with a source of the double--peaked He\\,II\\,$\\lambda 4686$ line observed in the disk \\cite{F97}.  \\begin{figure}[htb] \\centering \\epsfig{file=fabr_f4.ps,width=11.5cm} \\caption{ A sketch of the accretion disk with cocoons surrounding inner jet bases in active and passive states in two extreme precessional orientations} \\label{fig4} \\end{figure}  In passive states an amplitude of precessional modulation ($\\Delta I \\approx 0.4$) is about the same as amplitute of primary minima ($\\Delta I \\approx 0.5$). This means that the projected sizes of the outer disk rim and the companion star are the same. ",
        "conclusions": ""
    },
    "0207/astro-ph0207548_arXiv.txt": {
        "abstract": "{We present results from a deep imaging search for companions around the young bona-fide and candidate brown dwarfs Cha H$\\alpha$ 1 to 12 in the Cha I dark cloud, performed with HST WFPC2 (R, I, H$\\alpha$), VLT FORS1 (VRI), and NTT SofI (JHK$_{\\rm s}$). We find 16 faint companion candidates around five primaries with separations between 1.5$^{\\prime \\prime}$ and 7$^{\\prime \\prime}$ and magnitudes in R \\& I from 19 to 25 mag, i.e. up to 8 mag fainter than the primaries. While most of these companion candidates are probably unrelated background objects, there is one promising candidate, namely $1.5^{\\prime \\prime}$ SW off the M6-dwarf Cha H$\\alpha$~5. This candidate is 3.8 to 4.7 mag fainter than the primary and its colors are consistent with an early- to mid-L spectral type. Assuming the same distance (140 pc) and absorption (A$_{\\rm I}$ = 0.47 mag) as towards the primary, the companion candidate has $\\log$~(L$_{\\rm bol}/$L$_{\\odot}) = -3.0 \\pm 0.3$. At the age of the primary (1 to 5 Myrs), the faint object would have a mass of 3 to 15 Jupiter masses according to Burrows et al. (1997) and Chabrier \\& Baraffe (2000) models. The probability for this companion candidate to be an unrelated fore- or background object is $\\le 0.7\\%$, its colors are marginally consistent with a strongly reddened background K giant. One other companion candidate has infrared colors consistent with an early T-dwarf. In addition, we present indications for Cha H$\\alpha$ 2 being a close ($\\sim 0.2^{\\prime \\prime}$) binary with both components very close to the sub-stellar limit. Our detection limits are such that we should have detected all companions above $\\sim 1$~M$_{\\rm jup}$ with separations $\\ge 2^{\\prime \\prime}$ ($\\ge 320$~AU) and all above $\\sim 5$~M$_{\\rm jup}$ at $\\ge 0.35^{\\prime \\prime}$ ($\\ge 50$~AU). ",
        "introduction": "Extrasolar planets were detected by indirect methods, but no direct imaging was presented, yet. Imaging of sub-stellar companions around young very low-mass stars or brown dwarfs (rather than around normal solar-type stars) should be less difficult, because young sub-stellar objects are self-luminous due to ongoing contraction and very low-mass stars and brown dwarfs as primaries are intrinsically faint. Hence, the problem of dynamic range (large magnitude difference at small angular separation) is not that severe. Only a few brown dwarfs were detected directly as companions to normal stars so far, mostly around M-dwarfs, the first of which was Gl 229 B (Nakajima et al. 1995, Oppenheimer et al. 1995), and the youngest of which is TWA-5~B (Lowrance et al. 1999, Neuh\\\"auser et al. 2000).  Brandner et al. (2000) undertook an HST and AO survey for sub-stellar companions around T Tauri stars, which we now extend to the M6- to M8-type objects Cha H$\\alpha$ 1 to 12 in the Cha I dark cloud. We obtained H$\\alpha$ as well as R- and I-band images with the Hubble Space Telescope (HST) in order to search for wide visual companions around them. These observations are both deeper and have higher spatial resolution than any previous images of these objects. We complement the HST data with archived VRI data from the ESO 8.2m Very Large Telescope (VLT) and new JHK$_{\\rm s}$ data from the ESO 3.5m New Technology Telescope (NTT).  With deep imaging of bona-fide and candidate brown dwarfs as presented here, one can, in principle, find binaries comprised of (i) two low-mass stars (as some of the candidate brown dwarfs in the sample could be very low-mass stars), (ii) one brown dwarf and one low-mass star, (iii) two brown dwarfs, (iv) one giant planet and one low-mass star, (v) or one giant planet and one brown dwarf.  So far, only little is known about multiplicity of brown dwarfs: The first binary brown dwarf found was PPl 15, a spectroscopic binary made up of two brown dwarfs in the Pleiades (Basri \\& Mart\\'\\i n 1999). Then, faint companion candidates were detected directly near a few brown dwarfs, namely DENISJ1228.2 (Mart\\'\\i n et al. 1999), 2MASSWJ1146 (Koerner et al. 1999), and 2MASSJ0850 (Reid et al. 2001).  Multiplicity parameters of brown dwarfs (like binary frequency, orbit characteristics, mass functions of primaries and secondaries, etc.) will shed light on their as yet uncertain formation mechanism: Do brown dwarf companions form like stellar companions or like planets in circumstellar disks ? Do they form by fragmentation or core growth ? Do isolated brown dwarfs form in isolation or are they ejected from multiple systems ? If brown dwarfs are ejected in the early accretion phase (Reipurth \\& Clarke 2000), then one should not find any brown dwarfs as companions at ages $\\ge 1$ Myrs. If low-mass companions are ejected during the pre-main sequence phase (c.f. Sterzik et al. 2001), then young brown dwarf companions should be more frequent than old brown dwarf companions, which we can test with our sample of young objects in Cha I.  It is as yet unknown, whether brown dwarfs can have planets and, if so, what their typical separations from the primary objects would be. With several high-resolution spectra obtained with VLT/UVES, Joergens \\& Guenther (2001) found evidence for radial velocity variations in some of the Cha H$\\alpha$ objects, which could be due to giant planets or surface features.  In Sect. 2, we will present the properties of the targets. In Sect. 3, details about the HST observations and data reduction are given. The complementary VLT and NTT observations are presented in Sect. 4. The background contamination is estimated in Sect. 5. Properties of the most promising companion candidate are discussed in Sect. 6. In Sect. 7, we present indications for one of the primaries being a very close $\\sim 0.2^{\\prime \\prime}$ binary pair. We conclude with a brief discussion in Sect. 8. ",
        "conclusions": "From the images of Cha H$\\alpha$ 5 and 8 and their companion candidates observed on the PC chip, we determined the dynamic range achieved on the PC chip. From the deepest exposures with Cha H$\\alpha$ 4, 10, \\& 11 on WF4 and Cha H$\\alpha$ 8 on WF2, we determined the dynamic range achieved on the WF chips. In Fig. 13, the dynamic range is plotted as flux ratio versus separation, with flux ratio being the ratio between either a companion or the $3 \\sigma$ background level and the particular primary (I-band).  From the mean (and faintest, respectively) I-band magnitude of the primaries (being 16.0 mag and 17.5, respectively) and the dynamic range limit (log flux ratio being 4, i.e. a magnitude difference of 10 mag outside of 2$^{\\prime \\prime}$ with the PC), we can then obtain the magnitude limit for detectable companions, namely 26 mag (27.5 mag), or one mag brighter at 1$^{\\prime \\prime}$ separation (and 2 mag brighter for the WF chips). For an assumed I$-$K$_{\\rm s}$ color index of $\\sim 4.5$ mag for an L- or T- dwarf, this would correspond to a limit of 21.5 mag (23.0 mag) in K$_{\\rm s}$. This limiting magnitude at a distance of 140 pc and an age of 2 Myrs would correspond to a limiting companion mass of $\\le 1$M$_{\\rm jup}$ according to table 1 in Burrows et al. (1997) with B.C.$_{\\rm K}$=2 mag (as for Gl 229 B, Leggett et al. 1999).  Hence, we should have been able to have detected all companions with mass above $\\sim 1$M$_{\\rm jup}$ outside of 2$^{\\prime \\prime}$ (320 AU) and all companions down to a few M$_{\\rm jup}$ at $\\sim 100$ AU. Outside of $0.35^{\\prime \\prime}$ (50 AU), we should have detected all companions with masses above $\\sim 5$M$_{\\rm jup}$ (K$_{\\rm s} \\simeq 18$ mag).  \\begin{figure} \\vbox{\\psfig{figure=H3273F13.ps,width=12cm,height=7.5cm,angle=270}} \\caption{Dynamic range for our HST WFPC2 images: We plot the log of the flux ratio between primary and detected companion candidates (crosses) with I-band magnitudes from Table 3 (magnitude difference is given on the right-hand side y-axis). Also shown is the I-band dynamic range curve obtained on the PC and WF chips (determined as flux ratio between primary and flux being $3 \\sigma$ above the background level). The two dotted lines show the approximate flux ratio for 5 and 1 M$_{\\rm jup}$ mass companions (at 2 Myrs and 140 pc) around a primary (calculated for the mean primary I-band magnitude of 16 mag) according to Burrows et al. (1997) models. We indicate the locations of the very close binary candidate Cha H$\\alpha$ 2 (as star symbol in the upper left) and the most promising resolved companion candidate Cha H$\\alpha$ 5/cc 1.} \\end{figure}  If isolated brown dwarfs are ejected early in the accretion phase (Reipurth \\& Clarke 2001) or some time during the pre-main sequence evolution (Sterzik et al. 2001), one should expect more high-mass-ratio binaries among very young stars than among old stars, which we can check with our sample. We observed 11 M6-M8-type primaries with HST and found 16 clearly resolved wide companion candidates around five of them plus one additional very close ($0.2^{\\prime \\prime}$) candidate binary. From the background population, we estimated the probability for each candidate to be a true companion, so that we can expect a few real companions among the candidates. From the optical and IR colors, only two of the candidates could be L- or T-type objects (Cha H$\\alpha$ 5/cc 1 and Cha H$\\alpha$ 4/cc 2). Up to two true sub-stellar companions around 11 primaries correspond to a percentage of $18 \\pm 13~\\%$. All these detected companion candidates (as well as un-detected, but detectable companions) would have separations in the range from $\\sim 100$~AU (resolution limit) to 1000~AU (somewhat arbitrary upper limit) and high mass-ratios of 10 to 100 (with the ratio being the mass of the primary devided by mass of the companion). If Cha H$\\alpha$ 2 is a close binary (Sect. 7), then the secondary could also be a sub-stellar companion, so that we would then have $\\le 3$ sub-stellar companions around 11 primaries, i.e. $27 \\pm 16~\\%$.  This should be compared to the frequency and orbit characteristics of other high-mass-ratio binaries with large separations, i.e. brown dwarfs in orbit around a star. Recently, Gizis et al. (2001) estimated the frequency of wide, visual, old L- and T-type brown dwarf companions to normal stars to be $18 \\pm 14~\\%$. This number is consistent with our estimate given above, so that we find no evidence for the ejection scenario. However, because of the large error bars due to small-number-statistics, we do not have sufficient evidence to make a strong statement in this regard. In addition, the surveys discussed by Gizis et al. (2001) and our survey have different dynamical ranges and detection limits.  We also have to refrain from comparing our results with previous surveys for companions around young low-mass stars in Chamaeleon (Reipurth \\& Zinnecker 1993, Brandner et al. 1996, Ghez et al. 1997, K\\\"ohler 2002), because our sample of primaries (M6- to M8-type bona-fide and candidate brown dwarfs) is different from the previously surveyed samples (G- to early M-type T~Tauri stars) and because the previous surveys were restricted to smaller dynamical ranges, up to a magnitude difference of 5 mag between primary and companion candidate, while all but one (Cha H$\\alpha$ 5) or two (Cha H$\\alpha$ 2) of our companion candidates have larger magnitude differences.  In our observations, we found one promising companion candidate (Cha H$\\alpha$ 5/cc 1), which could be a brown dwarf (or even giant planet) companion. The other candidates are all fainter and often detected only in I. If they were true companions, they would all be giant planets, given their magnitudes, assuming the same age and distance as towards the primaries. However, they are all at physical separations $\\ge 200$ AU, which is not expected for planets, but should not be excluded. Follow-up 2nd epoch imaging and/or spectroscopy will show, whether and which of our companion candidates are truely cool and bound."
    },
    "0207/astro-ph0207581_arXiv.txt": {
        "abstract": "We study the cluster mass function and its evolution in different models with Dark Energy arising from a self--interacting scalar field, with Ratra--Peebles and SUGRA potentials. Computations are based on a Press \\& Schechter approximation. The mass functions we obtain are compared with results holding for open models or models with Dark Energy due to a cosmological constant. Evolution results, in some Dark Energy models, closely approach open models. ",
        "introduction": "One of the main puzzles in modern cosmology is the nature of Dark Energy (DE), whose presence seems to be required by SNIa data (see, e.g., Perlmutter et al 1999, Riess et al 1998). A joint analysis of CMB and LSS observations (see, e.g., Percival et al. 2002, Efstathiou et al 2002) also favor a flat Universe with a matter density parameter $\\Omega_m \\simeq 0.3$, mostly due to CDM and with a minor contribution of baryons ($\\Omega_b h^2 \\simeq 0.02$; $h$ is the Hubble constant in units of 100 km/s/Mpc; in this paper we shall take $h = 0.7$ and, unless differently specified, $\\Omega_m = 0.3$, anywhere). The residual energy content of the world, in the present epoch, should not be observable in the number--of--particle representation.  One of the most appealing possibilities is that such dark component arises from a self--interacting scalar field. With in the wide set of interaction potential suggested, a particular relevance is kept by Ratra--Peebles (1988, RP hereafter; see also Wetterich 1995) and SUGRA (Brax \\& Martin 1999, 2000) expressions: $$ V(\\phi) = \\Lambda^{4+\\alpha}/\\phi^\\alpha ~,~~~~~~~~~ V(\\phi) = (\\Lambda^{4+\\alpha}/\\phi^\\alpha) \\exp (\\kappa \\phi^2/2)~. \\eqno (1) $$ Here $\\Lambda$ is an energy scale, currently set in the range $10^2$--$10^{10}\\, $GeV, relevant for fundamental interaction physics; potentials depend also on the exponent $\\alpha$; fixing $\\Lambda$ and $\\alpha$, the DE density parameter $\\Omega_{DE}$ is automatically fixed; in this work we preferred to use as free parameters $\\Lambda$ and $\\Omega_{DE}$;  in SUGRA potentials, $\\kappa = 8\\pi G$ ($G$: gravitational constant).  In this work we try to determine some effects on galaxy clusters and their evolution, caused by replacing a simple cosmological constant with DE due to a scalar field self--interacting  according to the potentials (1).  The technique used to study non--linear evolution is the Press \\& Schechter (PS, Press \\& Schechter 1974) approach. It is based on the spherical collapse model, that Gunn \\& Gott (1972), Gott \\& Rees (1975), Peebles (1980) debated within the frame of pure CDM models, and Lahav et al (1991), Eke et al (1996), Brian \\& Norman (1998) and others generalized to the case of $\\Lambda $CDM. In spite of its approximation, such model, inserted in PS formulation, has been found to approach simulation results (see, e.g., Lacey \\& Cole 1993, 1994). Recent improvements of the method (Sheth \\& Tormen 1999, 2002, Sheth, Mo \\& Tormen, 2001, see also Jenkins et al 2001), allowing a better approximation, involve some more parameters and their use seems unnecessary when aiming just to compare different cosmological models. ",
        "conclusions": "Observable effects of the nature of DE have been considered by various authors. For instance, Cooray \\& Huterer (1999) discussed the relation between DE nature and gravitational lensing. Previous work on the value of $\\delta_c$ was made by Steinhardt, Wang \\& Zlatev (1999), although explicit outputs were not given. More recently, Lokas (2002) considered the behavior of $\\Delta_c$ and the mass function in DE models with constant $w = -p/\\rho$. This approximated treatment has been pursued in a number of recent papers and eases computations. Let us however remind that the value of $w$, when dynamical DE is considered, varies significantly in the relevant period. In Fig.~9 we report its variations between $z=0$ and 10, when structures form. Notice that the rate of $w$ variation is highest in the most critical redshift interval, between $z=0$ and 1--2. In RP models such variation is $\\sim 20\\, \\%$. In SUGRA models it is even greater, as $w$ passes from values $\\sim 0.8$, at $z=0$, up to values $\\sim 0.3$--0.4 in the above narrow $z$ interval. Taking constant $w$, therefore, is a dangerous approximation, whose reliability ought to be carefully inspected, in different problems.  \\begin{figure} \\begin{center} \\includegraphics*[width=9cm]{F9.eps} \\end{center} \\caption{Redshift dependence of $w$ for 4 RP ($\\Lambda/$GeV = $10^2$, $10^4$, $10^6$ and 10$^8$) and 2 SUGRA models ($\\Lambda/$GeV = $10^2$ and $10^8$). $\\Lambda$ values decrease from top to bottom curves. } \\end{figure}  The main results of this work are that: (i) the shape of the mass function of clusters, at $z=0$, is only mildly dependent on DE nature. On the contrary, (ii) the cluster evolution depends on the nature of DE in a significant way. More in detail, models with RP potentials closely approach the evolution expected in open CDM models. Only for $\\Lambda$ values as low as $\\sim 10^2\\, $GeV, the expected behavior in a RP model is appreciably different from an open CDM with the same $\\Omega_m$. On the contrary, the evolution of SUGRA models is intermediate between open CDM and $\\Lambda$CDM models.  Cluster data available within a few years were thought to be able to discriminate between open CDM and $\\Lambda$CDM, on the basis of the redshift dependence of cluster abundance. If independent data confirm that we live in a spatially flat world, finding an evolution closer to open CDM than to $\\Lambda$CDM will provide a precise information on the nature of DE.  As a by--product of the analysis leading to these conclusions, we also found the dependence of the virialization density contrast $\\Delta_c$ on $\\Omega_m$, on the nature of DE, and on redshift $z$. Such $\\Delta_c$ is to be used in various applications, e.g. to build SO algorithms able to find clusters in n--body simulations of models with DE."
    },
    "0207/astro-ph0207062_arXiv.txt": {
        "abstract": "{We report the discovery of a new gravitationally lensed QSO, at a redshift $z = 1.689$, with four QSO components in a cross-shaped arrangement around a bright galaxy. The maximum separation between images is $2\\farcs 6$, enabling a reliable decomposition of the system. Three of the QSO components have $g\\simeq 19.6$, while component A is about 0.6~mag brighter. The four components have nearly identical colours, suggesting little if any dust extinction in the foreground galaxy. The lensing galaxy is prominent in the $i$ band, weaker in $r$ and not detected in $g$. Its spatial profile is that of an elliptical galaxy with a scale length of $\\sim 12$\\,kpc. Combining the measured colours and a mass model for the lens, we estimate a most likely redshift range of $0.3 < z < 0.4$. Predicted time delays between the components are $\\la$ 10 days. The QSO shows evidence for variability, with total $g$ band magnitudes of 17.89 and 17.71 for two epochs separated by $\\sim 2$ months. However, the relative fluxes of the components did not change, indicating that the variations are intrinsic to the QSO rather than induced by microlensing. ",
        "introduction": "The majority of known gravitationally lensed QSOs displays image splitting into two components. Such systems offer relatively few constraints for their mass distributions -- usually just the two positions since the flux ratios might be changed by microlensing. Both for the purpose of studying lensing potentials (Keeton, Kochanek \\& Falco \\cite{keeton*:98:OPGL}) and for the purpose of measuring cosmological parameters (Schechter \\cite{schechter:01:GL}), quadruply imaged quasars are considerably more useful than their doubly imaged counterparts. The discovery of a new quadruply imaged is therefore welcomed by everyone except those who struggle to explain why their relative numbers are so high in radio lensing surveys (Rusin and Tegmark \\cite{rusi+tegm:01:FFI}).  In this paper we report the discovery of a new gravitationally lensed QSO with quadruple image splitting. The object was found to be multiple as part of a new, ongoing imaging survey for lensed quasars using the Magellan Consortium's 6.5~m Baade telescope on Cerro Las Campanas. We present the first spectrum of this QSO and analyse \\emph{gri} imaging data to establish a first suite of astrometric and photometric properties of the QSO components and the lensing galaxy. We then discuss the available constraints on the surface mass distribution in the lens and conclude with some prospects for future observations. ",
        "conclusions": "The new quadruple QSO \\he{} is an almost textbook example for gravitational lensing, with its four nearly identical components arranged symmetrically around a luminous early-type galaxy. Unlike most other known quadruple systems, photometric monitoring of this object should be relatively easy even in mediocre seeing conditions, because of its wide image separations. Furthermore, its location in the sky makes it accessible to both Northern and Southern observatories.  Owing to its symmetry, the time delay is expected to be short, and accurate measurement of differential time delays might therefore be difficult unless the QSO should prove to be variable on very short timescales. This could limit the usefulness of the object for cosmological purposes, but at the same time it makes it an attractive target for microlensing studies, because of the relative ease to separate intrinsic and microlensing-induced variations. Notice that compared to the Einstein Cross Q~2237$+$0305, the higher lens redshift in \\he{} implies a $\\sim 10\\times$ lower projected transverse velocity and hence a much longer characteristic timescale for high-amplification event from microlensing It will therefore be easier to obtain a well-sampled lightcurve, but unfortunately events will be rarer and take much longer to get covered.  We have presented evidence that the QSO experiences substantial flux variations on time scales of months and years. Whether microlensing could have a contribution in these variations is not yet clear, but it can already be said with certainty that monitoring of \\he{} will be a promising task."
    },
    "0207/astro-ph0207638_arXiv.txt": {
        "abstract": "We present numerical simulations of the electron-positron plasma creation process in a simple neutron star magnetosphere. We have developed a set of cascade `kernels', which represent the endpoint of the pair cascades resulting from monoenergetic photon seeds.  We explore two popular models by convolving these kernels with the seed photon distributions produced by curvature radiation and by inverse Compton scattering. We find that the pair plasma in either case is well-described in its rest frame by a relativistic Maxwellian distribution with temperature near $mc^2/k_B$.  We present cascade multiplicities and efficiencies for a range of seed particle energies and stellar magnetic fields.  We find that the efficiencies and multiplicities of pair creation are often lower than has been assumed in previous work. ",
        "introduction": "An electron-positron pair plasma is a key ingredient in most models of pulsar radio emission.  This plasma is assumed to come from a pair cascade which occurs in the open field line region close to the star's magnetic axis.  In this region,  rotation-induced electric fields pull charged particles from the polar cap and accelerate them to relativistic energies ($\\gam \\lap 10^7$).  These particles  radiate `seed' \\gray photons, either by curvature emission (Sturrock 1971; Ruderman \\& Sutherland 1975; Arons 1983), or by inverse Compton scattering of ambient thermal photons (Bussard, Alexander, \\& M\\'esz\\'aros 1986; Daugherty \\& Harding 1986; Sturner \\& Dermer 1994).  In the strong magnetic fields of pulsars, these primary photons are susceptible to magnetic one-photon electron-positron pair creation (Tsai \\& Erber 1974).  The newly formed leptons in turn radiate `secondary' photons, most commonly through synchrotron radiation (Harding \\& Preece 1987). The secondary photons may be capable of further pair production.  As Sturrock (1971) first pointed out, this cycle of energetic photon emission and further pair creation continues, forming a `pair production avalanche', which ends only when all remaining photons are transparent to pair creation.  Whether this pair plasma forms or not, and its properties when it does, are crucial issues in models of the radio emission region. The plasma properties determine the possible wavemodes which the plasma can support.  Excitations of these wavemodes ultimately become the radio emission we observe (after coupling to electromagnetic modes and escaping the magnetosphere). Propagation of these modes through the pair plasma may leave its signature on the observed signal (through dispersive effects).  In addition, the plasma itself may be the source of the excitation of the waves which lead to radio emission (if there is free energy available in the plasma distribution at birth to drive instabilities).  The need for a quantitative understanding of the pair plasma distribution function (DF) has been apparent in the literature for some time.  Existing calculations of plasma wavemodes and propagation have either assumed analytically convenient DFs, without physical justification, or have tried to quantify a plausible heuristic DF introduced by Arons (1981). (Examples include  Buschauer \\& Benford 1975, Arons \\& Barnard 1986, Beskin, Gurevich, \\& Istomin 1988; Kazbegi, Machabeli, \\& Melikidze 1991; Weatherall 1994, Lyubarskii 1996; Gedalin, Melrose, \\& Gruman 1998; Lyutikov, Blandford, \\& Machabeli 1999).  It is therefore critical to determine the properties of the plasma created by a pair-production avalanche in the pulsar magnetosphere.  This is the primary focus of our paper.  In addition to leptons, the pair cascade can also produce high-energy photons. A small number of pulsars exhibit such high-frequency emission.  Some authors ({\\it e.g.,} Romani 1996;  Horotani \\& Shibata 1999) believe that this emission comes from a high-altitude `outer gap' active region, others ({\\it e.g.}, Rudak \\& Dyks 1999;  Zhang \\& Harding 2000) believe that this emission originates in the polar cap region.  We therefore include photon spectra as a secondary focus in the cascade models we present in this paper.  \\subsection{The Setting: the Polar Flux Tube}  To place our calculation into a larger context, we briefly summarize the standard model of the pulsar magnetosphere. (We follow, for instance, Ruderman \\& Sutherland 1975, Arons 1992, or Melrose 1992, 1995).  Soon after pulsars were discovered, Goldreich \\& Julian (1969) pointed out that a pulsar magnetosphere would not be empty.  A rotating magnetized neutron star in vacuum generates electric fields strong enough to overcome the star's binding energy for electrons (and possibly light ions), contradicting the vacuum assumption.  Most of the inner magnetosphere is now assumed to corotate with a `force-free' electric field $\\be_{ff} = - c^{-1} (\\bom \\times {\\mathbf r}) \\times \\bb$, filled with the corotational (Goldreich-Julian) charge density $\\rho_{GJ} = (4 \\pi)^{-1} \\nabla \\cdot \\be_{ff}$, where $\\bom$ is the angular velocity of the star.  The exception is on those field lines which extend beyond the `light cylinder' (where the corotational speed is equal to $c$), defining the `polar flux tube' (Arons 1983). Radio emission is thought to originate in the plasma within this polar flux tube.  The polar flux tube's active properties are due to the extension of its $\\bb$ field lines beyond the light cylinder. Charged particles which stream outward along these lines are unable to maintain a corotational force-free state along the entire field line (even within the light cylinder). (In the `closed' magnetosphere, where $\\bb$ lines do not cross the light cylinder, $\\bb$ lines become electric equipotentials.)  Although individual particles in the flux tube may escape to form a stellar wind, polar currents are assumed to cross field lines somewhere near the light cylinder and complete a circuit back to the star to preserve its overall neutrality.  The details of how these polar currents return to the star are unknown.  This is unfortunate, because the global current structure is crucial for determining the accelerating electric fields, and thus the photon seeds for the pair production cascade (if it occurs).  If a cascade occurs, the pair plasma may allow the premature `shorting out' of the accelerating potential (\\eg Arons \\& Scharlemann 1979; Shibata, Miyazaki, \\& Takahara 1998).  This in turn modifies the conditions assumed to seed the cascade in the first place.  \\subsection{Modeling The Pair Cascade}  The pair cascade takes place within this setting.  We have already noted the complexity of a fully self-consistent solution.  Lacking this, important factors for the cascade development are uncertain.  What is the energy of the primary beam charges?  By what means do these primary charges produce seed \\gray photons? In contrast, the microphysics underlying the cascade is well-known. We know precisely the differential QED cross sections for lepton and photon production: the local $\\bb$ field, photon energy and impact angle determine the outcome.  Daugherty and Harding (1982; `DH82' hereafter) modeled the cascade numerically.  Their simulations began with a single particle streaming out along the $\\bb$ field, whose curvature radiation began the cascade. They studied the \\gray photon distribution produced by the cascade, with passing reference to the properties of the underlying plasma.  Our focus is different.  Because we are especially concerned with the properties of the pair plasma, we wanted to extend previous work to determine those properties, and their dependence on magnetospheric parameters (magnetic field and primary beam energy). Thus, we set up our calculation to determine both the momentum distribution function (DF) and the density of the pair plasma relative to that of the primary beam.  Because we were concerned about the uncertainties implicit in global magnetosphere models, we took a new approach to the cascade. We wrote a code which follows, in time and space, the cascade induced by a monoenergetic population of photons. We include pair production by the photons and synchrotron radiation by the leptons.  Each run terminates when all particle and photon production ceases. At the end of each run, we store  the photon and lepton distributions (typically about $10^5$ particles per run, binned into 50-100 momentum bins, depending on the number of particles available). We treat each such run as a `kernel'.  The final cascade is then formed as the composite of many such runs, each weighted by the relative distribution of seed photons of that particular energy.   In this paper we form composite cascades assuming the seed photons come from curvature radiation, or from inverse Compton scattering of ambient X-ray photons.  In the remainder of this paper, we describe  our code (\\S 2; with some details in the Appendix);  give a qualitative overview of the pair cascade process (\\S 2), and   present the results of our parameter-space survey (\\S 3). Our primary results are presented in \\S 4, where we form composite cascades, based on photon seeds from curvature radiation and inverse Compton scattering.  We close in \\S 5 with a summary and some final comments. The reader interested only in final plasma or photon distributions might skip the details and jump to \\S \\S 4 and 5. ",
        "conclusions": "We have simulated the pair production cascade which may occur in pulsar magnetospheres.  We did this by means of a numerical code which tracks leptons and photons as they propagate upward above the magnetic polar cap of a rotating neutron star with a dipolar magnetic field. The code simulates magnetic pair production by the photons and quantized synchrotron emission by the leptons, beginning with a population of monoenergetic seed photons.  We ran simulations with a range of seed energies, angles, and magnetic field strengths to obtain detailed results over a substantial parameter range relevant to pair cascade models. We used these results as kernels to model pair cascades produced by an initial particle beam, either by curvature radiation or by magnetic resonant inverse Compton scattering.  \\subsection{Our Results}  Our main results can be summarized succinctly.  \\textbf{Cascade Onset and Development.}   We find that $B \\sim 10^{11}$ G, or larger, is required in order for the cascade to go if standard values are assumed for the primary beam energies.  This is a consequence of the opacity condition, (6), and is {\\it independent of the magnetic field geometry}. This suggests that slow rotation-powered pulsars, and low-field millisecond pulsars, either have no pair production or have seed photon energies much higher than the standard models predict.  When the cascade occurs, we find that it develops and ends spatially within $\\delta r < r_*$ of the creation of the photon seeds, thus temporally within several microseconds after each photon injection event.  In an actual pulsar magnetosphere, it is possible that the seed photons are injected continuously, but gap models with rapid temporal variability of the accelerating potential and sparking (Ruderman \\& Sutherland 1975) are perhaps more likely due to the nonlinearity and magnitude of gap accelerating potentials.  If sparking occurs, the cascade timescale seen here may be directly connected to $\\mu$sec flickering known as microstructure.  \\textbf{Efficiencies and Multiplicities.}  Our computed efficiencies and multiplicities stand in contrast to  assumptions made in the literature.  In particular, it is difficult to transfer a large fraction of the primary beam's energy to the pair plasma; some, often most, of it escapes as photons. In addition, our computed multiplicities (number of pairs produced per primary beam particle) are often small, only a few tens. Only rarely do they reach the large values, $> 10^3$, which are often assumed in the literature; curvature radiation with a large beam energy seems to accomplish this best.  \\textbf{The Pair Plasma.} We find that the plasma DFs in all cases can be well described as relativistic Maxwellians {\\it in the plasma's comoving frame}.  They generally have temperature $k T \\sim m c^2$.  The plasma flows out along field lines with Lorentz factors  $\\gamma_{\\rm CM} \\sim 100 - 1000$.  The resulting Lorentz boost into the pulsar frame gives the characteristic, asymmetric shape seen in  figure 8.  \\textbf{The Escaping Photons.}  We find that the escaping photon spectra tend to be steep in lower magnetic fields, due to the abundance of cooler, secondary synchrotron photons.  Conversely, they tend to be flatter in higher fields, where synchrotron photons are rare and the spectra are dominated by those seed photons which escaped pair creation.  In addition to these general trends, the details of the cascade development are sensitive to the initial field strength and also the seed photon spectrum.  \\textbf{Effect of Magnetic Field Strength.} Qualitative differences in the pair cascade occur for initial magnetic field strengths of $B_* = 10^{13}$ compared to $B_* = 10^{12}$ G.  The `transition' field strength is likely near $B_* = 0.1 \\bc \\sim 4.4 \\times 10^{12}$ G. The differences are due to the abundance, or dearth, of secondary synchrotron photons.  Below this transition field strength, leptons tend to be born into excited Landau levels and copious synchrotron emission takes place.  This occurs because the pair formation events tend to be governed by opacity rather than energetics.  Above this transition field, pair formation events are generally controlled by energetics, so that leptons are usually born in their lowest Landau state, and synchrotron emission is rare or absent.  \\textbf{Effect of Seed Photon Source.} The overall cascade efficiency is sensitive to the seed mechanism.  We find that CR-seeded models give pair multiplicities of up to a few times $10^3$, but typically do not transfer much of the beam energy to the pair plasma. In contrast, ICS-seed models give multiplicities of fewer than $100$, but can efficiently transfer the beam energy into the pair plasma (however, this requires ambient thermal photons at the magnetic resonance of the beam, which we assumed at the outset). Furthermore, the ICS mechanism operates at much lower beam energies than does curvature radiation.  \\textbf{Independence of Seed Photon Source.} Despite the enormous range of cascade efficiencies, we find that the plasma and photon distributions are fairly similar for both seed photon  spectra which we used. Perhaps the likenesses should not be too surprising, since the general ingredients required for a pair creation cascade are somewhat independent of seed mechanism.  First, photons must be of sufficient energy to pair produce; the lower-end shape of the photon spectrum does not affect the plasma DFs at all. We chose photon seed spectra based on typical beam energies in the literature; only a small fraction of the seed photons these charges produce are able to seed the cascade.  Second, a high-energy cutoff is always present in the photon spectrum.  This cutoff, together with the low-energy cutoff where the primary photons no longer produce pairs, limits the effective (pair-producing) seed photon spectrum to a relatively narrow range of energies.  Thus, many qualitative features of the pair cascade are not very sensitive to the mechanism of radiation which seeds the cascade (provided that it seeds a cascade at all).  \\subsection{Comparison to Other Simulations}  Other authors have performed numerical pair cascade simulations; however, the focus has usually been on the final escaping \\grys rather than on the pair plasma.  DH82 simulated cascades seeded by curvature radiation.   Our CR-seeded photon spectra for high beam energies and field strengths $B_* = 10^{12}$ G agree qualitatively with those of DH82, but we find much softer spectra in the higher $B_*$ case than they do.  The single lepton DF shown in DH82 (their figure 7) is for $B_* = 10^{12}$ G, and a beam energy of $10^{13}$ eV ($\\gam_b \\sim 2 \\times 10^7$).  It shows a rough power-law above $\\ppar \\sim 100 mc$, and few pairs at lower momenta.  In addition, they predict that {\\it more} low-$\\ppar$ pairs will be created in higher fields, since photons of lower energy will be opaque to pair creation. We believe that their conclusions (in contrast with what we present here) can ultimately be traced to their assumption that both members of the pair share the direction and half the energy of the parent photon.  (Note that DH82 predates Daugherty \\& Harding 1983, where they present energy-differential pair production cross-sections.  In stronger fields, the lack of synchrotron radiation is more significant than the ability of lower-energy photons to pair create, and so the DFs still have fewer leptons at low $\\ppar$.)  Sturner \\etal (1995) investigated cascades seeded by inverse Compton radiation. They do not present the final plasma DFs at all.  Our ICS-seeded photon spectra for $B_* = 10^{12}$ agree qualtitatively with their results for $B_* = 4 \\times 10^{12}$.  Note that we find much flatter spectra for our $B_* = 10^{13}$ simulations (for which Sturner \\etal did not offer spectra).   \\subsection{Comparison to other Plasma DFs}  The plasma DF is  critical in determining the wavemodes the plasma can support.  Knowledge of the DF is therefore crucial to calculations of radio emission and signal propagation in the pulsar magnetosphere.  Previous work has had to assume some some analytic expression, not necessarily physically motivated.  Some work has chosen analytically convenient forms, such as cold plasma (delta function), boxcar and bell-curve DFs, and relativistic Maxwellians (both in the comoving and pulsar frames).  Our work supports one of these assumptions, namely, a comoving Maxwellian. Our final DFs are only marginally relativistically hot, with temperatures $k_B T \\sim mc^2$ in most cases, and lower temperatures in $B_* = 10^{13}$ fields when the beam energy is just above the threshold to seed a pair cascade.  The large spread in $\\ppar$ in the pulsar-frame DFs is simply due to of Lorentz boosting.  In addition, some authors have used an analytic representation of a DF which Arons (1981) presented as a cartoon model. The Arons cartoon-DF has a peak at $\\ppar \\sim$ a few $mc$. It drops off exponentially above $\\ppar \\sim 10^3$--$10^4$. Both of these features are roughly consistent with our lower-field results, where synchrotron photons are produced.  Our higher-field results have a qualitatively similar shape but at higher $\\ppar$ values. In addition, the Arons cartoon-DF has a flat region between these extremes. We do not see this feature in our DFs, even in the $B_* = 10^{12}$ case.   \\subsection{Impact on Future Work}  We anticipate that the results presented here will be most relevant to plasma-based modeling of the magnetosphere and of the radio emission mechanisms. We give two examples.  The pair plasma plays a critical role in the electrodynamics of the polar flux tube;  it is often assumed to short out the parallel electric field and terminate the acceleration region.  Ruderman \\& Sutherland (1975) estimated that the pair multiplicity should be around $\\gam_b/\\gam_{\\rm pairs}$. This is typically two orders of magnitude larger than the multiplicities we find (except for the high-field magnetic ICS runs). Shibata \\etal (1998) agree with Ruderman \\& Sutherland that these large multiplicities ($10^3$--$10^5$) are necessary to short out the accelerating $\\be_\\parallel$ field.  Our results suggest that the situation is not so simple.   If complete shorting-out cannot be maintained, we might expect a complex, non-steady plasma flow in the region.   Instabilities in the pair plasma are often assumed to give rise to coherent radio emission.  One such mechanism is a two-stream instability, driven by relative motion of the two signs of charge ({\\it e.g.}, Buschauer \\& Benford 1976).  The  high temperatures assumed by some authors tend to suppress this instability ({\\it e.g.}, Weatherall 1994);  our results however find a lower temperature which may be favorable for this mechanism. We are presently investigating the details of this instability in the pair plasmas found in our calculations (Arendt \\& Weatherall, in progress)."
    },
    "0207/astro-ph0207312_arXiv.txt": {
        "abstract": "{ Based on a power spectrum analysis of the IRAS ISSA maps, we present the first detection of the Cosmic far-Infrared Background (CIB) fluctuations at 60 and 100 \\ump. The power spectrum of 12 low cirrus emission regions is characterized by a power excess at spatial frequencies higher than $k \\sim 0.02$ arcmin$^{-1}$. Most of this excess is due to noise and to nearby point sources with a flux stronger than 1 Jy. But we show that when these contributions are carefully removed, there is still a power excess that is the signature of the CIB fluctuations. The power spectrum of the CIB at 60 and 100 \\um is compatible with a Poissonian distribution, at spatial frequencies between 0.025 and 0.2 arcmin$^{-1}$. The fluctuation level is $\\sim 1.6\\times10^{3}$ Jy$^2$/sr and $\\sim 5.8\\times10^{3}$ Jy$^2$/sr at 60 and 100 \\um respectively. The levels of the fluctuations are used in a larger framework, with other observationnal data, to constrain the evolution of IR galaxies (Lagache et al. 2002). The detections reported here, coupled with the level of the fluctuations at 170 $\\mu$m, give strong constraints on the evolution of the IR luminosity function. The combined results at 60, 100 and 170 $\\mu$m for the CIB and its fluctuations allows, on the CIB at 60 $\\mu$m, to put a firm upper limit of 0.27 MJy/sr and to give an estimate of 0.18 MJy/sr. ",
        "introduction": "The Cosmic Far-Infrared Background (CIB) was detected by \\cite{puget96}, exploiting COBE-FIRAS data, and has now been firmly established over a range of wavelengths (e.g. \\cite{dwek98,gispert2000,hauser2001}). The intensity is quite high with respect to predictions based on evolutionary models of star formation in galaxy populations inferred from optical data. Source counts obtained with SCUBA at 850 $\\mu$m and ISO at 170 and 15 $\\mu$m have partly resolved the CIB and shown a strong cosmological evolution. In the future, far-IR and sub-millimeter telescopes from ground and space will perform deep surveys over small areas, aimed at resolving a substantial fraction of the CIB and to shed light on the number density, luminosity and spectral evolution of the infrared galaxy populations.  However, investigation of the clustering of these populations requires surveys over much larger areas. One way to tackle the limitation on the number of detected galaxies per field is to search for CIB fluctuations. So far, CIB fluctuations has only been observed at 170 $\\mu$m in the FIRBACK fields \\cite[]{lagache2000} and at 90 and 170 $\\mu$m in the Lockman hole \\cite[]{matsuhara2000} surveys. These detections were probably dominated by the Poissonian contribution. By analyzing larger FIRBACK fields, power spectra seem to reveal correlated fluctuations, well above the cirrus contribution \\cite[]{puget2002}. These results are currently under further investigation. Above $170\\,\\mu$m, only SCUBA observations might currently be able to give information on the fluctuations. However, no significant CIB correlations have been detected in the SCUBA maps \\cite[]{peacock2000}. The CIB anisotropies are mainly contributed by moderate to high redshift star-forming galaxies, whose clustering properties and evolutionary histories are currently unknown. Since the clustering strength depends on the bias at the relevant redshift, observing the CIB correlated fluctuations will provide valuable informations on bulge and elliptical formation, as well as potentially QSOs, thereby providing clues on the physical relations between dark matter and starburst galaxies. \\\\ In this paper we present the results of a power spectrum analysis of the IRAS 60 and 100 \\um emission of 12 regions in the sky with very low interstellar emission. All these regions are characterized by an excess of power at high spatial frequencies with respect to the interstellar emission. We will show that this excess is not of instrumental origin and may be attributed to the CIB. The discovery of the CIB fluctuations at 60 and 100 $\\mu$m could give constraints on the number counts below the IRAS point source detection limit. However, instead of extrapolating the number counts at 60 and 100 $\\mu$m using the level of the fluctuations, which is of limited cosmological interest, we prefer to use the two detections in a more general framework of the modelisation of the IR galaxy evolution that combines all existing number counts, redshift distributions, and observations of the CIB and its fluctuations, in the whole IR and submm range \\cite{lagache2002}. \\\\ The paper is organized as follow. After a presentation (\\S~\\ref{data}) of the data used for this analysis we describe in \\S~\\ref{powerspectrum} the power spectrum at 60 and 100 \\um of the 12 fields selected. In \\S~\\ref{cib} we present the method used to separate the Galactic and extra-Galactic contributions to the power spectrum and we discuss our results in \\S~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  \\subsection{CIB fluctuations at 60 and 100 \\um}  We have shown here that the power spectrum of high latitude fields at 60 and 100 \\um is characterized by a break at small scales (near 0.02 arcmin$^{-1}$). As indicated by \\cite{gautier92}, the power spectrum of the dust emission is usually well described by a power law proportional to $k^{-3 \\pm0.2}$. A more detailed study of the statistical properties of the interstellar cirrus HI 21 cm emission has been carried out by \\cite{miville-deschenes99b}. In this analysis it is shown that there are limited variations of the spectral index of the power law from field to field but, what is most important for the present work, the power spectrum of cirrus emission for scales smaller than $12.5^\\circ$ is always characterized by a single power law with no break. Therefore, it is unlikely that the power excess observed here at small scales is of interstellar origin.  We could also wonder if this break is of instrumental origin. It was shown by \\cite{wheelock93} that the response of the IRAS detectors are affected by memory effects. This produces variations of the detector response as a function of scale. This effect is more important at small scales (under a few tens of arcminutes) but \\cite{wheelock93} have shown that the amplitude of the fluctuations at these scales were underestimated. This effect will thus produce a drop of the power spectrum at small scale and cannot explain the power excess observed here.  On the other hand, instrumental noise could produce such an excess in the power spectrum. But, as the IRAS ISSA maps result from redundant individual observations, we were able to estimate the contribution of the noise to the power spectrum. We are aware that our estimate of the noise rely on the fact that the individual HCONs are independant. This is not perfectly true as, in the construction of the HCONs, an offset was added to each scan to minimize the difference between different observations of the same position. Therefore the noise level estimated by subtracting HCONs may be underestimated at the scale of a scan, which is a few degrees. But at this angular scale the signal is completely dominated by the cirrus emission, even in the low brightness regions selected for our analysis. At the scale of a few arcminutes where the CIB is detected, the noise contribution to the power spectrum has been removed accurately.   \\begin{figure*} \\hspace{-0.7cm} \\includegraphics[width=\\linewidth]{H3453F5.ps} \\caption{\\label{z-distrib} Redshift distribution of the sources making the CIB and the fluctuations at 60, 100 and 170 \\ump. The first two-left panels are for the total contributions, the last two-right panels, for sources with flux below 1 Jy \\cite[]{lagache2002}} \\end{figure*}  In fact it appears that most of the power excess can be attributed to the numerous extra-galactic point sources that are present in such a low cirrus emission field. When the strong ($I_{100 \\mu m} > 1$ Jy) point sources are removed from the ISSA maps, we recover a power spectrum typical of cirrus emission at low spatial frequencies but with still a power excess at small scales (k$> 0.02$ arcmin$^{-1}$) that can be attributed to the unresolved cosmic infrared background. Moreover, the residue has homogeneous properties over the sky, consistent with CIB.\\\\  \\cite{knox2001} computed the expected power spectrum of the CIB at several frequencies ($\\nu \\le$ 1060 GHz), exploiting the far-IR volume emissivity derived from the count models of \\cite{guiderdoni98} and assuming a bias $b=3$, constant with redshift.  They concluded that the clustering-induced fluctuations can match those of the CMB at $\\ell\\lesssim$300. They also predict a shape of the CIB power spectrum peaking around scales of 1-3 degrees. This broad maximum, if present, is at the limit of our frequency range where the noise is becoming large, making the detection of the clustering very difficult. The power spectra of the CIB at 60 and 100 \\um are compatible with a Poissonian distribution with levels $\\sim 1.6\\times10^{3}$ Jy$^2$/sr and $\\sim 5.8\\times10^{3}$ Jy$^2$/sr respectively.  \\subsection{CIB intensity and anisotropy amplitudes color ratio}  The CIB rms fluctuations in the IRAS maps corresponding to the white noise power spectra are: \\begin{equation} \\sigma^{2}= \\int P_{\\rm cib}(k) 2 \\pi k dk \\quad Jy^2/sr^2 \\end{equation} giving $\\sigma$=0.048 MJy/sr and $\\sigma$=0.09 MJy/sr at 60 and 100 \\um respectively. As shown by \\cite{gispert2000}, the CIB at different wavelengths is dominated by sources at different redshifts, larger wavelengths being dominated by more distant sources. The same applies to the fluctuations: lower frequencies probe higher redshifts (e.g. \\cite{knox2001}) The fluctuations at 60 \\um are dominated by nearby bright objects. When we remove these objects, the residual fluctuations are quite low. At 100 \\ump, the contribution of higher redshift objects is increasing, leading to higher level of residual fluctuations. Therefore, the ratio of the 60 to 100 \\um fluctuation is qualitatively consistent with what is expected.  This is illustrated more quantitatively in Fig.\\ref{z-distrib} on the panels showing the redshift distribution of sources contributing to the CIB intensity and the fluctuations \\cite[]{lagache2002}. It is clear from these figures that the z-distribution of the fluctuations is bimodal, with one contribution at redshift lower than 0.25 and the other one centered at redshift around 1. For all sources the ratio of nearby to moderate-redshift source contribution to the fluctuations is equal to 3.8, 1.8 and 1.3 at 60, 100, 170 \\um respectively, illustrating that fluctuations at larger wavelengths are dominated by more distant sources. When the brightest sources are removed the ratio of nearby to moderate-redshift contribution becomes equal to 1.7, 1 and 0.74 at 60, 100, 170 \\ump. In this case, at 100 \\um the contribution to the fluctuations of nearby and moderate-redshift sources is the same, becoming lower at higher wavelength. At 60 \\ump, the fluctuations are still dominated by the nearby objects.   For the three wavelengths, the CIB is mainly due to sources at redshift around 1.  A detailed analysis of the CIB fluctuations at 100 and 170 \\um (which is beyond the scope of this paper) will give information on the distribution of sources at z$\\sim$1 making the bulk of the CIB. This is particularly true at 170 \\um where sources with flux lower than 4$\\sigma$=135 mJy can be removed \\cite[]{dole2001}, leading fluctuations highly dominated by the moderate-z sources.   We can compute the ratio of CIB fluctuations to intensity ($R_{\\lambda} = \\frac{\\sigma_{\\lambda}}{I_{\\lambda}}$) at 100 \\um and compare it with the previous determination at 170 \\ump. To compute R$_{100}$ and R$_{170}$, we use: \\begin{itemize} \\item a CIB intensity at 100 $\\mu$m of 0.5 MJy/sr \\cite[]{renault2001}, leading to R$_{100}$ = 0.18 \\item a CIB intensity at 170 \\um of 1 MJy/sr \\cite{lagache2001}. For the fluctuations, we assume that a cut of 1 Jy at 60 and 100 \\um corresponds to the same cut at 170 \\ump. We obtain rms fluctuations of 0.12 MJy/sr (corresponding to $\\sim$11000 Jy$^2$/sr, \\cite[]{puget2002}). This gives R$_{170}$ = 0.12 \\end{itemize} The ratio R is decreasing between 100 and 170 \\ump, as expected. At 60 \\ump, the minimal hypothesis is to consider that R$_{60}$ is equal to R$_{100}$=0.18 which gives an upper limit of 0.27 MJy/sr on the CIB intensity. An illustrative logarithmic extrapolation of R out to 60 \\um gives R$_{60}$=0.27 leading to an estimate of the CIB intensity at 60 \\um of 0.18 MJy/sr, which is significantly smaller than the previous determination of \\cite{finkbeiner2000} of 0.56 MJy/sr.  \\subsection{Implication for the component separation}  This work suggests that the high latitude IRAS maps, in the lowest cirrus regions, cannot be used as a tracer of the interstellar extinction structure as proposed by \\cite[]{schlegel98}. In fact it should be noted that it is only above an intensity of order of 10 MJy/sr at 100 \\um that the CIB fluctuations are lower than the cirrus contribution at the smallest scales (k$\\sim$0.2 arcmin$^{-1}$).  Present and future CMB observations, above 100 GHz, with high sensitivity bolometers need to remove foreground contributions (cirrus and CIB fluctuations). The CIB spectrum being significantly ``colder'' than the cirrus spectrum (${I_{\\rm cirrus}100\\mu m} / {I_{\\rm cirrus}1mm} \\sim$ 30; ${I_{\\rm CIB}100\\mu m} / {I_{\\rm CIB} 1mm} \\sim$ 5), the relative contribution of the CIB will increase with wavelength. It is thus expected that at 1 mm the range in l space dominated by the CIB will be much more extended than sees at 100 \\ump. This question will be dealt in a forthcoming paper."
    },
    "0207/astro-ph0207124_arXiv.txt": {
        "abstract": "{We present the complete photometric database and the color-magnitude diagrams for 74 Galactic globular clusters observed with the HST/WFPC2 camera in the F439W and F555W bands. A detailed discussion of the various reduction steps is also presented, and of the procedures to transform instrumental magnitudes into both the HST F439W and F555W flight system and the standard Johnson \\( B \\) and \\( V \\) systems. We also describe the artificial star experiments which have been performed to derive the star count completeness in all the relevant branches of the color magnitude diagram. The entire photometric database and the completeness function will be made available on the Web immediately after the publication of the present paper. ",
        "introduction": "\\label{intro}  Because of its excellent resolving power, the \\textit{Hubble Space Telescope} offers an exceptional opportunity to study the crowded centers of Galactic globular clusters (GGC). We therefore began several years ago a study of color--magnitude diagrams (CMD) of clusters, using HST's WFPC2 camera.  Our first program, GO-6095, examined 10 clusters in the HST \\( B \\) (F439W) and \\( V \\) (F555W) bands (Sosin et al.\\ 1997a, Piotto et al.\\ 1997), with accompanying ultraviolet images for better study of the extended blue tails that occur on the horizontal branches (HB) of a number of clusters. That program gave some quite interesting results, like the first discovery of extended horizontal branches (EHB) in the metal-rich clusters NGC~6388 and NGC~6441 (Rich et al.\\ 1997), and the discovery in NGC~2808 of an EHB extending down the helium-burning main sequence, with a multimodal distribution of stars along it (Sosin et al.\\ 1997b). GO-6095 also persuaded us that it would be more profitable to explore a larger number of clusters more rapidly in \\( B \\) and \\( V \\) alone, in order to delineate the upper parts of their CMDs, and especially their HBs, with the intention of returning to the more interesting clusters for more intensive follow-up studies.  We therefore changed our program to a snapshot study (GO-7470) aimed at producing CMDs down to a little below the main-sequence (MS) turnoff for all Galactic globular clusters with apparent \\( B \\) distance modulus \\( \\leq18 .0 \\) and whose centers had not yet been observed by HST in a comparable way---53 clusters in all. Since snapshot programs have no guarantee of being completed in a given year, we have resubmitted each year the list of clusters not yet observed (GO-8118, GO-8723). By the time we are writing this paper, only one of the GGCs in our original list remains to be observed: NGC~6779. There were also in the HST archive images of a number of clusters that could be treated similarly. We have measured these in the same way as our own images, and present here CMDs of a total of 74 GGCs, measured and reduced in a uniform way, all observed with WFPC2 with the same filter set, and with the PC centered on the cluster center. For all of the 74 GGCs, we ran artificial star experiments to measure the completeness of the star counts in all the relevant branches of the CMD. In this paper we describe the observations, the reduction procedures, the artificial star tests, and the steps followed to transform the instrumental magnitudes into magnitudes in both the HST flight and standard Johnson photometric systems.  The data set that we present has already been shown to be extremely valuable in attacking a number of still-open topics on evolved stars in GGCs.  In particular, these data have been used by Piotto et al.\\ (1999a) to investigate the problem of the EHBs and by Raimondo et al.\\ 2002) to study the properties of the red HB in metal-rich clusters; in Zoccali et al.\\ (1999) and Bono et al.\\ (2001) we have throughly discussed the red giant branch (RGB) bump, and compared the predicted position and dimension of this feature with the observed ones; in Zoccali et al.\\ (2000) we have used the star counts on the HB and on the RGB to gather information on the helium content and on the dependence of the helium content on the cluster metallicity; in Zoccali \\& Piotto (2000) we have made the most extensive comparison so far available between the model evolutionary times away from the main sequence and the actual star counts on the subgiant branch (SGB) and RGB; in Cassisi et al.\\ (2001) we have compared the observed and theoretical properties of the asymptotic giant branches; in Piotto et al.\\ (2000) we have used the CMDs in our database for the study of the GGC relative ages, and, finally, in Piotto et al.\\ (1999b), we have started to investigate the GGC blue straggler (BS) population, and have shown how the BS CMD and luminosity function can differ in clusters with rather different morphologies.  Future papers will carry out other detailed studies based on the same data.  Among these, we are presently working on the derivation of the relative ages, following the strategy already delineated in Rosenberg et al.\\ (1999) on a similarly photometrically homogeneous set of CMDs, but from ground-based data, and by Piotto et al.\\ (2000) on a small subsample of the present HST database. We are also working on the large BS database that came from the CMDs presented in the following Sections (Piotto et al.\\ 2002, in preparation).  As better described in Section 3, all the CMDs, the star positions, the magnitudes in both the F439W and F555W flight system, and the \\( B \\) and $V$ standard Johnson system will be made available to the astronomical community immediately after the publication of the present paper. ",
        "conclusions": ""
    },
    "0207/astro-ph0207642_arXiv.txt": {
        "abstract": "The spectral energy distributions (SEDs) of dust disks are widely used to infer dust properties (compositions and sizes) and disk structures (dust spatial distributions) which might be indicative of the presence or absence of planets or smaller bodies like asteroids and comets in the disk. Based on modelling of the SED of $\\hra$, a young main-sequence star with the largest fractional infrared (IR) emission, we show that the SED {\\it alone} is not a good discriminator of dust size, spatial distribution (and composition if no spectroscopic data are available). A combination of SED, mid-IR spectroscopy, and coronagraphic near-IR imaging of scattered starlight and mid-IR imaging of dust thermal emission provides a better understanding of these properties. ",
        "introduction": "} Over the past 2 decades, impressive evidence has been assembled for the existence of circumstellar dust disks around main-sequence (MS) stars as well as pre-MS stars (T Tauri stars and Herbig Ae/Be stars), post-MS stars (red giants), and a white dwarf (see Zuckerman 2001 for a review). A wide variety of observational techniques have been employed to study the formation/evolution and physical/chemical properties of dust disks: optical and near-infrared (IR) imaging of scattered stellar light, photometric measurements of dust thermal emission from near-IR to submillimeter, spectroscopic observations of mid-IR dust emission features and gas emission lines and ultraviolet (UV) and visible gas absorption lines. The spectral energy distribution (SED) is of particular interest in inferring the size and composition of dust grains and the disk structure (dust spatial distribution). However, the limitations of this method have not been adequately explored; for example, in modelling the $\\bp$ SED, Li \\& Greenberg (1998) found that the dust spatial distribution is coupled with the distribution of dust sizes, i.e., the distribution of dust in the disk and dust sizes cannot be uniquely determined simultaneously by the SED alone.  It is the aim of this {\\it Letter} to quantify the limitations on the information derived from the SED modelling concerning the dust properties (composition and size) and disk structure. We take the SED of $\\hra$, a nearby (distance to Earth $d\\approx 67\\pm 3\\pc$) young MS star (age $\\approx 8\\pm 3\\myr$) of spectral type A0\\,V (effective temperature $\\Teff\\approx 9500\\K$) for comparison with model results. $\\hra$ has the largest fractional IR luminosity relative to the total stellar luminosity ($L_{\\rm IR}/L_{\\star}\\approx 5\\times 10^{-3}$) among the $\\sim 1500$ A-type MS stars in the {\\it Yale Bright Star Catalogue} (Jura 1991). The dust disk in orbit around $\\hra$ has been extensively studied, both observationally and theoretically (see Zuckerman 2001 and references therein). With its relatively well determined dust and disk properties, $\\hra$ serves as a good comparison basis for the SED modelling efforts described in this {\\it Letter}. We stress that the main purpose of this {\\it Letter} is not to carry out a detailed study of the $\\hra$ dust disk which we defer to a subsequent paper (A. Li \\& J.I. Lunine 2002, in preparation). We will first outline our approach in \\S\\ref{sec:model}. We then discuss in \\S\\ref{sec:spasize} the degeneracy between the dust spatial distribution and dust sizes under the assumption of pure silicate dust. In \\S\\ref{sec:composition} we show that the dust composition is not well constrained by the observed SED unless mid-IR spectroscopic dust emission features are available. We discuss in \\S\\ref{sec:discussion} the possible dust composition and morphology from the evolutionary point of view that circumstellar dust disks around (pre-)MS stars are formed through the coagulation of interstellar solids. We also summarize our major conclusions in \\S\\ref{sec:discussion}. ",
        "conclusions": "} It is shown in \\S\\ref{sec:spasize} and \\S\\ref{sec:composition} that models with various compositions, sizes, and spatial distributions are able to reproduce the observed SED of the $\\hra$ disk reasonably well. It was also shown by Jura et al.\\ (1998) that icy grains with a typical radius near 100$\\mum$ are able to explain the $\\hra$ SED. Dust thermal emission depends on its absorption and emission properties which are determined by its size and optical properties. It would not be surprising for a wide range of dust materials with properly chosen sizes to be able to fit the SED. However, we should not be too pessimistic: such spectrum will be useful when combined with other constraints on the composition of the dust in circumstellar disks around (pre-)MS stars. The coagulation of interstellar grains that results in fluffy and inhomogeneous aggregates occurs in cold, dense molecular clouds and protostellar and protoplanetary dust disks and plays an important role in the formation of planetary systems (Weidenschilling \\& Cuzzi 1993). We can therefore approximately derive the proportional composition of the dust in circumstellar disks from the abundances of the condensable elements (C, N, O, Si, Fe, and Mg),\\footnote{% Some H will be present, mostly in combination with O, C, and N. } assuming protostellar activities impose little modification on protostellar grain compositions (see Beckwith, Henning, \\& Nakagawa 2000).  Let $\\xsun$ be the cosmic abundance of X relative to H (we assume the cosmic elemental abundances are those of the solar values: $\\csun \\approx 391$ parts per million (ppm), $\\nsun \\approx 85.2\\ppm$, $\\osun \\approx 545\\ppm$, $\\mgsun \\approx 34.5\\ppm$, $\\fesun \\approx 34.4\\ppm$, and $\\sisun \\approx 28.1\\ppm$ [Sofia \\& Meyer 2001]); $\\xgas$ be the amount of X in gas phase ($\\cgas\\approx 140\\ppm$, $\\ngas\\approx 61\\ppm$, $\\ogas\\approx 310\\ppm$; Fe, Mg and Si are highly depleted in dust; see Li \\& Greenberg 1997 and references therein); $\\xdust$ be the amount of X relative H locked up in dust ($\\cdust = \\csun-\\cgas \\approx 251\\ppm$, $\\ndust\\approx 24.2\\ppm$, $\\odust\\approx 235\\ppm$, $\\mgdust\\approx 34.5\\ppm$, $\\fedust\\approx 34.4\\ppm$, $\\sidust\\approx 28.1\\ppm$). Assuming a stoichiometric composition of MgFeSiO$_4$ for interstellar silicates, the total mass of silicate dust per H atom is $\\msil \\approx \\fedust\\mufe + \\mgdust\\mumg + \\sidust\\musi + \\osil\\muo \\approx 5.61\\times 10^{-3}\\,\\muh$ where $\\mux$ is the atomic weight of X in unit of $\\muh\\approx 1.66\\times 10^{-24}\\g$, and $\\osil\\approx 4\\,(\\fedust + \\mgdust + \\sidust)/3\\approx 129\\ppm$ is the amount of O in silicate dust per H atom (i.e., we assign 4 O atoms for the average of the Fe, Mg, and Si abundances). The carbonaceous dust component is dominated by C, with little H, N, and O (we assume H/C=0.5, O/C=0.1). The total mass of carbon dust per H atom is $\\mcarb \\approx \\cdust\\muc + \\ndust\\mun + 0.5 \\cdust\\muh + 0.1 \\cdust\\muo \\approx 3.88\\times 10^{-3}\\,\\muh$. The C, O, and N atoms left over after accounting for the silicate and carbon dust components are assumed to condense in icy grains in the form of H$_2$O, NH$_3$, CO, CO$_2$, CH$_3$OH and CH$_4$ (following Greenberg [1998], we assume CO:CO$_2$:CH$_3$OH:CH$_4$:H$_2$CO=10:4:3:1:1). The total mass of icy grains per H atom is $\\mice \\approx \\mice^{\\rm C} + \\mice^{\\rm N} + \\mice^{\\rm water}$, where the mass of C-containing ice $\\mice^{\\rm C}\\approx \\cgas\\muc + \\cgas\\,(22\\muo+18\\muh)/19\\approx 2.87\\times10^{-3}\\,\\muh$; the mass of NH$_3$ ice $\\mice^{\\rm N}\\approx \\ngas(\\mun+3\\muh) \\approx 4.10\\times 10^{-4}\\,\\muh$; the mass of water ice $\\mice^{\\rm water}\\approx \\owater(\\muo+2\\muh) \\approx 4.12\\times 10^{-3}\\,\\muh$; $\\owater \\approx \\osun-\\osil-0.1\\cdust-22\\cgas/19 \\approx 229\\ppm$ is the amount of O locked up in H$_2$O ice (we assume H$_2$O contains all the remaining available O).  Therefore, as a first approximation, we may assume a mixing ratio of $\\mcarb/\\msil \\approx 0.7$ and $\\mice/(\\msil+\\mcarb) \\approx 0.8$ for cold regions (for hot regions where ices sublimate the dust can be simply modelled as porous aggregates of silicate and carbon particles with $\\mcarb/\\msil \\approx 0.7$). This does not deviate much from the in situ measurements of cometary dust ($\\mcarb/\\msil \\approx 0.5$, $\\mice/[\\msil+\\mcarb] \\approx 1.0$; see Greenberg \\& Li 1999 and references therein) which is often suggested as porous aggregates of unaltered interstellar dust (Greenberg 1982; Greenberg \\& Li 1999). The porosity is a free parameter ranging from that of diffuse cloud interstellar dust ($\\sim 0.45$, Mathis 1996) to that of very fluffy cometary dust ($\\gtsim 0.9$, Greenberg \\& Li 1999). For illustration, we plot in Figure \\ref{fig:comet} the model spectrum calculated from a Gaussian distribution ($\\rin=0.15\\AU$, $\\rout=250\\AU$, $r_0=70\\AU$, $\\Delta=15\\AU$) of grains with (1) a power law size distribution ($\\amin=1\\mum$, $\\amax=1\\cm$, $\\alpha=3.6$), (2) $\\mcarb/\\msil = 0.7$ and a porosity of $P=0.6$ for $r<30\\AU$; and (3) compact icy grains with $\\mcarb/\\msil = 0.7$ and $\\mice/(\\msil+\\mcarb) = 0.8$ (porous grains of $P=0.6$ become compact after filled with ices of an amount of $\\mice/[\\msil+\\mcarb] = 0.8$) for $r>30\\AU$. This will be discussed in detail in a subsequent paper (A. Li \\& J.I. Lunine 2002, in preparation).  The fact that pure amorphous carbon grains are also able to account for the observed SED (see \\S\\ref{sec:composition}) reinforces the importance of combining observations with theoretical calculations of dust composition in the context of the formation and evolution of dust disks. We note that the non-detection of the silicate emission features in the $\\hra$ disk (Sitko et al.\\ 2000) does not necessarily imply the predominance of non-silicate dust in the disk. It may just imply the lack of small and hot silicate grains. On the other hand, the contradistinction between the various dust spatial distributions, which all provide a reasonably good fit to the observed SED (see \\S\\ref{sec:spasize} and \\S\\ref{sec:composition}), indicates the importance of direct disk imaging.  In summary, the spectral energy distributions of dust disks {\\it alone} are not necessarily able to constrain the dust compositions, sizes, and spatial distributions. The dust spatial and size distributions are coupled. Caution should be taken in discussing the presence/absence of planets, comets, and asteroids in the disk {\\it solely} based on the observed SED. We argue that grains in circumstellar disk around (pre-)MS stars are composed of silicate, carbonaceous dust (and ices in cold regions) and vacuum with a mixing ratio of $\\mcarb/\\msil \\approx 0.7$ and $\\mice/(\\msil+\\mcarb) \\approx 0.8$. A combination of compositional considerations, SED, mid-IR spectroscopy, coronagraphic near-IR imaging of scattered starlight and mid-IR imaging of dust thermal emission will allow us to better understand the properties of circumstellar dust disks."
    },
    "0207/astro-ph0207197_arXiv.txt": {
        "abstract": "Identifying the progenitors of Type~Ib and Type~Ic supernovae requires knowing, among other things, whether SNe~Ib eject hydrogen, and whether SNe~Ic eject helium, and perhaps even hydrogen.  Recently it has become clear that some SNe~Ib do eject hydrogen, and it may be that all SNe~Ib do.  Two arguments that have been made in the past that SNe~Ic eject helium are difficult to confirm, but I discuss other possible evidence that SNe~Ic eject helium, as well as hydrogen.  If so, these elements extend to {\\sl lower} ejection velocities than in SNe~Ib.  The spectroscopic differences between SNe~Ib and SNe~Ic may depend on the radial distributions of the helium and hydrogen as well as on the ejected masses of helium and hydrogen.  We should consider the possibility that SNe~Ic are more mixed up. ",
        "introduction": "Type~Ia supernovae are thought to be nuclear--powered disruptions of accreting or merging white dwarfs.  Most if not all other supernovae (SNe) are thought to involve expulsion of the envelopes of massive stars following the gravitational collapse of their highly evolved cores.  Of these, Type~II SNe have conspicuous features due to hydrogen in their optical spectra.  Type~Ib SNe lack conspicuous hydrogen features but have conspicuous He~I features during their photospheric phase.  In the spectra of Type~Ic SNe, neither hydrogen nor He~I features are conspicuous.  For reviews of SN classification see Filippenko (1997) and Turatto (2002).  The optical spectrum of a typical SN~II evolves from a nearly featureless continuum when the temperature is high to one that contains mainly hydrogen features and then gradually develops features of lower excitation due to Ca~II, Fe~II, O~I, Na~I, and Ti~II, plus Ba~II and Sr~II if the temperature gets sufficiently low.  Apart from the defining differences involving the hydrogen and He~I features, the spectra of SNe~Ib and SNe~Ic evolve in much the same way.  Among the questions that those who make explosion models of SNe~Ib and SNe~Ic would like spectroscopists to answer include two very basic ones.  Do SNe~Ib eject any hydrogen?  And do SNe~Ic eject any helium --- and perhaps even some hydrogen?  I will be able to provide only partial and tentative answers to these questions --- but I should at least be able to explain why they are so hard to answer.  \\vfill\\eject ",
        "conclusions": "So --- {\\sl do SNe~Ib eject hydrogen?}  Yes, the deep--H$\\alpha$ events certainly do; they appear to be closely related to SNe~IIb.  In fact, the differences between SNe~IIb and the deep--H$\\alpha$ SNe~Ib can be so small that the classification of some events probably depends on how early the first spectrum is obtained.  [In the nomenclature of Clocchiatti \\& Wheeler (1997) the deep--H$\\alpha$ events could be included among the ``transition supernovae'', SNe~IIt.]  Whether typical SNe~Ib eject hydrogen is much more difficult to decide.  Given the strong similarities between the deep--H$\\alpha$ and typical SNe~Ib (B02), the most natural assumption would seem to be that the 6300~\\AA\\ absorption in typical SNe~Ib also is produced by H$\\alpha$.  This is supported by the probable presence of intermediate--strength 6300~\\AA\\ absorption in the shallow--helium SN~1999ex.  On the other hand, nonthermally excited Ne~I lines are not implausible, and because Ne~I $\\lambda$6402 would not need to be detached it would require one fewer free parameter (the detachment velocity of H$\\alpha$).  The 6300~\\AA\\ absorption tends to drift redward with time like the other features; this may seem to favor undetached Ne~I, but it is not decisive since the detached He~I lines in SNe~Ib also drift to the red.  (B02 offer an explanation.)  How are we to decide about the H$\\alpha$ identification?  The most direct confirmation would be to detect Paschen--alpha (P$\\alpha~\\lambda$18751), but in LTE its Sobolev optical depth is lower than that of H$\\alpha$ by a factor of 34 at 5000~K and 3.4 at 10,000~K, so this may not work.  A larger sample of high quality optical spectra of typical SNe~Ib will allow the following test to be applied: if the 6300~\\AA\\ feature {\\sl always} is consistent with undetached Ne~I, then it probably is Ne~I, because it would be hard to believe that hydrogen always is detached by the same amount with respect to the velocity at the photosphere.  The current data are inconclusive on this point.  {\\sl Do SNe~Ic eject helium?}  As we have seen, the 1~$\\micron$ feature in SN~1994I does not provide convincing evidence for He~I $\\lambda$10830, and the presence of weak optical absorptions with highly detached He~I lines is unlikely.  However, it may be that the optical spectra contain undetached He~I lines.  How are we to decide? Since He~I $\\lambda$10830 can be confused with lines of other ions, the most direct test might be to look for $\\lambda$20581, which in LTE has about the same Sobolev optical depth as $\\lambda$10830 above 10,000~K.  (But be careful --- $\\lambda$20581 is a singlet transition.)  Infrared spectra also could be used to sort out the contributions of C~I and Si~I to the 1~$\\micron$ feature because both ions have other strong lines that should be detectable (Figure~12 of Millard et~al. 1999).  {\\sl Do SNe~Ic eject hydrogen?}  Identifying the weak 6360~\\AA\\ absorption in SN~1987M with H$\\alpha$ may be questionable, but the possibility of emission near the rest wavelength of H$\\alpha$ looks interesting. The hydrogen would need to be undetached.  This, together with the possibility that SNe~Ic contain undetached He~I lines, means that if SNe~Ic do eject helium and hydrogen then these elements extend to {\\sl lower} ejection velocities than in SNe~Ib. The appearance of the hydrogen and He~I features in SNe~Ib and SNe~Ic may be affected by the radial distributions of the helium and hydrogen in addition to the total ejected masses of helium and hydrogen.  We should consider the possibility that the composition structures of SNe~Ic are more mixed up than those of SNe~Ib.  This discussion of line identifications has been based on calculations carried out with the simple SYNOW code.  Detailed NLTE calculations for a grid of explosion models for SNe~Ib and SNe~Ic are needed, especially to test the possibilities of nonthermally excited Ne~I lines in SNe~Ib (and SNe~Ic), and of rounded H$\\alpha$ emission in SNe~Ic.  It also is needed in order to estimate, or place upper limits on, the masses of helium and hydrogen in SNe~Ib and SNe~Ic.  Clocchiatti et~al. (1997) have emphasized that light--curve shapes imply significant differences in the ejected masses of supernovae of the same spectral type.  This, together with our uncertainty about the composition structures of the ejected matter, and our still scanty knowledge of the relative occurrence frequencies of the various kinds of SNe~Ib and SNe~Ic, means that there is much work to be done before we will be able to match the various kinds of SNe~Ib and SNe~Ic with their stellars progenitors in a detailed way."
    },
    "0207/astro-ph0207474_arXiv.txt": {
        "abstract": "{We present a comprehensive X-ray study of four years of pointed {\\it RXTE} observations of \\grs\\ in the $\\chi$-state. We interpret the behavior of the hard power law tail spectrum as coming from inverse Compton scattering of soft disk photons on a thermally dominated hybrid corona above the accretion disk. \\grs\\ shows a strong, variable reflection amplitude. As in other BHC and in Seyfert galaxies, a correlation between the power law slope and the reflection was found. Also, the radio fluxes at 2.25\\,GHz and 15\\,GHz correlate with the power law slope, thus revealing a connection between the outflowing matter and the comptonizing region in the $\\chi$-states. ",
        "introduction": "With the discovery of superluminal ejections (Mirabel \\& Rodriguez 1994) from the galactic transient source \\grs\\ and \\gro, much observational and theoretical attention has been directed to the field of microquasars. Microquasars are thought to be downscaled analogs to quasars exhibiting much smaller time scales and are therefore potential laboratories for studying accretion and relativistic jets near black holes (Mirabel \\etal 1992). The most prominent object of this type is the galactic X-ray binary system \\grs\\ which shows dramatic variability (Greiner \\etal 1996) in light curve, quasi-periodic oscillations, phase lags and coherence behavior (Morgan \\etal\\ 1997, Muno \\etal 2001). \\grs\\ is the most energetic object known in our galaxy with a luminosity of $\\sim$10$^{38}$ erg/s in the low state.  \\grs\\ harbors a black hole of 14\\,\\msun\\ (Greiner \\etal 2001b) making it the most massive stellar black hole known. Being in the galactic plane and at a distance of $\\sim$12~kpc the source suffers a very high extinction in the optical band of the order of 25-30\\,mag (Greiner \\etal 1994, Chaty \\etal 1996). Greiner \\etal\\ (2001a) found the donor to be an K-M III late type giant as determined from absorption line measurements in the near infrared.  Black hole transients generally exhibit different states of intensity and spectrum: the high/soft state with a prominent disk component and a weak or negligible steep power law tail, the low/hard state with negligible accretion disk and flat power law. In the intermediate and very high state both the accretion disk and a power law are seen.  Originally discovered by {\\it Granat} (Castro-Tirado \\etal 1992), \\grs\\ was extensively monitored by the {\\it RXTE} since 1996 and a number of investigations of these data have been published (e. g. Morgan \\etal 1997, Belloni \\etal 1997 \\& 2000, Muno \\etal 1999 \\& 2000)  The source evades simple classification, although it seems to spend most of the time in the very high state. Several attempts to categorize the behavior of \\grs\\ have been made in the past. Belloni \\etal (2000) defined 12 different states depending on light curve variation and hardness colors. One of these states, the so-called $\\chi$-state is characterized by a lack of obvious variations in the light curve and spectrum and is associated with continuous radio emission of differing strength. $\\chi$-states correspond to the low/hard state of \\grs, exhibiting relatively low flux from the accretion disk and showing a hard power law tail. Theses states resemble long variants of the lulls in $\\beta$-states, when part of the inner accretion disk is proposed to be absent. Depending on the strength of the radio emission, differing phase lag and power spectrum behavior are seen. During the so-called radio quiet $\\chi$-states the phase lag of the hard X-ray photons to the soft X-ray photons is partly negative, whereas it is always positive when the radio emission is strong (Muno \\etal 2001). Also, the frequency of the 0.5--10\\,Hz QPO decreases with increasing radio flux.  In this paper we present a comprehensive study of the X-ray spectral behavior of \\grs\\ in the $\\chi$-state with over nearly four years of observation with {\\it RXTE}. The variation of the spectral properties is then compared to those of the radio emission. Valuable information about the disk geometry and properties of the compact object is derived. ",
        "conclusions": "The analysis of 139 {\\it RXTE} observations of \\grs\\ in the $\\chi$-state revealed variable components and parameters of the X-ray spectra. The structure of the hard X-ray power law depends e. g. on the radio flux. Further, a two-branch correlation of the power law slope and the power law normalization was found. Theses branches show different pivoting behavior.  The most probable geometry is that of a hot corona above the accretion disk. A bulk motion comptonization seems to be ruled out for the $\\chi$-states because of the time lags of hard and soft X-ray photons.  The continuous outflow of matter is not correlated with the accretion disk parameters as measured at X-rays. Neither the disk temperature nor the inner disk radius are connected with the radio flux. Therefore, the outflowing matter should not be provided by a disruption of the inner part of the accretion disk, as thought to be the case for the sporadic jets in \\grs, and a positive correlation between radio emission and power law slope is found. Because both the base of the jet and the corona are probably located near the black hole the correlation indicates an interaction of both structures. Whereas during low radio emission the hard X-ray component originates due to comptonization on a thermally-dominated corona, in radio loud states the comptonization should appear in the outflowing matter.  \\grs\\ shows a positive $R$($\\Gamma$)-correlation as seen in other X-ray binaries and AGN. The large reflection amplitude suggests a highly anisotrophic inverse Compton scattering with the dominant part of the soft photons being scattered back into the disk plane."
    },
    "0207/astro-ph0207018_arXiv.txt": {
        "abstract": "{ We present multi-frequency VLA continuum observations towards 8 star  forming regions with molecular and optical outflows: L1489, \\hhs, \\hhn, \\ngc, L1681B, L778, MWC~1080 and V645~Cyg.  We detect three thermal radio jets, L1489, YLW~16A in L1681B and NGC~2264D VLA~7, associated with molecular and/or HH outflows.  The L1489 and NGC~2264D VLA~7 thermal radio jets appear elongated in the direction of the larger scale outflow.  We report the first tentative detection of a non-thermal radio jet, L778 VLA~5, associated with a low mass Class I protostar and powering a molecular outflow.  For \\hhs, \\hhn\\ and the molecular outflow in \\ngc\\ we could not identify a candidate of the exciting source of these outflows.  The radio emission associated with \\vcy\\ is quite extended, $\\sim 0.1$~pc, and time variable.  We detect three radio sources in the \\mwc\\ that could be associated with YSOs. }   \\begin{document} ",
        "introduction": "}   \\begin{table*}[t] \\scriptsize \\caption{Observed regions} \\label{obs} \\[ \\begin{tabular}{lcccrcrrc} \\hline \\multicolumn{1}{c}{} & \\multicolumn{2}{c}{Phase Center} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{2}{c}{Synthesized Beam} & \\multicolumn{1}{c}{rms} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{2}{c}{\\hrulefill} & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{Phase} & \\multicolumn{2}{c}{\\hrulefill} & \\multicolumn{1}{c}{Noise} \\\\ \\multicolumn{1}{c}{Region} & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} & \\multicolumn{1}{c}{cm} & \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{Calibrator} & \\multicolumn{1}{c}{HPFW} & \\multicolumn{1}{c}{PA} & \\multicolumn{1}{c}{$\\mu$Jy~beam$^{-1}$} \\\\ \\hline L1489 & 04 04 42.9 &  +26 18 56 & 3.6 & 28/05/90   &  0403+260  & $0\\farcs42\\times0\\farcs29$ & $-84\\arcdeg$ & 15 \\\\ \\hhs\\ & 05 41 38.2 &$-$06 27 44 & 3.6 & 07/07/95   &  0550+032  & $0\\farcs33\\times0\\farcs25$ & $-19\\arcdeg$ & 24 \\\\ \\hhn\\ & 05 43 39.0 &$-$02 35 09 & 2 & 26/01/90   & 0541$-$056 & $6\\farcs36\\times4\\farcs87$ & $-11\\arcdeg$ & 63 \\\\ &&& 3.6 & 28/05/90   & 0541$-$056 & $0\\farcs39\\times0\\farcs36$ & $+72\\arcdeg$ & 15 \\\\ &&& 3.6 & 05/01/97   & 0539$-$057 & $0\\farcs32\\times0\\farcs27$ & $+25\\arcdeg$ & 10 \\\\ &&& 6 & 31/07/89   &  0550+032  & $6\\farcs88\\times3\\farcs95$ & $-30\\arcdeg$ & 34 \\\\ \\ngc\\ & 06 41 04.5 &  +09 36 20 & 2 & 26/01/90   &  0725+144  & $5\\farcs74\\times4\\farcs94$ & $+10\\arcdeg$ & 59 \\\\ &&& 3.6 & 16/07/92   &  0629+104  &$10\\farcs15\\times8\\farcs12$ & $ -9\\arcdeg$ & 42 \\\\ &&& 3.6 & 07/07/95   &  0550+032  & $0\\farcs30\\times0\\farcs25$ & $-45\\arcdeg$ & 24 \\\\ &&& 6 & 31/07/89   &  0550+032  & $4\\farcs74\\times4\\farcs16$ & $-34\\arcdeg$ & 41 \\\\ L1681B& 16 27 28.0 &$-$24 39 33 & 3.6 & 28/05/90   & 1626$-$298 & $0\\farcs51\\times0\\farcs21$ & $+23\\arcdeg$ & 17 \\\\ L778  & 19 26 28.9 &  +23 56 53 & 2 & 26/01/90   &  1923+210  & $5\\farcs27\\times4\\farcs89$ & $ -3\\arcdeg$ & 57 \\\\ &&& 3.6 & 28/05/90   &  1923+210  & $0\\farcs36\\times0\\farcs27$ & $+49\\arcdeg$ & 16 \\\\ &&& 6 & 27/07/89   &  1923+210  & $4\\farcs43\\times3\\farcs91$ & $-25\\arcdeg$ & 33 \\\\ \\vcy\\ & 21 39 58.1 &   +50 14 20 & 2 & 26/01/90   &  2200+420  & $5\\farcs37\\times4\\farcs95$ & $+12\\arcdeg$ & 66 \\\\ &&& 3.6 & 31/01/94   &  2146+608  & $11\\farcs1\\times10\\farcs6$ & $-66\\arcdeg$ & 30 \\\\ &&& 6 & 27/07/89   &  2200+420  & $4\\farcs72\\times4\\farcs30$ & $+15\\arcdeg$ & 32 \\\\ \\mwc\\ & 23 17 27.4 &   +60 50 48 & 2 & 26/01/90   &  0014+612  & $5\\farcs88\\times4\\farcs90$ &  $-2\\arcdeg$ & 80 \\\\ &&& 6 & 27/07/89   &  2352+495  & $5\\farcs50\\times5\\farcs03$ & $+55\\arcdeg$ & 35 \\\\ \\hline \\end{tabular} \\] \\end{table*}    It is well established that in the early stages of star formation there is a large mass loss in young stellar objects (YSOs). The two  most spectacular manifestations of this phenomenon are the Herbig-Haro (HH thereafter) outflows and the bipolar molecular outflows. HH outflows are shocks excited by highly collimated, fast winds coming from YSO (\\eg\\ Hartigan \\et\\ 2000), while bipolar molecular outflows are likely ambient gas swept up by those highly collimated winds (\\eg\\ Richer \\et\\ 2000).  The energy sources of most molecular and HH outflows are surrounded by large amounts of gas and dust, which contribute significantly to their spectral energy distribution and produce such a large extinction that YSOs are generally invisible at optical wavelengths (\\eg\\ Andr\\'e 1997). Sometimes, they are so deeply embedded that in spite of the recent developments in instrumentation at near-infrared wavelengths, they are not easily detected even at these wavelengths (\\eg\\ Lada \\& Lada 1991). Therefore, observations in the mid and far infrared (and longward) wavelengths are needed to identify the YSOs driving the molecular and HH outflows. Submillimeter and millimeter observations are probably the most useful techniques to identify YSOs, since they usually exhibit strong dust emission at these wavelengths, and the mm interferometers can achieve high angular resolution (\\eg\\ Wilner \\& Lay 2000; \\ro\\ \\et\\ 1998).   \\begin{table*}[t] \\scriptsize \\caption[]{Sources Detected from the Matching-Beam Observations\\tablenotemark{\\lowercase{a}}} \\label{tmatch} \\[ \\begin{tabular}{lcllccl} \\hline \\noalign{\\smallskip} \\multicolumn{4}{c}{} & \\multicolumn{1}{c}{6 cm Flux} & \\multicolumn{1}{c}{2 cm Flux} & \\multicolumn{1}{c}{} \\\\ \\multicolumn{1}{l}{Region} & \\multicolumn{1}{l}{VLA} & \\multicolumn{1}{c}{\\mmm$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} & \\multicolumn{1}{c}{(mJy)} & \\multicolumn{1}{c}{(mJy)} & \\multicolumn{1}{l}{Identification} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\hhn: &(1)& 05 43 28.36 &$-$02 35 26.3 & $0.50\\pm0.09$ &$opb$&$bg$ \\\\ &(2)& 05 43 39.18 &$-$02 35 10.1 & $0.30\\pm0.07$ & $0.39\\pm0.12$ & \\\\ &(4)& 05 43 57.60 &$-$02 34 54.4 & $1.53\\pm0.07$ &$opb$& \\\\ &(5)& 05 44 15.50 &$-$02 45 35.1 &\\tablenotemark{b}&$opb$&$bg$ \\\\ &(6)& 05 44 47.19 &$-$02 35 06.1 &\\tablenotemark{b}&$opb$& PMN J0544-0234?\\\\ \\ngc: &(1)& 06 40 48.14 &  +09 33 05.0 &  $3.6\\pm0.3$  &$opb$&$bg$ \\\\ &(2)& 06 40 50.47 &  +09 32 19.6 &  $0.9\\pm0.2$  &$opb$&$bg$ \\\\ &(3)& 06 40 51.96 &  +09 31 54.0 & $18.6\\pm0.3$  &$opb$&$bg$ \\\\ &(4)& 06 40 52.64 &  +09 29 54.0 & $20.0\\pm0.5$  &$opb$&$bg$ \\\\ &(5)& 06 40 56.91 &  +09 38 42.1 & $2.78\\pm0.08$ &$opb$&$bg$ \\\\ &(7)& 06 41 04.52 &  +09 36 20.5 & $0.72\\pm0.04$ & $1.12\\pm0.08$ & \\\\ &(12)&06 41 45.14 &  +09 47 03.0 &\\tablenotemark{b}&$opb$ & 4C09.25 \\\\ L778: &(1)& 19 26 02.69 &  +23 55 07.1 & $10.8\\pm0.1$  &$opb$&$bg$ \\\\ &(2)& 19 26 02.76 &  +23 54 57.1 & $11.0\\pm0.1$  &$opb$&$bg$ \\\\ &(3)& 19 26 12.64 &  +23 54 01.8 &  $1.0\\pm0.1$  &$opb$&$bg$ \\\\ &(4)& 19 26 22.84 &  +23 53 46.8 &  $0.3\\pm0.1$  &$opb$&$bg$ \\\\ &(5)& 19 26 28.77 &  +23 56 52.7 & $1.01\\pm0.05$ & $0.22\\pm0.06$ &IRAS~19243+2350 \\\\ &(6)& 19 26 29.15 &  +23 56 18.6 & $0.71\\pm0.05$ & $0.79\\pm0.09$ &    \\\\ &(7)& 19 26 30.11 &  +23 55 14.0 & $8.83\\pm0.06$ & $2.66\\pm0.24$ &$bg$ \\\\ V645~Cyg: &(1)& 21 39 22.60 &  +50 16 58.7 &  $8.2\\pm0.2$  &$opb$&$bg$ \\\\ &(2)& 21 39 22.86 &  +50 10 24.7 &  $4.4\\pm0.3$  &$opb$&$bg$ \\\\ &(3)& 21 39 32.92 &  +50 09 08.5 &  $3.1\\pm0.2$  &$opb$& IRAS~21377+4955 \\\\ &(4)& 21 39 33.36 &  +50 09 10.8 &  $2.1\\pm0.3$  &$opb$& \\id\\ \\\\ &(5)& 21 39 43.62 &  +50 15 16.8 & $1.72\\pm0.06$ &$opb$&$bg$ \\\\ &(6)& 21 39 58.24 &  +50 14 21.5 & $0.57\\pm0.04$ & $1.04\\pm0.22$ & V645~Cyg \\\\ &(7)& 21 40 00.32 &  +50 06 51.6 &  $3.9\\pm0.5$  &$opb$&$bg$ \\\\ &(8)& 21 40 00.87 &  +50 13 39.6 & $8.43\\pm0.06$ & $2.89\\pm0.15$ &$bg$ \\\\ MWC~1080: &(1)& 23 16 46.83 &  +60 53 19.9 &  $4.1\\pm0.2$  &$opb$&$bg$ \\\\ &(2)& 23 17 20.36 &  +60 48 20.8 & $0.43\\pm0.08$ &$opb$&$bg$ \\\\ &(3)& 23 17 24.14 &  +60 50 44.9 & $0.17\\pm0.04$ &$\\la0.32$ & \\\\ &(4)& 23 17 25.52 &  +60 50 42.9 & $0.21\\pm0.07$ &$\\la0.32$ & MWC~1080? \\\\ &(5)& 23 17 27.43 &  +60 50 48.9 & $0.20\\pm0.05$ &$\\la0.32$ & \\\\ &(6)& 23 17 34.48 &  +60 56 43.1 &     $\\sim18$  &$opb$&$bg$ \\\\ \\hline \\end{tabular} \\] $^a$ ``$opb$'' indicates sources outside of the 2~cm primary beam. ``$bg$'' means a likely background source. \\\\ $^b$ Sources outside the primary beam at 6~cm. \\\\ \\end{table*}    An alternative way to identify this type of sources is through interferometric radio continuum observations at centimeter wavelengths, carried out mainly with the VLA (and more recently with MERLIN and the Australia Telescope). The VLA allows to map large regions (e.g. the primary beam at 6~cm is  9$'$) with a very high sensitivity. In addition, the spectral indices of the  sources can be measured from multi-frequency radio continuum observations  and, therefore, it allows to elucidate the nature of the radio emission. Combining this type of observations with other criteria, such as the source coinciding with the geometrical center of the outflow and its association with an infrared and/or millimeter counterpart, has proven  to be a very useful tool to discriminate among the candidates of the energy  source of the outflow (\\eg\\ HH~1-2: Pravdo \\et\\ 1985; L1448:  Curiel \\et\\ 1990; L1287: Anglada \\et\\ 1994).   Recently, a number of surveys at centimeter wavelengths have been  carried out in order to identify the powering sources of molecular outflows (\\eg\\ Anglada \\et\\ 1992, 1998;  Beltr\\'an 2001) and of HH objects (\\eg\\ \\ro\\ \\& Reipurth  1994, 1998; Avila, \\ro, \\& Curiel 2001).  For wavelengths longer than $\\sim$1~cm, the continuum emission of the energy sources of molecular and HH outflows is often dominated by free-free emission from partially ionized outflows (\\eg\\ Anglada 1996).  These cm radio observations allow to determine with great accuracy the position of the exciting source and to determine the morphology and other physical parameters of the ionized gas at small angular scales.  Thus, the  radio continuum sources associated with YSOs, powering molecular and/or HH outflows, are usually found to have the following characteristics (\\eg\\  Anglada 1996, \\ro\\ 1997): {\\it (1)} Relatively weak flux  densities in the cm regime (around 1 mJy or less); {\\it (2)} Spectral indices  that are flat or rise slowly with frequency (typically between -0.1 and 1);  {\\it (3)} No evidence of large time variability; {\\it (4)} No evidence of  polarization; and {\\it (5)} in some of the best studied cases they exhibit a  jet-like morphology, with their orientation, in most cases, along the  molecular or HH outflow direction. Because of these properties, these objects  are known in the literature as ``thermal radio jets''.  In this paper we present matching--beam VLA observations at 2 and 6~cm and sub-arcsecond angular resolution observations towards several star forming regions with associated molecular outflows and/or  Herbig-Haro objects. Most of the selected regions were previously observed with  the VLA at only one wavelength and with lower angular resolution. ",
        "conclusions": "Table~\\ref{tsize} shows the radio continuum candidates to power the molecular and/or HH outflows in the regions observed with the VLA.  Of the observed regions, HH~68-69 is the only one that is not shown in this table because of the lack of radio continuum emission from our observations. The main results from our observations are:  (1) We detected three radio sources, whose properties are consistent with being thermal radio jets: L1489, \\ngc\\ VLA~7 and YLW~16A. The length of these radio jets are $\\sim$40~AU for L1489 and YLW~16A and $\\sim$370~AU for \\ngc\\ VLA~7, which are typical values for this type of objects (Anglada 1996). These three radio jets are associated with HH and molecular outflows. L1489 coincides with a compact molecular outflow (Hogerheijde \\et\\ 1998) and with HH~360 (G\\'omez \\et\\ 1997). \\ngc\\ VLA~7 is associated with the system HH~125/225/226 (Walsh, Ogura, \\& Reipurth 1992).  For these two regions, the thermal radio jet is well aligned with the molecular or HH outflow.  The thermal radio jet in YLW~16A is associated with a pole-on molecular outflow (Sekimoto \\et\\ 1997).  (2) We detected a non-thermal radio jet in L778 that appears to be associated with a Class I infrared source, IRAS~19243+2350. This radio jet is elongated in the direction of a pair of red and blue high velocity CO lobes (Myers \\et\\ 1988) centered roughly in IRAS~19243+2350.  Thus, we suggest the first tentative detection of a  non-thermal radio jet associated with a low mass protostar.  (3) Our observations could not find clear candidates for the \\hhn\\ system and the molecular outflow in \\ngc. In both regions there are several radio sources that could trace the powering source of the outflows, but the lack of known counterpart or the incomplete information of their radio emission properties does not allow to discriminate between them.  (4) \\vcy\\ shows radio emission with striking properties: its emission is quite extended, $\\sim23000$~AU, but at the same time it is variable. There is also an extended source that appears to be an optically thin HII region.  (5) There is a $\\sim 3\\farcs5$ offset between the optical source \\mwc\\ and VLA~4, so it is not clear if VLA~4 is directly associated with the star.  There are two other radio sources, VLA~3 and 5, that are associated with mm sources and could be tracing protostars."
    },
    "0207/astro-ph0207043_arXiv.txt": {
        "abstract": "We describe the method we have used to parallelize our spherically symmetric special relativistic short characteristics general radiative transfer code PHOENIX. We describe some possible parallelization strategies and show why they would be inefficient. We discuss the multiple parallelization strategy techniques that we have adopted. We briefly discuss generalizing these strategies to full 3-D (spatial) radiation transfer codes. ",
        "introduction": "In general parallelization is a subject to be avoided (as nicely described by P.~H\\\"oflich, this volume), however in order to take advantage of the enormous computing power and vast memory sizes of modern parallel supercomputers, allowing both much faster model calculations as well as more detailed models, we have implemented a parallel version of the general model atmosphere code {\\tt PHOENIX} (Hauschildt \\& Baron 1999 and references therein). Since the code uses a modular design, we have implemented different parallelization strategies for different modules in order to maximize the total parallel speed-up of the code. Our implementation allows us to change the load distribution onto different processor elements (PEs) both via input files and dynamically during a model run, which gives a high degree of flexibility to optimize the performance on a number of different machines and for a number of different model parameters.  Since we have both large CPU and memory requirements we have implemented the parallel version of the code on  using the {\\tt MPI}\\ message passing library (MPI Forum 1995).  We have chosen to work with the {\\tt MPI}\\ message passing interface, since it is both portable (public domain implementations of {\\tt MPI} are readily available cf.~Gropp et~al. 1996), running on dedicated parallel machines and heterogeneous workstation clusters and it is available for both distributed and shared memory architectures.  For our application, the distributed memory model is in fact easier to code than a shared memory model, since then we do not have to worry about locks and synchronization, etc.\\ on {\\em small} scales and we, in addition, retain full control over interprocess communication. This is especially clear once one realizes that it is fine to execute the same code on many PEs as long as it is not too CPU intensive, and avoids costly communication. We have added a few simple OpenMP directives to our code, but do not discuss SMP parallelization further here.   An alternative to an implementation with {\\tt MPI} is an implementation using High Performance Fortran ({\\tt HPF}) directives (in fact, both can co-exist to improve performance). However, the process of automatic parallelization guided by the {\\tt HPF} directives is presently not yet generating optimal results because the compiler technology is still very new. {\\tt HPF} is also more suited for problems that are purely data-parallel (SIMD problems) and would not benefit much from a MIMD approach. An optimal {\\tt HPF} implementation of {\\tt PHOENIX} would also require a significant number of code changes in order to explicitly instruct the compiler not to generate too many communication requests, which would slow down the code significantly. The {\\tt MPI} implementation requires only the addition of a few explicit communication requests, which can be done with a small number of library calls. ",
        "conclusions": "While the parallelization of a large numerical code such as {\\tt PHOENIX} would ideally be performed at the compiler level, it is not possible with modern compilers to construct efficient code with good scaling properties. Thus, individual modules of the code have to be examined and parallelized separately. With {\\tt MPI} this task is manageable and allows the efficient use of both commodity parallel clusters and modern massively parallel supercomputers.  We found good speedup up to about 64 PEs for a typical supernova calculation.  However, for more than 64 PEs the communication, synchronization, and loop overheads begin to become significant and it is not economical to use more than 128 PEs For static models and opacity tables we are able to use very large numbers of PEs with scaling at close to the theoretical maximum. Fully 3-D calculations will present yet another parallelization challenge."
    },
    "0207/astro-ph0207569_arXiv.txt": {
        "abstract": "{We present new evidence, based on faint {\\it HST} proper-motions, for a bi-modal kinematic population of old white dwarfs, representative of the Thick-Disk and Halo of our Galaxy. This evidence supports the idea of a massive Halo comprised of faint and old white dwarfs, along with an extant population of Thick-Disk white dwarfs. We show how most of the required dark matter in the solar vicinity can be accounted for by the remnants from these two components together. ",
        "introduction": "The claim by Oppenheimer et al. (\\cite{oppenheimer}) that they had found a significant population of Halo stars from a kinematic survey toward the South Galactic Pole opened up an interesting discussion regarding the nature of dark matter in the solar neighborhood. Their discovery seemed to corroborate earlier findings using very deep {\\it HST} photometry that also pointed out to the existence of an as yet unobserved component to the Halo, identified as very cool and ancient white dwarfs (M\\'endez and Minniti 2000). However, several authors have indicated that Oppenheimer's sample could be also interpreted as the tail of a 'warmer' white dwarf component (in the sense of having shorter cooling ages and therefore higher surface temperatures), better ascribed to the intermediate Thick-Disk population of the Galaxy (see e.g., Reid et al. 2001; Reyl\\'e et al. 2001).  Obviously, disentangling the true nature of Oppenheimer's objects is important to understand their contribution to the dark matter problem in the Galaxy. In this paper we use {\\it HST} data, deep photometry and proper-motions, to demonstrate that there is indeed a very important population of what appears to be Thick-Disk white dwarfs with a large mass density in the solar vicinity. These Thick-Disk objects could not have been identified before on the Hubble Deep Fields South and North due to their high Galactic latitude ($|b| \\ge 50^o$), which prevented the appearance on these fields-of-view of an important population of these objects, characterized by a steep concentration toward the Galactic plane.  A fundamental discovery presented in this paper is that, regardless of whether the objects found are Thick-Disk or Halo white dawrfs, they together account for most of the dark matter in the solar neighborhood. ",
        "conclusions": "The stellar and mass densities derived above imply that the Thick-Disk white dwarf stars, a population {\\it unaccounted for in the Hubble Deep Fields South and North due to their high Galactic latitude}, contributes to about 85\\% of the missing mass in the solar neighborhood, whose density in the solar neighborhood is estimated to be $\\rho_{{\\sun}_{DM}} \\sim 1.26 \\times 10^{-2} \\, M_{\\sun}/pc^3$.  Additionally, the value derived for only one possibly Halo star, albeit having a large error, is similar to the value derived by MM00 from the larger sample of Halo white dwarfs found in the HDF-S\\&N. They found a density of $\\rho_{{\\sun}_{WDH}} = 4.64 \\times 10^{-3} \\, M_{\\sun}/pc^3$. Their value, however, has to be taken with caution because the lack of kinematic information prevented a clear separation of the populations (compare Figs.~\\ref{rpmfield} and~\\ref{cmd} above). It has been mentioned already that one might expect one or two Thick-Disk stars on the HDF-S and one or none in the HDF-N. If these objects were inadvertently identified by MM00 as Halo stars, then the estimated density for these objects on the Halo would have been too large. It is interesting to point out that on HDF-S there are two stars on the WD tracks that are quite a bit brighter than the rest of the blue-faint group. In the HDF-N there seems to be a grouping of faint WDs, and a brighter one, which could also be ascribed to the Thick-Disk (see Fig.~3 on MM00). In any case, proper-motions for both the HDF-S\\&N sample would help settle this uncertainty.  If we combine the present results for Thick-Disk stars, with the earlier results from MM00 for Halo stars, the overall density on WD stars for both populations seems to be able to account for the whole of the (local) missing mass in the Galaxy. The existence of a massive component of Thick-Disk evolved white dwarfs has been proposed recently by Gates and Gyuk~(\\cite{gates}) to help explain some of the (many) difficulties with a massive Halo of ancient stars. However, our results suggest that, in addition to this Thick-Disk component, there would also be a quite massive Halo of remnant stars. Even though their local densities are similar, the mass locked in each population is quite different, in account of their different stellar density functions (double exponential {\\it vs.} power-law). The simple integration of these density functions over the whole Galaxy lead to masses of:  \\[ M_{WD-TD} = 4 \\pi \\rho_{{\\sun}_{WDTD}} h_z h_R^3 \\sim 10^{9} \\, M_{\\sun} \\] ,and,  \\[ M_{WD-H} = 4 \\pi \\rho_{{\\sun}_{WDH}} R_{\\sun}^3 \\, \\ln (R_{max}/R_{min}) \\sim 9 \\times 10^{10} \\, M_{\\sun} \\] for the Thick-Disk and Halo components respectively. $R_{max} \\sim 20$~kpc and $R_{min} \\sim 1$~kpc in the last equation are the overall extension of the Halo and the Halo core-radius respectively.  The total mass in these components is consistent with the mass required to produce the MACHO events, as indicated by Gates and Gyuk~(\\cite{gates}), and is also not inconsistent with the total mass estimates for the Galaxy out to a distance of 20~kpc. We also note that with the small scale height of the Thick-Disk population ($h_z \\leq 1.5$~kpc), their contribution to the optical depth of microlensing events would be quite small, and thus the primary sources of the MACHO events would still be primordially the Halo white dwarfs.  The possibility of a white dwarf dominated (dark) halo has been criticized from many fronts (see e.g. Richer~1999 and Gates and Gyuk~2001). Among the most important criticisms is that the precursor of these stars would have produced metals at a rate greater than observed (Gibson and Mould~1997) or that Galaxy halos at high redshift would be brighter than observed due to the the white dwarf precursors (Charlot and Silk~1995). While these are no doubt important issues, Chabrier (1999) and more recently Fontaine et al. (2001, see also Fontaine's C.S. Beals Lecture at http://www.astro.ubc.ca/WD\\_workshop/talks/index.html) have shown that many of these points can either be overcome, or are doubtful criticisms of this scenario. One example is the chemical evolution problem: The yield of a zero-metallicity stars is largely uknown, nor has there been much modelling of the structure of these stars (e.g., a zero metallicity star never undergoes a helium flash, which must have an effect on the the yield of metals in the PNe phase).  Even if the problems described above with this scenario persist, as pointed out by Lynden-Bell and Tout~(\\cite{lynden}) in the first 'Russell lecture' of the new millennium, they could be avoided if the objects found are 'pristine' white dwarfs, objects that never underwent nuclear reactions nor followed the 'normal' path of stellar evolution. These objects (with masses in the range $ 0.2 \\leq M/M_{\\sun} \\leq 1.1$) could have formed gradually and very slowly by accretion onto planetary-sized precursor bodies in low-density regions such as cooling flows and galaxy halos. By radiating their energy before collapsing, these bodies would grow in mass resting on the zero-point energy of confined electrons following the uncertainty principle, and never getting a temperature high enough to start nuclear burning. The difficulty with this scenario is that, although theoretically possible, preliminary results suggest that the required accretion rates are too slow to allow the formation of these objects in less than a Hubble time, which keeps the puzzle regarding the origin of these objects still open.  As we have seen before, the assignement of the bulk of the faint blue stars found here to the Thick-Disk is actually irrelevant in the sense that one could as well have used the classical $1/V_{\\rm max}$ method to derive space densities, without regard to the origin or association to a given stellar population of these stars. This is an important issue in the context of whether some of the stars found by Oppenheimer et al.~(\\cite{oppenheimer}) are either Thick-Disk or Halo stars. The controversy in this regard seems to come from using white dwarf cooling ages while ignoring the main sequence lifetimes of the progenitors, which for many of these stars will be substantial. Indeed, Oppenheimer's sample seems to come from a population of stars formed in a single burst of stars formation between 11 and 14~Gyr ago, as indicated by the luminosity function of these stars, which is a much more direct indication of the age of the population than the cooling ages (see discussion on http://research.amnh.org/users/bro). Another indication that perhaps most of Oppenheimer's stars actually do belong to the Halo is provided by the recent work of Koopmans and Blanford (2001) that have used a maximum-likelihood analysis to show that these stars are consistent with a kinematically distinct flattend Halo population at the more than 99\\% confidence level.  Albeit our objects are faint, it would be quite interesting to acquire U-band photometry with HST. In this case, true white dwarfs will stand out from other objects in a UBV diagram. Additionally, at least the two brightest blue objects on Fig.~\\ref{cmd} are whithin the capabilities of current 8-m class telescopes for an spectroscopic follow-up. This is an interesting possibility that deserves further work."
    },
    "0207/astro-ph0207275_arXiv.txt": {
        "abstract": "We report on a detailed morphological and kinematic study of the isolated non-barred nearby Seyfert 2 galaxy NGC 2110. We combine Integral Field optical spectroscopy, with long-slit and WFPC2 imaging available in the HST archive to investigate the fueling mechanism in this galaxy. Previous work (Wilson \\& Baldwin 1985) concluded that the kinematic center of the galaxy is displaced $\\sim$ 220 pc from the apparent mass center of the galaxy, and the ionized gas follows a remarkably normal rotation curve. Our analysis based on the stellar kinematics, 2D ionized gas velocity field and dispersion velocity, and high spatial resolution morphology at V, I and H$\\alpha$ reveals that: 1) The kinematic center of NGC 2110 is at the nucleus of the galaxy. 2) The ionized gas is not in pure rotational motion. 3) The morphology of the 2D distribution of the emission line widths suggests the presence of a minor axis galactic outflow. 4) The nucleus is blue-shifted with respect to the stellar systemic velocity, suggesting the NLR gas is out-flowing due to the interaction with the radio jet. 5) The ionized gas is red-shifted $\\sim$ 100 km/s over the corresponding rotational motion south of the nucleus, and 240 km/s with respect to the nuclear stellar systemic velocity. This velocity is coincident with the HI red-shifted absorption velocity detected by Gallimore et al (1999). We discuss the possibility that the kinematics of the south ionized gas could be perturbed by the collision with a small satellite that impacted on NGC 2110 close to the center with a highly inclined orbit. Additional support for this interpretation are the radial dust lanes and tidal debris detected in the V un-sharp masked image. We suggest that a minor-merger may have driven the nuclear activity in NGC 2110. ",
        "introduction": "The two most important questions concerning active galactic nuclei (AGN) are the fundamental nature of their energy source and the mechanism by which the nuclear activity is fueled. Even though it is widely accepted that the energy source originates in the accretion of mass onto a central super-massive black hole (SMBH), starbursts could also play an important role in the same galaxies (e.g. Gonz\\'alez Delgado 2001; Veilleux 2001, and references therein). In fact, many examples have been reported in which the two phenomena co-exist together. At the high-luminosity regime, ULIRGs are the best examples (Lutz \\& Tacconi 1999). Powerful starbursts are also found in the central $\\sim$ 100 pc of less luminous AGN, such as Seyfert 2 galaxies (Heckman et al 1997; Gonz\\'alez Delgado et al 1998), and even at smaller spatial scales ($\\sim$ few pc), as it is the case of the low luminosity AGN NGC 4303 (Colina et al 2002).  The origin of the gas that fuels the central engine and the mechanism that brings the gas to the very central region are still controversial. Observations of high luminosity AGNs indicate that QSOs and ULIRGs have been gas-rich mergers (Sanders et al 1988; Canalizo \\& Stockton 2001). However, in low luminosity AGNs, the fueling mechanism remains unclear. Early investigations in Seyfert galaxies have also suggested that galaxy-interaction could provide the gas fueling onto the SMBH (e.g. Noguchi 1988; Mihos \\& Hernquist 1996). Early research suggested that there is an excess of companions in Seyfert galaxies (Dahari 1984). However, this result has not been confirmed more recently (De Robertis, Yee \\& Hayhoe 1998; Rafanelli et al 1995). Thus, the galaxy-interaction hypothesis can be ruled out as the main fueling mechanism in Seyfert galaxies.  The non-axisymmetric gravitational potential originated by a large scale bar in the host galaxy has also been suggested as an efficient way to drive the gas to the central region (Shlosman, Begelman \\& Frank 1990; Athanassoula 1992). However, it can be ruled out as a general mechanism because an excess of bars in Seyfert galaxies has not been confirmed observationally (Adams 1977; Moles, M\\'arquez \\& P\\'erez 1995; Hunt et al 1997; Maiolino, Risaliti \\& Salvati 1998). Studies based on near-infrared imaging, known to trace the stellar bars better than the optical observations, arrive also at contradictory results. Based on HST (NICMOS) observations Mulchaey \\& Regan (1997), Regan \\& Mulchaey (1999) and Martini \\& Pogge (1999) find no evidence for an excess of stellar bars in Seyferts. On the opposite, Knapen, Shlosman, \\& Peletier (2000) found an excess in the sample of CfA Seyfert galaxies with respect to a control sample. However, this excess of bars in CfA Seyfert galaxies disappears when galaxies that are in interaction are excluded from the statistics (M\\'arquez et al 2000).  Minor merger between a  disk galaxy and a nucleated less massive galaxy is also proposed as an alternative fueling mechanism (Garc\\'\\i a-Lorenzo et al 1997; de Robertis et al 1998; Taniguchi 1999). In this context, it is interesting to note that some kinematic peculiarities found in NGC 1068 have been interpreted by Garc\\'\\i a-Lorenzo et al (1997 and 1999) as the result of a minor merger event. Since most galaxies have satellites (Zaritsky et al 1997), it is likely that minor mergers drive the fueling mechanism for most of the isolated Seyfert galaxies. However, there is not yet observational proof of this scenario.   NGC 2110 is a nearby ($\\sim$ 31 Mpc), early type S0, isolated galaxy, showing a Seyfert 2 nucleus (McClintok et al. 1979; Shuder 1980). This galaxy has been extensively studied at radio, NIR, optical, UV and X-ray wavelengths. These studies have revealed a number of interesting properties in this object, making it a good target to investigate the minor-merger hypothesis as the mechanism that drives the gas to the center and feeds the AGN in NGC 2110. Its main properties can be summarized as follow:  \\begin{itemize}   \\item{It is classified as a narrow-line X-ray galaxy (Bradt et al 1978) due to the strong X-ray emission, having a hard X-ray luminosity (2-10 KeV) comparable to those of Seyfert 1 galaxies (Weaver et al 1995). Soft X-ray emission is also detected at 4 arcsec North of the nucleus (Weaver et al 1995).}  \\item{ VLA radio continuum maps reveal a symmetrical jet-like radio emission  extended 4 arcsec along the North-South direction (Ulvestad \\& Wilson 1983; Nagar et al 1999). More recent VLBA observations at 3.6 cm also show a compact ($\\leq$ 1 pc) radio core aligned with the 400 pc scale radio jet (Mundell et al 2000).}  \\item{ Ground-based narrow band ([OIII] and H$\\alpha$+[NII]) images reveal extended emission up to 10 arcsec showing an S-shape morphology in the North-South direction (Wilson, Baldwin and Ulvestad 1985; Pogge 1989). HST pre-Costar [OIII] and H$\\alpha$+[NII] images confirm the alignment of the optical and radio jet, but only a loose anti-correlation between the brightest optical and radio knots was found (Mulchaey et al 1994). Ferruit et al (1999) have obtained HST+FOC UV, and optical spectra of the jet to investigate the mechanism that ionizes the extended gas. They suggest that the ionized high excitation gas is produced by photoionization by fast-shock waves probably generated by the propagation of the radio ejecta into the ISM of the host galaxy. Shocks, driven by the radio jet, could be also responsible for the strong [FeII] $\\lambda$1.2567 $\\micron$ nebular emission detected in its nuclear and circumnuclear region (Storchi-Bergmann et al 1999; Knop et al 2001).}  \\item{ From the morphological point of view, the optical continuum (Malkan et al 1998) and NICMOS/WFPC color maps (Quillen et al 1999) show clear asymmetries in the central region. Dust lanes at 1 and 4 arcsec to the west of the nucleus are detected.}  \\item{The nuclear stellar population was studied by Gonz\\'alez Delgado, Heckman, \\& Leitherer (2001) to look for signatures of starbursts. Its optical continuum (stellar features and spectral energy distribution) is well explained by an old stellar population with a small contribution from a power-law. Thus, no obvious signs of a young (age $\\leq$ 1 Gyr) stellar population were found. }  \\item{The kinematics of the circumnuclear ionized gas region was studied by Wilson and Baldwin (1985) and Wilson et al (1985). They reported a rotation pattern remarkably similar to that expected for spiral galaxies. They also found that the rotation center was not coincident with the optical (and radio) nucleus, but was located 1.7 arcsec (220 pc) to the south of the nucleus. They also reported the presence of asymmetric profiles in the [OIII] nebular lines, south and east of the nucleus.}  \\item{HI absorption has been detected in NGC 2110 by Gallimore et al (1999) located south of the maximum of the radio continuum source. This absorption has a velocity of 2626 km/s; red-shifted  by about 290 km/s with respect to the systemic velocity, 2335 km/s, calculated by Strauss et al (1992) using the nuclear optical emission lines.}  \\end{itemize}    Here, we present new Integral Field Spectroscopic (IFS) data at optical wavelengths of the circumnuclear region of NGC 2110. This technique is ideal for studying an asymmetric and complex object like NGC 2110 since it provides 2D spectral information of relevant spectral features (i.e., velocity fields, velocity dispersion maps, line intensity maps, continuum maps, etc) simultaneously and, therefore, without the uncertainties associated with centering, reference systems, etc.  In addition to these IFS data, we also present new longslit spectroscopy to study the stellar kinematics. These data, together with  some WFPC2 images retrieved from the HST archive, allow us to investigate the kinematic peculiarities found in NGC 2110 and their possible connection with the fueling mechanism of the AGN in NGC 2110. The paper is organized as follows: In Section 2 we present the observations and data reduction. Section 3 presents the analysis of the stellar and ionized gas distribution, and section 4 the gas and stellar kinematics. We discuss our results and their implications in section 5. The conclusions and summary are in section 6. ",
        "conclusions": "We have obtained Integral field optical spectroscopy in the range 3700-7000 \\AA, and long-slit optical spectroscopy (3400-5500 \\AA\\ and 6600-9100 \\AA) along P.A= 6$^{\\circ}$ of the Seyfert 2 galaxy NGC 2110. These observations allow us to get the distributions of the emission lines (H$\\alpha$, H$\\beta$, [OIII], [OI], [NII], and [SII]), the ionized gas velocity field, the velocity dispersion map, and the stellar velocity curve along P.A=6$^{\\circ}$. These data combined with HST+WFPC2 (I, F606W and H$\\alpha$) have reported the following results:  \\begin{itemize}  \\item{At large scales, the morphology of NGC 2110 is not perturbed. The red continuum is well fitted by ellipses that have their major axis oriented at P.A=163.5$^{\\circ}$, and an inclination of  42$^{\\circ}$. The centers of these ellipses are located at the maximum of the continuum within 0.13 arcsec in RA and 0.07 arcsec in Dec. This result indicates that the photometric center is at the nucleus of the galaxy. }  \\item{V and I un-sharp masked images have revealed dust lanes that spiral toward the nucleus and they may represent the gas flows that feed the AGN. The V (F606W) un-sharp masked image shows also radial dust lanes and emission arcs at 5 and 8 arcsec at the south and north-east, similar to the tidal debris predicted by numerical simulations of the merger of a small galaxy with a galaxy disk.}  \\item{The ionized gas is extended in the North-South direction, such as the linear radio jet. The morphology of [OI], [NII] and [SII] is similar to [OIII], that is more extended to the north than to the south. This distribution suggests that the gas at the north may be photoionized by the AGN. At larger scales, H$\\alpha$ and H$\\beta$ images show an S-shape that is similarly extended at the north and at the south.}  \\item{The ionized velocity field shows clearly a north-south asymmetry, with a larger velocity gradient in the south than in the north. This asymmetry is seen also along the kinematic major axis. This result suggests that the ionized gas is not normally rotating gas in the disk of the galaxy. The ionized [NII] nuclear gas velocity is 2315 km/s.}  \\item{The ionized gas velocity dispersion map shows a bar like morphology oriented near the photometric minor axis. The maximum of the velocity dispersion is also coincident with the maximum of the optical continuum. Thus, the kinematic center is coincident with the nucleus of NGC 2110. This alignment between the maximum of the emission line widths and the galaxy minor axis may be due to a wind driven by the AGN.}  \\item{The stellar velocity curve inferred from the CaII triplet observed along P.A=6$^{\\circ}$ is symmetric with respect to the nucleus that has a systemic velocity of 2380 km/s. The amplitud of the stellar velocity curve is lower than the amplitud of the ionized gas velocity curve.}  \\item{The nuclear ionized gas of NGC 2110 is blueshifted with respect to the stellar systemic velocity. Two velocity components that are shifted $\\pm125$ km/s from the stellar rotation curve have been detected at the central 2 arcsec. These components indicate that the NLR is out-flowing, possibly driven by the radio jet.}  \\item{At larger spatial scales (up to 5 arcsec),  the ionized gas on the south of NGC 2110 is red-shifted by about 100 km/s with respect to the stellar velocity field, and 240 km/s with respect to the nuclear stellar systemic velocity. This excess of the velocity with respect to the rotation indicates that the gas in the south of NGC 2110 is not following the pure gravitational potential of the galaxy. It is coincident with the HI red-shifted (240 km/s) absorption extended to the south of the nucleus detected by Gallimore et al (1999). }  \\end{itemize}  We conclude that the ionized gas on the south is kinematically perturbed. This conclusion together with the morphological peculiarities detected (arcs, radial dust lanes) suggests that the nuclear activity could be the result of a minor-merger. The collision with a nucleated small satellite (or with an intergalactic HI cloud) with a very high inclined orbit that impacts close to the nucleus may be responsible of the red-shifted velocity of the ionized gas in the south of the galaxy and the arcs and radial dust lanes observed in the central host disk of the galaxy. This interpretation of the data is of relevance in relation with the fueling mechanism and the triggering of nuclear activity in isolated non-barred Seyfert galaxies.     {\\bf Acknowledgments}  We are grateful to Jack Gallimore, the referee, for his suggestions and corrections that helped to inprove the paper. We also thanks to Luis Cuesta for making his {\\it inter2r} program available that allows to built the two-dimensional images, Jaime Perea for SIPL, and Luis Colina for his comments from a thorough reading of the paper. This work was supported by the Spanish DGICYT project AYA2001-3939-C03-01.  The 4.2m William Herschel Telescope is operated by the Isaac Newton Group at the Observatorio de Roque de los Muchachos of the Instituto de Astrof\\'\\i sica de Canarias.  Based observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at ST-ECF.  \\clearpage"
    },
    "0207/astro-ph0207105_arXiv.txt": {
        "abstract": "We use the Gaussian-fit results of Paper I to investigate the properties of interstellar HI in the Solar neighborhood. The Warm and Cold Neutral Media (WNM and CNM) are physically distinct components. The CNM spin temperature histogram peaks at about 40 K; its median, weighted by column density, is 70 K.  About $60\\%$ of all HI is WNM; there is no discernable change in this fraction at $z=0$. At $z=0$, we derive a volume filling fraction of about 0.50 for the WNM; this value is very rough. The upper-limit WNM temperatures determined from line width range upward from $\\sim 500$ K; a minimum of about $48\\%$ of the WNM lies in the thermally unstable region 500 to 5000 K. The WNM is a prominent constituent of the interstellar medium and its properties depend on many factors, requiring global models that include all relevant energy sources, of which there are many.  We use Principal Components Analysis, together with a form of least squares fitting that accounts for errors in both the independent and dependent parameters, to discuss the relationships among the four CNM Gaussian parameters. The spin temperature $T_s$ and column density $N(HI)$ are, approximately, the two most important eigenvectors; as such, they are sufficient, convenient, and physically meaningful primary parameters for describing CNM clouds.  The Mach number of internal macroscopic motions for CNM clouds is typically about 3 so that they are strongly supersonic, but there are wide variations. We discuss the historical $\\tau_0$-$T_s$ relationship in some detail and show that it has little physical meaning.  We discuss CNM morphology using the CNM pressure known from UV stellar absorption lines. Knowing the pressure allows us to show that CNM structures cannot be isotropic but instead are sheetlike, with length-to-thickness aspect ratios ranging up to about 280. We present large-scale maps of two regions where CNM lies in very large ``blobby sheets''.  We test the McKee/Ostriker model of the interstellar medium by explicitly modeling our data with CNM cores contained in WNM envelopes. This modeling scheme works quite well for many sources and also predicts the WNM filling factor reasonably well. However, it has several deficiencies. ",
        "introduction": "\\label{introduction}  This paper discusses  the astronomically oriented results of a new  Arecibo\\footnote{The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under a cooperative agreement with the National Science Foundation.} 21-cm absorption line survey; it is the comprehensive version of the preliminary report by Heiles (2001b). Paper I (Heiles and Troland 2002) discusses the observational and data reduction techniques.  We took great care in accounting for instrumental gain fluctuations and angular structure of HI so that we could derive accurate opacity and expected emission profiles, including realistic uncertainties.  (An expected profile is the emission profile towards the source that would be observed if the source flux were zero).  The opacity profiles come from the Cold Neutral Medium (CNM) and are characterized by distinct peaks; we decomposed them into Gaussian components.  The expected profiles are produced by both the Warm Neutral Medium (WNM) and the CNM. We fit them using a simple but physically correct radiative transfer equation that includes both the emission and absorption of the CNM and, in addition, one or a few independent Gaussians for the WNM emission. We discussed the fitting process and its uncertainties in detail, and presented many examples of the technique.  We derived spin temperatures for the CNM using the opacity and expected profiles.  We derived upper-limit temperatures for the CNM using the line widths.  We presented all results in tabular, graphical, and electronic form.  Table \\ref{sourcelist} summarizes the sources observed and the column densities of CNM and WNM. Here, by ``WNM'', we mean Gaussian components detected only in emission, and by ``CNM'' we mean Gaussians that were detected in absorption.  Paper I presents the full table of Gaussian component properties. We have a total of 79 sources, 202 CNM components, and 172 WNM components. 13 sources have $|b|<10^\\circ$, and we exclude these from some of our discussion below because their profiles are complicated or the WNM linewidths might be significantly broadened by Galactic rotation.  \\S \\ref{tempdist} shows that the division between WNM and CNM is not only observational, but also physical; \\S \\ref{cnmwnmsummary} summarizes the statistics on CNM/WNM column densities for the Gaussians. \\S \\ref{coldenstatistics} presents column density statistics for the lines of sight for the CNM and WNM. \\S \\ref{volfill} discusses the volume filling fraction of the WNM, both at high and low $z$.  The next few sections discuss the basic statistical properties of the Gaussian components. \\S \\ref{vlsrstatistics} presents statistics on $V_{LSR}$. \\S \\ref{correlations} presents correlations among the four parameters that describe the CNM components. The reader interested in these correlations should consult the two subsequent sections: \\S \\ref{pequality} shows that inadequate angular resolution might affect these correlations, and \\S \\ref{againstraisin} shows that CNM features are sheetlike and not isotropic with the consequence that angular resolution effects are far less important than found in \\S \\ref{pequality}.  \\S \\ref{momodel} re-reduces all the data of Paper I in terms of the McKee \\& Ostriker (1977) (MO) model, with each CNM component surrounded by an independent WNM component; it is gratifyingly successful for most sources but some MO predictions are not quantitatively fulfilled. \\S \\ref{descmodel} presents two descriptive models; the second, the clumpy sheet model for the CNM, applies to our data.  \\S \\ref{summary} is a summary, and \\S \\ref{commentary} is a commentary on the importance of the WNM for understanding not only the ISM but also its multiplicity of energy sources and the Universe at large. ",
        "conclusions": "For sources at $|b|>10^\\circ$, the global ratio of WNM to total HI column density is $\\langle R(HI)_{WNM} \\rangle =0.61$.  Mass is equivalent to column density if the distances are the same.  The WNM is systematically more distant than the CNM because it has a larger scale height (Kulkarni \\& Heiles 1987), so this is a lower limit for the mass fraction.  The $N(HI)$ fraction of WNM having $T_{kmax}$ in the unstable region 500 to 5000 K is $0.48$; the true fraction of gas in this unstable regime might be higher because $T_{kmax}$ is an upper limit on temperature derived from the linewidth.  It's conceivable, but unlikely in our opinion, that much of this gas has temperature $T_k<500$ K.  The $N(HI)$ fraction of CNM having $T_s$ in the range 25 to 70 K (the main peak in the histogram) is $0.46$.  At low latitudes, $|b| \\lesssim 10^\\circ$, the line of sight doesn't leave the HI layer for nearby gas. We can use low-latitude sources as a test to determine whether the fraction of WNM gas $R(HI)_{WNM}$ decreases at lower $|z|$ where the pressure is higher, as is theoretically predicted. We have 8 sources with reasonably accurate Gaussian fits (and 5 with unusable fits; Table \\ref{sourcelist}). These 8 sources have $\\langle R(HI)_{WNM}\\rangle = (0.67 \\pm 0.08)$. This is indistinguishable from the $|b|>10^\\circ$ mean value $\\langle R(HI)_{WNM}\\rangle = 0.61$. Thus, there is no evidence for the predicted decrease in $R(HI)_{WNM}$. However, we stress that our low-latitude results are generally less accurate than the others because it is more difficult to obtain accurate expected profiles and to perform Gaussian fits. Accurate results for low latitudes probably requires high-sensitivity interferometric observations.  \\subsection{Comparison of CNM temperatures with other results}  \\label{comparison}  \\subsubsection{ Previous CNM temperatures from the 21-cm line}  Our spin temperatures are colder than previously obtained ones. Histograms of CNM temperatures have been given by Dickey, Salpeter, \\& Terzian (1978, DST), Payne, Salpeter, \\& Terzian (1983, PST), and Mebold \\ et al 1982, among others. They find broader histograms than ours with temperatures extending to much higher values and median values in the neighborhood of 80 K; for example, Mebold et al find a median (by components) of 86 K. Our histogram is narrower and peaked near 40 K (Figure \\ref{histoplot4new}) and our median (by components) is 48 K. In contrast, our median (weighted by column density) is 70 K. When quoting medians, it is important to distinguish between the component median and the column-density median.  Our lower temperatures do not arise because the older data were incorrect (although some were); it is because the analyses were incorrect. In contrast to the previous treatments, our Gaussian technique, which is thoroughly discussed in Paper I \\S 4 and \\S 5, properly accounts for the two-phase medium and the associated radiative transfer. Recent measurements of temperatures in the Magellanic Clouds (Mebold et al 1997; Marx-Zimmer et al 2000; Dickey et al 2000) use the slope technique, which also properly treats radiative transfer for simple profiles (Paper I, \\S 4, \\S 6); they find smaller temperatures, consistent with ours, and show that the older incorrect technique yields incorrect higher temperatures.  \\subsubsection{ Temperatures from H$_2$}  Temperatures are also derived from the ratio of populations in the two lowest rotational states of H$_2$.  Unfortunately, these are not directly comparable to our CNM temperatures, for two reasons. First, the H$_2$ lines of sight are chosen to maximize column density; in contrast, ours are random with respect to column density. Second, the H$_2$ lines are saturated, which means that the derived temperatures are a weighted average over all velocity components and all the gas, both CNM and WNM; one cannot know which phase dominates the results because the fractional H$_2$ abundances in the two phases are unknown. Because the H$_2$ measurements refer to all gas, a median derived therefrom is more akin to a column-density median than a component median.  Recent FUSE measurements (Shull et al 2000) confirm the large survey of Savage et al (1977), who found the range of temperatures to be $T_{H_2} = 77 \\pm 17$ (rms) K.   This is comparable to our component median for the CNM. However, because the H$_2$ sample is biased to large column density lines of sight, the results are not directly comparable. We further explore the comparison by considering four of our sources that are fairly close to stars in three regions studied by Savage et al. This by no means guarantees that the physical regions sampled are identical, but one hopes that the lines of sight are physically similar.Table \\ref{h2tbl} shows radio sources and stars in these three areas; in each area the radio and optical positions are close, within a few degrees.  The first two regions have high $N(HI)$ and are cold, with CNM temperatures lying near the peak of our histogram; the H$_2$ temperatures are higher than the HI temperatures.  We detected the 21-cm line in absorption in the third region but we would not classify the $510$ K gas as CNM; the H$_2$ temperature of 377 K is smaller than the HI temperature, although realistic uncertainties may mean that the results are consistent.  The upshot is that the H$_2$ temperatures do not agree with the HI CNM temperatures: This conclusion needs confirmation via observations of HI and H$_2$ absorption along identical lines of sight. Such observations require a background source such as 3C273 with significant radio and UV emission.  \\label{summary}  Paper I discusses the observational and data reduction techniques. In particular, it devotes considerable attention to the Gaussian fitting process, which is subjective and nonunique. Concerned readers should see \\S 5 of that paper.  The present paper treats the astronomically oriented implications of the Gaussian components from Paper I and includes the following topics: \\begin{enumerate}  \\item \\S \\ref{tempdist} discusses the statistics of the Gaussian components. It shows that the CNM and WNM are not only observationally distinct, but also physically distinct. The median column density per CNM Gaussian component is about $0.5 \\times 10^{20}$ cm$^{-2}$, and per WNM component about $1.3 \\times 10^{20}$ cm$^{-2}$ (Table \\ref{medianmeannh}).  The CNM temperature histogram peaks near $T_s = 40$ K (Figure \\ref{histoplot4new}), about half the temperature obtained by previous workers. Its median by components is 48 K and, weighted for $N(HI)$, 70 K. CNM temperatures range down to $\\sim 15$ K, which can be attained only if grain heating is not operative.  CNM temperatures appear to be smaller than those derived from UV absorption line observations of H$_2$, but the comparison means little because H$_2$ temperatures refer to all velocity components and all phases along the line of sight.  A significant fraction of the WNM, $\\gtrsim 48\\%$, lies in the thermally unstable range $T_k = 500$ to 5000 K.   \\item \\S \\ref{coldenstatistics} summarizes the statistics of WNM and CNM column densities for entire lines of sight instead of individual Gaussian components. There are many lines of sight having no CNM; these form a distinct class and are confined to particular areas of the sky. Column densities depart very markedly from those expected from a plane-parallel distribution. $61\\%$ of the HI we observed is WNM; at $z=0$, it fills $\\sim 50\\%$ of the volume, but this number is {\\it very rough}. In \\S \\ref{volfill} we show that this is in reasonably good agreement with MO, when the WIM-associated HI is included.  Figure \\ref{mapnhrawratio2} shows the factor $R_{raw}$ by which $N(HI)$ calculated from the optically thin approximation (i.e.\\ from the line profile area) underestimates the true $N(HI)$; this can be significant even at high Galactic latitudes.  \\item \\S \\ref{vlsrstatistics} shows that the component velocities that we observe are not significantly affected by Galactic rotation. The column-density weighted rms velocities are about 7 and 11 km s$^{-1}$ for the CNM and WNM Gaussian components, respectively.  \\item \\S \\ref{correlations} uses Principal Components Analysis, together with a form of least squares fitting that accounts for errors in both the independent and dependent parameters, to discuss the relationships among the four CNM Gaussian parameters. The spin temperature $T_s$ and column density $N(HI)$ are, approximately, the two most important eigenvectors; as such, they are convenient, physically meaningful primary parameters for describing CNM clouds.  The Mach number of internal macroscopic motions for CNM clouds is typically $\\sim 3$, but there are wide variations and a weak increase with $T_s$. Most CNM clouds are strongly supersonic. We discuss the historical $\\tau_0$-$T_s$ relationship in some detail and show that it has little physical meaning.  \\item \\S \\ref{pequality} discusses the possible effect of angular resolution on the relationships among observed CNM parameters. These effects are important if CNM clouds are isotropic. However, \\S \\ref{againstraisin} shows that CNM clouds are definitely not isotropic. CNM features are sometimes large sheets with aspect ratios measured in the hundreds. These sheets contain blobs, which themselves are sheetlike but with much smaller aspect ratios.  \\item \\S \\ref{momodel} directly compares our data with the McKee/Ostriker model by re-reducing all Gaussian components in terms of that model, i.e.\\ with each CNM cloud having an associated WNM envelope. This fitting scheme works very well for many sources, but not for all. The MO model greatly underpredicts the WNM abundance and, also, the fraction of WNM that is thermally unstable.  \\item In \\S \\ref{descmodel} we argue that there is so much WNM that CNM clouds probably don't have individual WNM halos, but rather that many CNM clouds exist within a common WNM halo. We discard the raisin pudding model as a commonly envisioned descriptive model and replace it by the blobby sheet model, in which the CNM consists of sheetlike structures with sheetlike blobs or cloudlets embedded within. Each WNM cloud probably contains a few CNM large sheets.  \\item \\S \\ref{againstraisin} uses our knowledge of the CNM pressure to derive the morphological shape of CNM structures: they are sheetlike. In two regions of the sky the CNM is organized into large sheets with length-to-thickness aspect ratios $\\sim 280$ and 70; the latter is permeated by small sheetlike structures.  \\item In the following section we provide comments on the importance of the WNM for understanding not only the ISM but also the full range of its energy sources.  \\end{enumerate}"
    },
    "0207/astro-ph0207333_arXiv.txt": {
        "abstract": "{A significant  fraction of extended radio  sources presents a peculiar  X-shaped radio  morphology: in addition  to the  classical double lobed structure, radio emission is also observed along a second axis of symmetry in the form of  diffuse wings or tails.  We re-examine  the   origin  of  these  extensions   relating  the  radio morphology to the properties  of their host  galaxies.  The orientation of the wings shows  a striking connection with the structure of the  host galaxy as they are  preferentially aligned with its minor axis. Furthermore, wings  are only observed in galaxies of high projected ellipticity. Hydrodynamical simulations of the radio-source evolution show that X-shaped radio-sources naturally form in this geometrical situation: as a jet propagates in a non-spherical gas distribution, the cocoon surrounding the radio-jets expands laterally at a high rate producing wings of radio emission, in a way that is reminiscent of the twin-exhaust model for radio-sources. ",
        "introduction": "\\label{intro}  Extended radio sources have  been historically classified on the basis of their radio morphology, the main division being based on their edge darkened or edge brightened  structure that leads to the identification of the Fanaroff-Riley classes I  and II (Fanaroff \\& Riley 1974).  The characteristic  structure of  FR~II sources  is dominated  by  two hot spots located  at the edges  of the radio  lobes that, in  most cases, show  bridges of  emission linking  the core  to the  hot  spots.  The presence  of significant  distortions  in the  bridges was  recognized since early interferometric  imaging of 3C sources (see  e.g. Leahy \\& Williams 1984). Distortion  in FR~II can be classified  in two general classes: mirror  symmetric ( or  C-shaped) when the bridges  bend away from the galaxy  in the same direction, or  centro symmetric when they bend  in opposite  directions forming  an X-shaped  or  Z-shaped radio source, depending  on the  location of the  point of insertion  of the wings.   In  many  X-shaped  sources  the  radio  emission  along  the secondary axis, although more  diffuse, is still quite well collimated and can be even more extended than the main double lobed structure.  C-shaped  morphologies  are   observed  also  in  FR~I  radio-galaxies although  in these sources  the distortions  affect their  jets rather than their  lobes and they  give rise to  the typical shape  of Narrow Angle  Tails, where  the opposite  jets bend  dramatically  and become almost parallel one  to the other. Another common  morphology for FR~I is that  of centro-symmetric S-shaped sources.   Conversely, there are no examples of X-shapes among FR~I.  There is now a general  agreement that the C-shaped radio sources form when they are  in motion with respect to the  external medium: jets or bridges are bent  by the ram pressure of  the surrounding gas.  Models successfully reproduced the morphology of FR~I Narrow Angle Tails (see e.g. O'Dea \\&  Owen 1986) and the extension of  this scenario to FR~II bridges appears quite natural.  Concerning  the X-  or Z-shaped  sources several  mechanism  have been proposed  for their  origin. Ekers  et al.  (1978) suggested  that the tails of  radio emission  in one  of these sources,  NGC 326,  are the result of the trail caused by  a secular jet precession (see also Rees 1978).  A similar model accounts for the morphology of 4C 32.25 (Klein et al.  1995). In a  similar line, Wirth  et al.  (1982) noted  that a change in the jet direction can be caused by gravitational interaction with a  companion galaxy. Recently Dennett-Thorpe et  al. (2002), from the analysis of spectral variations along the lobes, proposed that the jet reorientation  occurs over short time  scales, a few  Myr, and are possibly associated to instabilities  in the accretion disk that cause a rapid  change in the  jet axis. In  all these models,  the secondary axis of radio emission represents a  relic of the past activity of the radio source.  An  alternative  interpretation was  suggested  by  Leahy \\&  Williams (1984) and  Worrall et  al. (1995).  They  emphasize the  role of  the external  medium in  shaping radio  sources, suggesting  that buoyancy forces can bend the back-flowing  material away from the jet axis into the direction of decreasing external gas pressure.  In this Paper we present a different scenario based on evidence for a strong  connection between the  properties of  the radio  emission and those of  the host galaxy, a  comparison that has been overlooked in the  past but  that provides  crucial new  insights on  the  origin of X-shaped radio-sources.   In fact, in Sect. \\ref{host}  we compare the radio and  host galaxy orientation  of X-shaped radio  sources showing that there is a close alignment between the radio wings and the galaxy minor axis.  We  also show that X-shaped sources  occur exclusively in galaxies  of high  ellipticity.  Numerical  simulations,  presented in Sect.  \\ref{simulations}, provide  support to  the idea  that X-shaped radio-sources form  naturally in this  geometrical situation. Finally, in  Sect. \\ref{discussion}  we  discuss the  implications of  our results that are summarized in Sect. \\ref{summary}.   \\input tab1.tex ",
        "conclusions": "\\label{summary}  We presented evidence that the orientation of the secondary axis of radio emission in X-shaped radio-sources is closely linked to the orientation of the host galaxy. More specifically, considering the effects of projection, the radio wings appear to be parallel to the galaxy minor axis. This relationship suggests a causal connection between the galaxy property and the origin of the X-shaped radio morphology. Such a connection is strengthened noting that X-shaped radio sources are found only in galaxy of high ellipticity. We then  argue that X-shaped  radio-sources naturally form when  a jet propagates  in a  non-spherical  gas distribution. In  this case  the cocoon expansion along the direction of maximum pressure gradient (the galaxy  minor axis) can  occur at  a similar  (or higher)  rate than along the  radio source main  axis. The situation is reminiscent of the twin-exhaust scenario where a bubble of hot and light gas inside an elliptical distribution of density and pressure expands laterally at a high rate producing wings of radio emission. Hydrodynamical simulations  of the radio-source  evolution in this situation support this scenario.  From the point of numerical simulations, a three dimensional analysis is clearly needed to investigate in more detail the radio source behaviour, but it might be envisaged that with more realistic simulations it might be possible to constraints the jet's properties based on the wings formation. In particular it will be important to perform  a thorough exploration of the jet's parameters space. In fact, it is clear that the effects of the asymmetric gas distribution will affect more jets of lower kinetic power: at higher power the crossing time of the galaxy's central regions will be shorter and this will in turn shorten the phase during which the cocoon can expand laterally. We thus expect that only jets in a limited region of the parameters space might produced winged radio-sources.  Finally, a very valuable information will come from mapping of the external gas from high resolution X-ray images such as those that can now be obtained with the new generation of X-ray satellites such as Chandra of Newton-XMM."
    },
    "0207/astro-ph0207619_arXiv.txt": {
        "abstract": "There  are three  important aspects  concerning the  study of  the red giant and in particular of  the asymptotic giant branch (AGB) stars in the Magellanic Clouds. These are: the surface distribution, the luminosity function and the variability. The  spatial distribution of AGB stars is  an efficient tool to study the structure of the galaxies  and their metalicity by analysing  the ratio between carbon--  and  oxygen--rich AGB  stars. The shape of the luminosity function carries informations about  the star formation rate in  the Clouds and it  can  be mathematically  related  to  their  history. Most  AGB  stars vary  their magnitude in a few to several  hundred years time; the one epoch DENIS magnitudes for both Large and Small Magellanic Cloud AGB stars outline the same relations as a function of period. ",
        "introduction": "Studying  the  stellar  content  of  the Magellanic  Clouds  has  some advantages: these  are nearby galaxies  that can be fully resolved in stars, they are relatively un--obscured and all their stars are at about the  same distance. Before the publication  of the catalogues and follow--up studies  from large scale surveys what  was known about red  giants  in  the  Clouds  were  spatially  and  magnitude  limited informations,  often  focalised   on  particular  objects  (i.e.  Mira variables).   The past  few  years  have seen  the  release of the  DENIS catalogue   towards  the   Magellanic   Clouds  (DCMC   --  Cioni   et al.~\\cite{cioni1}) which provides photometry in the broad $I$, $J$ and $K_S$ bands.  These filters are particularly suitable for the study of the red  giant population (Cioni  et al.~\\cite{cioni2}, \\cite{cioni3}, \\cite{vdm1}).   The  release  of  the  2MASS  catalogue  (Nikolaev  \\& Weinberg \\cite{2mass}) provides photometry in the  $J$, $H$ and $K_S$  broad bands and  reaches slightly  fainter magnitudes  than the DENIS measurements. The three micro-lensing surveys OGLE, MACHO \\& EROS provide light--curves over several years in two very broad blue and red filters that allow  to characterise the complex variability of the red giants,  as discussed also in the contribution by Wood (this proceeding).  The OGLE  and MACHO  datasets are  at present  publically  available.  The combination of the near--infrared  observations and light--curves is a key  to understand  the  evolution and  properties of  long--period variables.   In    addition   Zaritsky   (\\cite{zari}) published his $UBVI$ survey of the SMC providing also an extinction map of the Cloud. Finally, Massey (\\cite{massey}) presented less sensitive $UBVR$ measurements in both Clouds to characterise the more massive and brighter stellar populations. \\begin{figure} \\begin{minipage}[t]{6cm} \\resizebox{6cm}{!}{\\includegraphics{iijlmc.ps}} \\end{minipage}\\   \\ \\begin{minipage}[b]{6cm} \\resizebox{6cm}{!}{\\includegraphics{kjklmc.ps}} \\end{minipage} \\caption{Colour--magnitude diagrams of DCMC sources detected simultaneously in three DENIS wave bands. Left: ($I-J$,$I$). The TRGB is at $I=14.54$. Right: ($J-K_S$, $K_S$), the TRGB is at $K_S=11.98$, O--rich AGB stars have $(J-K_S)=1.2$ and C--rich AGB stars have $(J-K_S)>1.4$.} \\label{cmds} \\end{figure} In the following sections the  discussion is concentrated on the DENIS measurements described  in my Ph.D. thesis (\\cite{cionit}),  I am also using EROS  and MACHO light--curves, and ISO  (LW2, LW10) measurements covering small fields in the Clouds (Loup et al.~{\\it in preparation}). ",
        "conclusions": "We  have seen that different groups  of red giants can       be      distinguished       in       the      near--infrared colour--magnitude--diagrams, we have also  seen that the AGB stars are very good  indicators of the extended  structure of a  galaxy and that from their luminosity function we  can obtain indications on the stars formation history  and the metalicity,  the latter from the  ratio of C-- over  O--rich AGB  stars. Finally, these  long period  variables follow the same period--luminosity relations in both Clouds."
    },
    "0207/astro-ph0207455_arXiv.txt": {
        "abstract": "We present the results of completeness simulations for the detection of point sources as well as redshifted elliptical and spiral galaxies in the $K'$-band images of the Munich Near-Infrared Cluster Survey (MUNICS).  The main focus of this work is to quantify the selection effects introduced by threshold-based object detection algorithms used in deep imaging surveys. Therefore, we simulate objects obeying the well-known scaling relations between effective radius and central surface brightness, both for de Vaucouleurs and exponential profiles. The results of these simulations, while presented for the MUNICS project, are applicable in a much wider context to deep optical and near-infrared selected samples. We investigate the detection probability as well as the reliability for recovering the true total magnitude with Kron-like (adaptive) aperture photometry.  The results are compared to the predictions of the visibility theory of Disney and Phillipps in terms of the detection rate and the lost-light fraction.  Additionally, the effects attributable to seeing are explored. The results show a bias against detecting high-redshifted massive elliptical galaxies in comparison to disk galaxies with exponential profiles, and that the measurements of the total magnitudes for intrinsically bright elliptical galaxies are systematically too faint. Disk galaxies, in contrast, show no significant offset in the magnitude measurement of luminous objects.  Finally we present an analytic formula to predict the completeness of point-sources using only basic image parameters. ",
        "introduction": "\\label{s:introduction}  The luminosity function as well as the mass function, and to a lesser degree the number counts of galaxies provide an important observational toolset for understanding the evolution of galaxies. Due to their statistical nature, these methods rely on the understanding of sample selection effects, i.e.\\ the knowledge of what fraction of the true number of galaxies, as a function of their intrinsic properties, is actually present in the sample. In order to extract this information from the data, two different approaches are commonly used in the literature, creation of artificial objects or modification of observed objects.  The synthetic objects are, as described by e.g. \\citet{Martini2001}, usually point-sources or objects with galactic profiles, that are inserted into the observed images. Then the fraction of objects recovered by the applied detection algorithm as a function of, e.g, the assigned apparent magnitude is computed. This approach exhibits two major drawbacks: The galaxies are created using a discrete set of half-light radii and a continuous range of magnitudes, thus ignoring known relations between surface-brightness and effective radius, like the fundamental plane \\citep{BBF92,GFSBP2000}. The results show the completeness limits for distinct types of objects, but as the actually observed ratio of these types is unknown, predictions about the total completeness of the data cannot be made.  In the second case, as conducted by e.g.\\citet{BLK1998}, \\citet{ADCZFG99} or \\citet{Saracco99}, images of the observed objects are dimmed or brightened to produce artificial objects. The advantage of this approach is, that it preserves the observed mix of galaxies, assuming that it remains unchanged with time, yet the resulting completeness fractions for extended objects are questionable, as long as the resolution of images is seeing-limited.  When the bright extended objects are dimmed, they are still profile dominated, whereas in reality a faint object would be seeing dominated. The same holds true for artificially brightened faint sources. The resulting objects would still be seeing dominated.  Furthermore, the observed size of a local object of given magnitude is different from the size of a distant object of the same brightness.  Consequently this method only yields information about the probability to detect an object if it was fainter, but not about the probability to detect an intrinsically faint object.  Another drawback of both approaches is that the obtained results can only be used to define the completeness of the survey in apparent magnitudes. The magnitudes and radii of the artificially constructed objects all are physically plausible, but the objects do not occupy one plane in the $\\langle \\mu_e \\rangle - r_e - z$ parameter space. Therefore the results derived using galactic profiles are not usable to correct absolute magnitudes, as effects of the completeness of different galaxy types at given redshifts are unknown.  In deep extra-galactic surveys, the observed galaxies span a wide range in intrinsic profile shape, intrinsic brightness, and intrinsic size. The apparent quantities vary with cosmological distance, such that the fraction of galaxies visible is also a function of redshift. The need to simulate objects obeying the known scaling relations for galaxies was pointed out several times in the literature.  Using profiles obeying the magnitude-radius relations,~\\citet{Yoshii1993} predicted the number of objects lost in number count analysis. He finds a strong dependence of the detection rate on the applied detection criteria and magnitude measurement algorithm, leading to a larger number of undetected faint galaxies at high redshifts.  \\cite{Dalcanton1998} analysed the biases of the luminosity function introduced by the cosmological effects of size variation with redshift, and cosmological dimming for galactic profiles in dependence of size and magnitude, taking into account effects of seeing. Starting from the deficiency that magnitudes are usually measured as some sort of isophotal magnitude, that is directly influenced by the above mentioned effects, she finds the possibility of a severe underestimation of the true luminosity function introduced by the fact that an object and distance dependent part of the light is lost outside the limiting isophote.  The simulations presented in this paper were performed by adding artificially created objects into the $K'$-band images of the Munich Near-Infrared Cluster Survey (MUNICS; \\citealp{MUNICS1}). We analysed the detection probability and the reliability for recovering the assigned magnitude for point-like sources (Moffat profiles), elliptical galaxies (de Vaucouleurs profiles), and spiral galaxies (exponential profiles).  The radii and magnitudes of the simulated galaxies were distributed according to the projected fundamental-plane relation~\\citep{BBF92} for ellipticals and a Freeman law type relation~\\citep{Freeman1970} constructed from observed data for spirals.  The galaxies were simulated at five distinct redshifts between $z=0.5$ and $z=1.5$, taking into account size variation with redshift as well as cosmological dimming and K-correction. A flat universe with $\\Omega_M = 0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=65$~km~s$^{-1}$~Mpc$^{-1}$ was assumed.  The results of the completeness simulations for the different fields of the MUNICS survey presented here, illustrate the principal limitations of imaging surveys. A full correction of the incompleteness would only be possible if the type mix of galaxies was known.  The presentation of the results here is limited to the $K'$-band data, but as the selection biases are caused by the physical nature of the objects, they are applicable in a much wider context to deep extragalactic surveys spanning the optical and near-infrared wavelength regimes. The conclusions drawn here result from the use of a threshold-based detection algorithm.  Accordingly they will hold true for other similarly created datasets as well.  Finally, a reliable and handy analytic formula to estimate the completeness limit of a survey for point-like sources is presented. In Section~\\ref{s:implementation} the implementation of the simulations and the generation of the artificial objects is described. Section~\\ref{s:discussion} presents the results of the simulations, and a discussion of these. To compare the results with analytic predictions, we present the results of an analysis using the visibility theory devised by \\citet{DP1983} in Section~\\ref{s:visibility} for the detection probability and in Section~\\ref{s:lostlight} for the lost-light fraction. In Section~\\ref{s:seeing} we discuss the effects of seeing in the above mentioned analysis. Finally, in Section~\\ref{sec:values} we present the results of the completeness simulations for the MUNICS fields.  \\subsection{The Munich Near-Infrared Cluster Survey (MUNICS)}  The Munich Near-Infrared Cluster Survey (MUNICS) is a wide-field medium-deep survey in the near-infrared $K'$ and $J$ pass-bands (Drory et al.\\ 2001; MUNICS1 hereafter). The survey consists of a $K'$-selected catalogue down to $K' \\le 19.5$ covering an area of 1 square degree. Additionally, 0.35 square degrees have been observed in $I$, $R$, and $V$. The layout of the survey, the observations and data reduction are described in \\citet{MUNICS1}. ",
        "conclusions": "\\label{s:discussion}  \\begin{figure} \\centering \\includegraphics[width=8.4cm]{fig2} \\caption{Probability $P_D(m)$ to re-detect an artificially created point-source (upper panel), an extended object with a de Vaucouleurs profile (middle panel) or exponential profile (lower panel) as a function of the objects' $K'$-band input magnitude $m_{K'}$. The different line types in the middle and lower panel show the completeness for the five analysed redshifts: $z=0.5$ (solid), $0.75$ (dash), $1.0$ (dot), $1.25$ (dashdot) and $1.5$ (dashdotdot).} \\label{f:compl_results} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8.4cm]{fig3} \\caption{Mean magnitude difference $\\delta m = m_{K'} - m_{meas}$ between the assigned magnitude $m_{K'}$ of an object and the Kron-magnitude measured by the photometry software $m_{meas}$ in bins of $0.25$ mag as a function of the input $K'$-band magnitude $m_{K'}$. The top panel shows the results for point-sources, the middle panel for de Vaucouleurs profiles and the lower panel for exponential profiles. The errorbars show the standard deviation of $\\delta m$} \\label{f:compl_magmag} \\end{figure}  The results for one of the MUNICS fields with a seeing of approximately one arcsec FWHM are shown in Fig.~\\ref{f:compl_results}. The upper panel shows the results for the point-sources, the middle and the lower panel for the de Vaucouleurs and exponential profiles, respectively, for the five redshifts simulated. An interesting effect seen in the figure is that, for higher redshifts $z \\ga 1$, the detection probability does not reach unity even for the brightest objects, forming a plateau at some lower value. This effect is mainly caused by the cosmological surface brightness dimming and discussed in detail in Section~\\ref{s:discussion} below. Fig.~\\ref{f:compl_results} shows that the completeness for point-like sources provides a rough but reasonable approximation for the completeness of the analysed extended objects up to a redshift of $z=1.0$.  For each detected object, the difference between the input magnitude and the measured Kron-magnitude was computed. Fig.~\\ref{f:compl_magmag} shows the mean magnitude differences for the analysed profiles, averaged in bins of 0.25~mag and with the standard deviation of the measurements in the bin indicated as errorbars.  The object recovery probabilities for the high-redshift de Vaucouleurs profiles, and to a lesser degree for the exponential profiles exhibit a significant detection bias compared to lower redshift objects, as shown in Fig.~\\ref{f:compl_results}. The fact that the objects at these redshifts never reach a detection probability of one is caused by the distribution of their physical parameters along the fundamental-plane relation -- in case of the ellipticals -- and according to the Freeman law -- in case of the disks.  In both cases, as the object's size increases its radius grows and therefore the average surface brightness decreases. As a result, even the brightest objects of the distribution fail to produce a large enough area above the threshold in surface brightness to be detected with high probability in the presence of noise. In Section~\\ref{s:visibility} these results are compared with the theoretical predictions of the visibility theory, which will confirm the above statement.  The detection deficiencies at low redshifts found for the exponential profiles are caused by the flat light distribution of the objects, resulting in a low central surface brightness even for the luminous objects. The image noise can then scatter these objects below the detection threshold. This result is in general agreement with the predictions of~\\citet{GBS1995}. In their paper they re-derive the formalism devised by~\\citet{DP1983} using two parameters for the disk galaxies. They predict that disk galaxies would suffer from detection biases introduced by their low central surface brightness.  \\citet{Martini2001} performed simulations similar to the ones presented here, analysing the detection probabilities of Moffat, de Vaucouleurs and exponential profiles. The galactic profiles were created for a discrete set of half-light radii and a continuous range of apparent magnitudes. The simulations show a decrease of the limiting magnitude with increasing half-light radius. This effect can be found in the simulations presented here, by comparing the completeness magnitude at different redshifts. With increasing redshift, the sampled range in absolute magnitudes is shifted toward fainter apparent magnitudes, resulting in an increase of effective radius at given apparent magnitude. The detection bias against high redshifted elliptical galaxies we showed cannot be found by the simulations performed by Martini, as these effects are only visible when distributing the objects along a fundamental-plane relation.  The bias first shown by~\\citet{DP1983}, that with increased distance more low-luminosity system would be lost from the survey, is reproduced by our simulations as well. With increasing redshift, the magnitude of the completeness limit stays approximately constant, or gets brighter, while the range of sampled absolute magnitudes is shifted to fainter apparent magnitudes. This means, that with increasing redshift, the observable range in absolute magnitudes moves toward higher luminosities, the low-luminosity systems become unobservable.  \\begin{figure} \\centering \\includegraphics[width=8.4cm]{fig4} \\caption{Intensity integrated over image pixels for a de Vaucouleurs profile (solid line) and an exponential profile (dashed line) at redshift $z=0.5$ with similar total magnitude and effective radius. The upper left panel shows the case in the absence of seeing, the upper right one, with a seeing of $0.5$ arcsec and the lower left with a seeing of $1.6$ arcsec. Plotted is the logarithm of the intensity against the radius relative to the effective radius.} \\label{f:prof} \\end{figure}  The deviations of the measured magnitude from the true magnitude for the de Vaucouleurs profiles as seen in Fig.~\\ref{f:compl_magmag}, can be explained as resulting from the estimate of the Kron-like aperture radius under the conditions of a surface brightness limited detection procedure and the involved intrinsic brightness profile of the objects. The upper left panel of Fig.~\\ref{f:prof} shows the intensity of a de Vaucouleurs and an exponential profile with similar total magnitude and effective radius, as a function of the radius in units of the effective radius in absence of seeing. Assuming that both profiles are detected at a similar intensity level, the Kron-radius of the de Vaucouleurs profile would be smaller compared to the exponential profile, even though the effective radius of both objects is the same. The measured light within the underestimated radius results then in a too faint object magnitude. Combined with the previously mentioned surface brightness distribution along the fundamental plane, this leads to an increased amount of lost flux and a larger error in the output magnitude for brighter objects, as these have lower surface brightnesses. In Section~\\ref{s:lostlight} we will confirm this explanation with predictions based on computations of the visibility function and the lost-light fraction.  Our results confirm the predictions of~\\citet{McGaugh1994}, that the magnitudes of distant spiral galaxies would be measured correctly for luminous galaxies, and slightly underestimate them for the systems with low-luminosity.  The deviations of the measured from the input magnitude we find, are compatible with the simulations of \\citet{Martini2001}, who compared the reliability of different photometric methods for point-sources and exponential profiles. For Kron-like magnitude measurements both shapes would be measured correctly at the bright end, and slightly underestimated at the faint end.  \\begin{figure} \\centering \\includegraphics[width=8.4cm]{fig5} \\caption{Results of a simulation assuming luminosity evolution with redshift. Shown are the simulation results for de Vaucouleurs profiles (left column) and exponential profiles (right column) at a redshift $z=1$ for a seeing of $0.8\"$ (upper row) and $1.6\"$ (lower row). The solid curve shows the results obtained from the normal simulations, the dashed line for a simulation with objects brightened artificially by $0.75$ mag.} \\label{f:lumevol} \\end{figure}  To explore the effects caused by the increase of the objects' surface brightness introduced by luminosity evolution with redshift, an additional set of simulations for objects at $z=1$ was created.  The brightnesses of the created artificial de Vaucouleurs and exponential profiles was increased by $0.75$ magnitudes~\\citep{PSEDE99}.  The results shown in Fig.~\\ref{f:lumevol} show no significant change of the magnitude of the 50\\% completeness limit. The reason for this is that the increase of the detectable area is small in the magnitude range of the 50\\% limit. The increase in size is larger for brighter magnitudes, but the positive effect is mainly lost due to the already high detection probability, or may result in a slightly higher plateau level.    We presented the results of extensive completeness simulations for imaging surveys of faint galaxies, taking into account their known size -- surface-brightness relations. The absolute magnitudes of the elliptical galaxies simulated as de Vaucouleurs profiles following a local fundamental-plane relation. The disk galaxies represented by pure exponential profiles were modelled according to a Freeman law. These physical parameters were converted into apparent sizes, depending on the simulated redshift of the object. The simulations were carried out for the $K'$-band images of the MUNICS survey, but the deficiencies found should remain valid in other comparable surveys as long as a threshold based detection algorithm and an adaptive isophotal magnitude measurement are used.  We find a strong bias against bright $r^1/4$ profiles at higher redshifts in the rate of detection as well in the photometry of the objects. At low redshifts the completeness fraction is comparable to point-like sources. The measured magnitudes underestimate the true magnitudes at the bright end of the distribution. The sharp maximum in the centre of the profile results in an underestimate of the objects' radius leading to a too small isophotal radius in the magnitude measurement. At high redshifts the detection probability does not reach one even for the brightest objects, forming a plateau at lower values. This stems from the correlation of surface brightness and effective radius along the fundamental-plane relation.  Toward the bright end the object's surface brightness decreases, although the total magnitude increases.  The detection of spiral galaxies at high redshift shows similar deficiencies. Here the effective radius increases with magnitude, but the intrinsic mean surface brightness stays constant as predicted by the Freeman law. The rather flat light distribution of the exponential profile helps to estimate the object's radius correctly and accordingly to measure the magnitudes without large errors.  These results are in general agreement with the prediction of the visibility function theory~\\citep{DP1983} and the brightness defects predicted by~\\citet{Dalcanton1998} -- except in the two most distant redshift-bins, where the predictions break down -- using calculations of the lost-light fraction, as long as the effects of seeing are taken into account.  With equations~(\\ref{eq:get_compl}) and~(\\ref{eq:get_compl2}) we provide a rather simple way to estimate the $\\sim 50 \\%$ completeness limit of point-like sources represented by Moffat profiles."
    },
    "0207/astro-ph0207349_arXiv.txt": {
        "abstract": "It has recently been argued that the unidentified SCUBA objects (USOs) are a thick disk population of free-floating dense, compact galactic gas clumps at a temperature of about 7 K. The characteristic mass scale is constrained to be on the order of a Jupiter mass, and the size is about 10 AU. A typical galactic USO is located at a distance from the sun of about 300 pc.  We have calculated the molecular emission lines from these low temperature clouds. We consider three molecules: HD, LiH, and CO. HD is optically thin in the cloud, LiH is a molecule with a large electric dipole moment, and CO is an abundant molecule that is observed in dusty clouds.  Our estimate for the typical object shows that LiH may be  detectable by the future sub-mm array project, ALMA; its expected flux is at the mJy level and the line width is about $10^5$ Hz. Although typical galactic USOs are chemically and dynamically transient, the younger USOs will be recognisable via LiH emission if about a hundred USOs are observed. If USOs are confirmed to be of galactic origin, the total baryonic budget will need to  be reevaluated. ",
        "introduction": "Standard cosmic nucleosynthesis calculations together with the current observational measurements of the abundance of $^2$D, in addition to limits on $^3$He, $^4$He, and $^7$Li, constrain the cosmic baryon density to be (O'Meara et al. 2001) $$ \\Omega_{\\rm b} h^2 = 0.02\\pm 0.002. \\eqno(1) $$ The Hubble constant, $H_0 = 100 h$ km s$^{-1}$ Mpc$^{-1}$, is estimated to be $h=0.72 \\pm 0.08$ (Freedman et al. 2001).  Hence the cosmic baryon density is between 0.03 and 0.05.  There is a similar constraint on $\\Omega_{\\rm b} h^2$ from the CMB (e.g. Ferreira 2001; Bartlett 2001 for reviews), the most recent result being $\\Omega_{\\rm b} h^2 = 0.033 \\pm 0.013$ (Stompor et al. 2001).  Stringent constraints from the CMB anisotropies limit the extreme possibility that baryons account for all of the matter in the universe (Griffiths, Melchiorri, \\& Silk (2001).  In fact, the baryon budget is poorly known (cf.  Persic \\& Salucci 1992; Gnedin \\& Ostriker 1992; Bristow \\& Phillipps 1994; Fukugita, Hogan \\& Peebles 1998). Table 3 of Fukugita et al. summarises the results.  While the stellar contributions are minor, hot plasma in clusters and groups of galaxies is the dominant contributor to the observed baryon budget.  However, for example, the contribution from the intracluster medium is considered to be only in the form of hot plasma. Any warm gas or small, cold clouds in the clusters would contribute to the baryon budget.  While the conclusion of Fukugita et al. is consistent with $\\Omega_{\\rm b}$ as constrained above, the possibility remains that the contribution of small cold clouds is significant.  This has indeed been argued to be the case for the intracluster medium, but a stronger case can perhaps be made for the presence of such clouds in less hot environments, such as the intra-group medium and galaxy halos, where observations are less constraining.  For example, if there is undetected galactic HI and CO, the galactic baryon budget in the form of tiny  very cold clouds may be underestimated.  Possible dynamical structures of such  tiny very cold clouds have been examined by Gerhard \\& Silk (1996). According to these authors, in order to stabilize these  low temperature clouds against collapse, two mechanisms are of interest. One  is gas-dust collisional heating, and the other is the effect of the gravitational potential wells of collisionless dark matter. Other aspects are discussed in many papers (e.g. de Paolis et al. 1995a, 1995b; Henriksen \\& Widrow 1995, Draine 1998; Kalberla et al. 1999; Wardle \\& Walker 1999).  In this $Letter$, we also comment on the dynamical properties very briefly in the next section, focussing on the  observational claim that some very low temperature clouds may have retained or acquired dust (e.g. Lawrence 2001).  Indeed, it has been suggested observationally that halo dark matter could reside in the form of cold dark clouds and so be undetected. Pfenniger \\& Combes (1994) and Pfenniger et al. (1994) argued that such clouds would be at or near the traditional hierarchical fragmentation limit and further argued that flat rotation curves could be explained by baryonic dark matter in this form. Walker \\& Wardle (1998) examined observational constraints in the radio bands for the clouds. The observational possiblity of an emission line forest (i.e. quasi-continuum) from the clouds was first examined by Sciama (2000). An extensive  review of the small--scale interstellar medium, which may be related to these tiny cold clouds, is presented in Heiles (1997). In the same context, Lawrence (2001) has proposed a  possible galactic origin for  the unidentified SCUBA objects (USOs).  In order for the USOs to be galactic, the far IR/mm background and SCUBA source counts at 400 and 800 $\\mu$m, together with dynamical limits on galactic dark matter, constrain the cloud parameters, if distributed throughout the galaxy, to have low temperatures $\\sim 7$K, Jupiter-like masses $M_{\\rm J}$, and to be very small, about 10 AU in size. As long as these characteristics are imposeed because of the dynamical limits, it may be  that USOs contribute to the baryon budget in the intracluster and intra-group medium although they  have a small covering factor over the sky. In this $Letter$, we propose a simple observational strategy for the direct detection of galactic USOs. ",
        "conclusions": "It has been suggested that some halo and disk dark matter could reside in the form of cold dark clouds and so be undetected. In particular, Lawrence (2001) has constrained the physical properties of such halo baryonic dark matter candidate objects by appealing to submm-mm observations. According to our observational predictions for Lawrence's objects, LiH emission lines can in principle be detected by ALMA. This will provide a strong constraint on models for the galactic unidentified SCUBA sources, which we refer to as USOs in this $Letter$. According to Gerhard \\& Silk (1996), galactic USOs can be dynamically transient. In the current work, furthermore, since the chemical depletion of CO and presumably LiH is significant, then they are also chemically transient. Hence, we may need to observe around 100 USOs for clear evidence of the galactic USO candidates that we have postulated. If their reality is established, the contribution of USOs as well as MACHOs to the galactic baryon budget will need to be re-examined."
    },
    "0207/astro-ph0207380_arXiv.txt": {
        "abstract": "s{ We develop a new method for estimating the redshift of galaxy clusters through resolved images of the Sunyaev-Zel'dovich effect (SZE). Our method is based on morphological observables which can be measured by actual and future SZE experiments. The method is tested using a set of high resolution hydrodynamical simulations of galaxy clusters at different redshifts. The method combines the observables in a principal component analysis. We show how this can give an estimate of the redshift of the galaxy clusters. \\\\ Although the uncertainty in the redshift estimation is large, the method should be useful for future SZE surveys where hundreds of clusters are expected to be detected. A first preselection of the high redshift candidates could be done using our proposed morphological redshift estimator.} ",
        "introduction": "The advent of new experiments dedicated to the observation of the Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich, 1972) (SZE hereafter), demands the development of new techniques to best analyze these new and exciting data. \\\\ Through the SZE it is possible to trace the hot plasma in the galaxy clusters which distorts the spectrum of the cosmic background radiation. This distortion is redshift independent and it is proportional to the temperature of the plasma and its electron density ($n_e$). This quality ($z$-independent, and $\\propto n_e$) makes the SZE effect an ideal way to explore the high redshift population of galaxy clusters. However, the fact that the distortion induced by the cluster in the CMB is independent of the redshift of the cluster, makes the determination of the cluster redshift from SZE observations a challenging task. \\\\ Redshifts can be easily measured for relatively nearby clusters but for distant clusters one should use other approaches, for instance photometric redshifts. However, photometric redshifts for clusters above redshift $\\approx$ 0.5 require large telescopes. Ongoing and future SZE experiments will detect hundreds or thousands of clusters. The redshift estimation of all these clusters using 10-m class telescopes is unfeasible. An optimal solution could be to combine small, and medium-sized telescopes to determine the redshifts of the low and intermediate $z$ clusters respectively. Then, we could leave the estimation of the redshifts of the most distant clusters to the large telescopes. However, to select the low, intermediate and high z clusters we need an estimate of their redshift. The motivation of this work is to show how it is possible to make this preselection of the low, intermediate and high redshift clusters using SZE data alone. Then, this preselection can be used to determine more accurately the redshifts of the clusters combining small (for the low z clusters) medium-sized (for the intermediate z clusters) and large (for the high z clusters) telescopes. \\\\ Our method is based only on observed quantities of the SZE. These quantities are associated with the observed shapes of the 2D surface brightness profile of the clusters which does have some dependence on the redshift. The observed size of a particular cluster, for instance, will decrease with increasing redshift. So, given an observed size we could, in principle, get a probability for the cluster to be at a given redshift. However, the size of the cluster will also depend on its total mass. This means that two clusters with different redshifts and masses could have the same apparent size provided the most distant cluster had a larger mass that compensates the decrease in the apparent size due to the increase in redshift. There are, therefore, degeneracies between the redshift of the cluster and its mass. The question now is, can we break this degeneracy by including more information in our analysis ? A resolved SZE image of a cluster provides information, not only about its size, but also about the shape of the entire profile. The total observed flux of the cluster, for instance, will depend on the total mass of the cluster (and its redshift and temperature). The central SZE decrement will depend on the core radius and electron central density but not on the redshift. By adding these and other additional observables, it should be possible to break the degeneracies between the mass and redshift. \\\\ Our method will have, however, one limitation: it works with resolved SZE images. Therefore it should be useful for sub-arcmin experiments but not for experiments like Planck where the best resolution will be 5 arcmin. \\\\ ",
        "conclusions": ""
    },
    "0207/astro-ph0207663_arXiv.txt": {
        "abstract": "Deep spectroscopy of the two millijansky radio galaxies LBDS 53W069 and LBDS 53W091 has previously shown them to have old ($\\ga3$~Gyr) stellar populations at $z\\simeq1.5$.  Here we present the results of {\\it Hubble Space Telescope (HST)\\/} observations with the Wide Field and Planetary Camera 2 (WFPC2) in F814W and with the Near-Infrared Camera and Multi-Object Spectrograph (NICMOS) in F110W.  We find that 53W069 has a de Vaucouleurs $r^{1/4}$ profile in both the F814W \\& F110W data with a mean effective radius of $0\\farcs30\\pm0\\farcs06$ ($2.7\\pm0.5$~kpc).  The restframe $U-B$ colour gradient is consistent with that of present-day ellipticals, requiring a stellar population of super-solar ($3Z_{\\sun}$) metallicity that formed on a very short timescale at high redshift ($z>5$).  53W091 has a regular $r^{1/4}$ profile in F110W with an effective radius of $0\\farcs32\\pm0\\farcs08$ ($2.9\\pm0.7$~kpc).  The F814W profile is more extended and is consistent with the presence of a blue exponential disk that contributes $20\\pm10$\\% of the flux within $r_e$.  We find a restframe $U-B$ colour gradient which is significantly larger than that observed in field ellipticals at $z\\le1$, implying a stellar population of mixed metallicity (1--$3Z_{\\sun}$) that formed in a high-redshift rapid burst.  We have compared these two LBDS radio galaxies with the Kormendy relations of ten 3CR radio galaxies at $z\\simeq0.8$ and a sample of cluster ellipticals at $z\\sim0.4$.  The LBDS galaxies follow the Kormendy relation for the more radio-luminous 3CR galaxies, assuming passive evolution of their stellar populations, although they are smaller than the 3CR galaxies whose mean effective radius is 12~kpc. Their sizes and radio luminosities are consistent with scaling relations applied to the 3CR galaxies, in which both radio power and effective radius scale with galaxy mass.  Compared with the sample of cluster ellipticals, 53W069 \\& 53W091 lie well within the scatter of the Kormendy relation.  We conclude that the hosts of these millijansky radio sources at $z\\simeq 1.5$ are passively-evolving elliptical galaxies that will evolve into ordinary $L^*$ ellipticals by the present day. ",
        "introduction": "How did elliptical galaxies form and evolve?  There are two mechanisms that have been proposed.  First is the monolithic collapse model, in which elliptical galaxies form at high redshifts in an intense burst of star formation \\citep{Eggen62}.  Essentially all the final mass of the galaxy is already present within its potential well at the time of formation.  The second mechanism is the merger model, in which ellipticals form at relatively lower redshifts from the merger of many smaller galaxies \\citep[e.g.,][]{Kauffmann98}.  In this model, the mass of the galaxy increases with time as more and more smaller galaxies are canibalized by the forming elliptical.  Recent observations suggest that both mechanisms may have a role to play.  For example, \\citet{Ellis97} found that spheroidal galaxies in clusters at $z\\simeq0.5$ required a formation redshift (by which we mean the redshift at which the dominant stellar population formed) of $z_f\\ge3$.  In contrast, \\citet{Zepf97} found that the small number of very red galaxies in deep optical/infrared surveys suggested either that ellipticals form at moderate redshifts (and are enshrouded by dust), or that they assemble through the merging of smaller galaxies. Similarly \\citet{Pascarelle96} identified a group of Lyman-$\\alpha$ emitters at $z=2.39$ that they suggested may subsequently merge at a lower redshift into one or more luminous galaxies.  \\citet{Treu99} concluded from their study of NICMOS parallel fields that a significant fraction (10--66\\%) of the elliptical galaxy population formed at $z_f\\ge3$, but the rest may have been formed (or at least assembled) at lower redshifts.  One important method of studying high-redshift ($z\\ga1$) ellipticals is via radio selection.  At low redshifts, luminous radio sources are almost exclusively hosted by giant elliptical galaxies containing old stellar populations \\citep*[e.g., ][]{Kron85,Taylor96,Nolan00a}.  To the extent that radio galaxies at high redshifts are also hosted by giant ellipticals, they can be used to study the evolution of the elliptical galaxy population.  \\citet{Lilly84} obtained infrared $K$ magnitudes for a subsample of the 3CR radio survey, from which they constructed a $K$-band Hubble diagram. They concluded from this $K$--$z$ relation that luminous radio galaxies at $z\\sim1$ are giant ellipticals with passively-evolving stellar populations. \\citet{Lilly85} and \\citet{Dunlop90} found that the $K$--$z$ relation for less-powerful radio galaxies led to a similar conclusion.  At redshifts $z\\ga0.6$, \\citet{Eales97} found that the 6C radio galaxies, with radio luminosities a factor of $\\sim5$ lower than the 3CR, were also on average 0.6~mag fainter than 3CR in the $K$-band. \\citet{Roche98} investigated the $K^\\prime$-band morphologies of ten of these 6C sources at $z\\sim1.1$, showing that seven of them were normal ellipticals and the other three were ongoing or recent mergers. The radii of the 6C galaxies were significantly smaller than those of the 3CR sources at similar redshifts \\citep{Best98,Zirm99,McLure00}, indicating that their fainter $K$ magnitudes were due to their smaller size, not solely due to the difference in power of their AGN.  The Leiden-Berkeley Deep Survey \\citep[LBDS; ][]{Windhorst84,Waddington00}, with a flux density limit of 1~mJy at 1.4~GHz, is sensitive to radio sources a factor of $\\sim200$ fainter than the 6C survey, thus probing lower radio luminosities and higher redshifts.  \\citet*{Windhorst98} showed that the radio galaxy LBDS 53W002 has a weak AGN and a dominant $r^{1/4}$ profile with the colours of a $\\sim 3\\times 10^8$~year young stellar population at $z=2.39$, suggesting that these weak radio sources could have formed rather mature ellipticals at high redshifts.  Two more of these sources have proved particularly useful for the study of galaxy evolution, due to the deep spectra that we obtained with the Keck telescope.  LBDS 53W091 has a redshift of 1.552, and its restframe UV spectrum is best modeled by a stellar population $\\ge3.5$~Gyr old \\citep{Dunlop96,Spinrad97}.  LBDS 53W069 is at $z=1.432$ and has a best-fitting stellar population of age $\\ge4$~Gyr \\citep{Dey97,Dunlop99b}.  \\citet{Spinrad97} investigated the age determination in detail and showed how the minimum age of the universe at $z\\simeq1.5$, as required by the age of 53W091, placed limits on the allowed values of the cosmological parameters ($H_0$, $\\Omega_{\\rm M}$, $\\Omega_\\Lambda$).  \\citet{Stockton95} similarly showed that the age inferred from the Keck spectrum of the radio galaxy 3C65 ($\\sim4$~Gyr at $z=1.175$) required a formation redshift of $z_f\\ge5$.  Several authors have questioned the reliability of these age estimates, for example \\citet{Bruzual99} and \\citet{Yi00} derive ages of 1.5--2~Gyr for 53W091.  However, \\citet{Dunlop99} argued that such young ages are only deduced if the near-infrared photometry is included in the model fitting; if the fitting is confined to the spectroscopic data then a variety of stellar population sythesis codes consistently produce ages of $\\ge2.5$~Gyr (see also Nolan \\etal\\ 2001b\\nocite{Nolan00b}).  We will not discuss this age controversy further here, but simply note that nothing in the current paper depends crucially on the precise ages of these sources.  In this paper we will investigate the morphologies of these two radio galaxies with the {\\it Hubble Space Telescope's (HST)} Wide Field and Planetary Camera 2 (WFPC2) and Near-Infrared Camera and Multi-Object Spectrograph (NICMOS).  At $z\\simeq1.5$, the 4000-\\AA\\ break is straddled by the F814W and F110W filters, and thus the emission observed through these two filters is dominated by the young and old stellar populations respectively.  In section 2, we describe the observations and discuss the processing steps that were applied to the data.  In section 3, we investigate the surface brightness profiles and colour gradients of the two sources.  The location of the galaxies on the Kormendy relation is the topic of section 4 and finally, in section 5, we comment on the apparent relations between radio luminosity, black-hole mass and effective radius of these sources. Our results are summarized in section 6.  We use AB magnitudes unless otherwise noted, and denote magnitudes in the three \\hst\\ filters by $I_{814}$ for F814W (WFPC2), $J_{110}$ for F110W (NICMOS) and $H_{160}$ for F160W (NICMOS).  We assume a flat cosmology with $H_0=65$~\\kmsmpc, $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "We have used \\hst\\ observations of the two millijansky radio galaxies LBDS 53W069 and 53W091 to investigate their optical morphologies at $z\\simeq 1.5$.  In both F814W (restframe $U$-band) and F110W (restframe $B$-band) 53W069 is best described by an elliptical (de Vaucouleurs) model, of effective radius $0\\farcs3$ or 3~kpc.  The $U-B$ colour gradient indicates that the galaxy formed in a rapid burst of star formation at high redshift ($z>5$) and has a high metallicity ($3Z_{\\sun}$) stellar population.  In F110W, 53W091 is similarly modelled as an elliptical galaxy of effective radius $0\\farcs3$ (3~kpc).  At shorter wavelengths (F814W), 53W091 is more extended, consistent with there being a faint blue disk contributing $\\sim20$\\% of the flux within $r_e$.  The $U-B$ colour of the galaxy within a radius of 0\\farcs4 indicates a high metallicity ($3Z_{\\sun}$), high-redshift ($z>5$) burst of star formation, whereas the metallicity falls to solar at larger radii.  Assuming passive evolution of their stellar populations, these two elliptical galaxies lie on the $B$-band Kormendy surface brightness--effective radius relations of both 3CR radio galaxies and cluster ellipticals.  Their sizes and radio luminosities are consistent with scaling relations applied to the 3CR radio galaxies, in which both radio power and effective radius scale with galaxy mass.  Our analyses of the restframe UV spectra of these two galaxies demonstrated that they were passively-evolving ellipticals, with ages $\\ga 3$~Gyr at $z\\simeq 1.5$ \\citep{Dunlop96,Spinrad97,Dunlop99b,Nolan02}.  The age of the universe at this redshift is 4.5~Gyr in the flat, lambda-dominated cosmology that we assume here, requiring a formation redshift of $z_f\\ga 5$.  In \\citet{Peacock98} we compared the primordial density fluctuation spectrum required by the existence of these galaxies with that inferred from Lyman-$\\alpha$ absorbers and Lyman-break galaxies. Those results indicated a similarly high redshift of gravitational collapse, $z_c\\simeq 6$--8.  Here we have demonstrated that the morphologies and internal colours of these galaxies similarly show them to be ellipticals whose last major episode of star formation was at very high redshift.  These results imply that the last major merger was also some while in the past, since such an event would produce significant star formation and an irregular morphology, for which we see no strong evidence at the epoch at which we see them.  Although it may be that the faint blue disk component in 53W091 is the remains of the galaxy's last merger. In this context, we note that 53W091 has several companions with similar colours, one of which has been spectroscopically confirmed to have the same redshift of 1.55 \\citep{Spinrad97,Bunker00}.  It is possible, even likely, that one or more of these companions will merge with 53W091 in the future.  In contrast, 53W069 has no such companions.  We began the paper by asking how ellipticals formed -- via monolithic collapse or hierarchical mergers?  We end by concluding that we cannot be certain, even in the particular case of these two LBDS radio galaxies.  The most likely interpretation of the data is that they have not evolved significantly since high redshifts ($z\\sim 5$), save for passive stellar evolution.  However, their formation could still have been via mergers at early times, and at least in the case of 53W091, it seems probably that further mergers lie ahead.  What is clear, is that these radio galaxies are quite ordinary ellipticals, and if they continue to evolve passively until the present day they will become typical $L^*$ elliptical galaxies."
    },
    "0207/astro-ph0207039_arXiv.txt": {
        "abstract": "{ The $\\sim$15.1 years period found in the long-term $UBV$ photoelectric and photographic photometry of the symbiotic nova V1016 Cyg is detected also in the $(J-K)$ colour index and in the UV continuum and emission line fluxes from IUE and HUT spectra. It could be interpreted either as the effect of recurrent enhanced mass loss episodes from the Mira type variable companion to a hot component along its ultra-wide orbit (proposed from recent HST observations) or the true orbital period of the inner, unresolved binary of a triple system. A 410-day delay of the maximum of UV emission lines fluxes with respect to the maximum of continuum was found. The pulsation period of the Mira type variable was improved to 474$\\pm$6 days. ",
        "introduction": "V1016 Cyg (MH$\\alpha$ 328-116) is a member of a small subgroup of symbiotic stars, called symbiotic novae, also including V1329 Cyg and HM Sge, whose outbursts lead to a nebular spectrum (M\\\"urset \\& Nussbaumer \\cite{mur}). Symbiotic novae are wide interacting binaries, where matter from a late-type giant is transferred onto the surface of the more compact companion. The nova-like optical outburst ($\\Delta m \\sim 5-7$ mag), lasting decades, is caused by a thermonuclear runaway on the surface of a wind-accreting white dwarf after the critical amount of material has been accumulated (cf. Mikolajewska \\& Kenyon, \\cite{mik}). V1016 Cyg underwent such nova-like outburst in~1964 (McCuskey \\cite{mcc}). The object is classified as a D-type symbiotic, the cool component being a Mira type variable embedded in a dust envelope whose pulsation period turned out to be $\\sim$478 (Munari \\cite{mun}). The onset of a dust formation episode in 1983 is reported by Taranova \\& Yudin (\\cite{tar2}).  The orbital period of V1016 Cyg is not yet established. Taranova \\& Yudin (\\cite{tar1}) made use of the increase of Balmer emission lines in combination with the appearance and disappearance of FeII lines to derive an orbital period of $\\approx$20 years. Afterwards, Nussbaumer \\& Schmid (\\cite{nuss2}), though unable of recording two consecutive maxima, proposed an orbital period of 9.5 years on the basis of the apparent periodicity seen in the flux of OI and MgII UV emission lines by means of the IUE satellite. At the same time Munari (\\cite{mun}), by resorting to IR observations taken over two decades, proposed instead a 6-year orbital period by modeling the sequence of dust obscuration episodes, likely related to the passage of the Mira type variable at the inferior conjunction in the system.   Much longer periods have been proposed by Wallerstein (\\cite{wall}) and Schild \\& Schmid (\\cite{schi}). In the former paper the author, assuming that the sharp FeII emission lines are formed in the chromosphere of the cool star so as to reflect its orbital motion, concludes that their observed radial velocities between 1978 and 1985 limit any high inclination orbit to a period greater than 25 years or to a large eccentricity. The latter analysis, based on spectropolarimetric data taken from 1991 to 1994, indicates that the orbital period is about 80 $\\pm$ 25 years, though later observations obtained in 1997 put this result into question (Schmid, \\cite{schm}). Finally, Brocksopp et al. (\\cite{bbe}), adopting a projected separation of the two stellar components as large as 84 AU on the basis of their HST/WFPC2 images, are forced to propose an astonishingly long orbital period of $\\sim$544 years.  \\begin{figure*} \\centerline{\\resizebox{18cm}{!}{\\includegraphics{h3590F1.eps}}} \\vspace*{-4.5cm} \\hfill \\caption{Photographic (left) and $UBV$ (right) light curves of V1016 Cyg. Arrows $a,b,c$ and $d$ mark the epochs of subsequent activity episodes of the system (see text).} \\end{figure*}  In the context of this still-to-be-settled issue, one has also to interpret the evidence for the 15-year periodic activity (not necessarily representing the sought {\\em orbital} period) recently presented by Parimucha et al. (\\cite{pac}) (hereafter PAC) who gathered and analysed long-term photographic, photoelectric and visual photometry of the object. The aim of the present paper is indeed to give further (i.e., multiwavelength) support to the existence of such a periodicity making use of both IR photometry and UV spectroscopy, as well as to investigate its own origin. ",
        "conclusions": "No doubt the symbiotic nova V1016 Cyg includes a pulsating Mira type variable and an accreting white dwarf that underwent a thermonuclear outburst in 1964, leading straight to a nebular spectrum (FitzGerald et al. \\cite{fitz}). According to the ionization model of symbiotic binaries, the hot luminous component ionizes the neutral wind of the giant giving rise to the nebula in the system. Such a mechanism is confirmed by the strict coincidence of the epochs of maxima in the $U$ passband (observed in 1980 and 1994) and those recorded in the space-UV continuum combined with the 410-day delay of the OIII], CIII], NIII] and SiIII] emission line maxima which, in turn, assures that the lines are collisionally excited in the surrounding nebula when the fast wind from the hot object interacts with the slow wind from the Mira type variable.  At present, the above well-understood phenomenology leaves still open two alternative scenarios for V1016 Cygni.  If this symbiotic nova is a wide-orbit binary with a 544-year period as proposed by Brocksopp et al. (\\cite{bbe}), one is forced to interpret the periodic (15-yr) variations of the optical and UV continuum as induced by flares of the hot component triggered by the recurrent enhanced mass loss episodes from the Mira type variable companion. According to Fleischer et al. (\\cite{fgs}), this mass loss is most pronounced every few periods of the Mira type variable pulsations. It can trigger individual flares in the accretion disk of the hot white dwarf. The existence of the disk is supported by the 3D simulation of the wind accretion by the compact star (Theuns \\& Jorissen, \\cite{thjo}).   Alternatively, V1016 Cygni could be interpreted as a {\\em triple} system, whose {\\em inner}, unresolved binary has an {\\em orbital} period of 15 years. If this is the case, flares would be triggered by an enhanced mass transfer from a cool giant {\\em during its periastron passage on the 15-year orbit}. In this respect, a hint of enhanced mass transfer, coming just before the weak flare in 1994, can be found in the long-term infrared photometry performed by Taranova \\& Schenavrin (\\cite{tar3}) showing a brightness increase in the $J$ and $H$ passbands in 1992. Moreover, the coincidence of the observed maximum of the $J-K$ index in 1988 (Fig.~3) and the wide minima of the UV continua (Fig.~5) would constrain to that epoch the inferior conjuction of a cool giant along its 15-year orbit."
    },
    "0207/astro-ph0207513_arXiv.txt": {
        "abstract": "Most dusty winds are described by a set of similarity functions of a single independent variable that can be chosen as \\tV, the overall optical depth at visual. The self-similarity implies general scaling relations among the system parameters, in agreement with observations. Dust drift through the gas has a major impact on the structure of most winds. ",
        "introduction": "The complete description of a dusty wind should start with a full dynamic atmosphere model and incorporate the processes that initiate the outflow and set the value of \\Mdot. These processes are yet to be identified with certainty, the most promising are stellar pulsation (e.g. Bowen 1989) and radiation pressure on the water molecules (e.g. Elitzur, Brown \\& Johnson 1989). Proper description of these processes should be followed by grain formation and growth, and subsequent wind dynamics. Two ambitious program attempting to incorporate as many aspects of this formidable task as possible have been conducted over the past few years by groups at Berlin (see J.M.\\ Martin, these proceedings) and Vienna (see Dorfi et al 2001). While much has been accomplished, the complexity of this undertaking necessitates simplifications such as a pulsating boundary. In spite of continuous progress, detailed understanding of atmospheric dynamics and grain formation is still far from complete.  Fortunately, the full problem splits naturally to two parts, as recognized long ago by Goldreich \\& Scoville (1976). Once radiation pressure on the dust exceeds all other forces, the rapid acceleration to supersonic velocities decouples the outflow from the earlier phases---{\\em the supersonic phase would be exactly the same in two different outflows if they have the same mass-loss rate and grain properties even if the grains were produced by entirely different processes.} Furthermore, these stages are controlled by processes that are much less dependent on detailed micro-physics, and are reasonably well understood. And since most observations probe only the supersonic phase, models devoted exclusively to this stage should reproduce the observable results while avoiding the pitfalls and uncertainties of dust formation and the wind initiation. ",
        "conclusions": ""
    },
    "0207/hep-th0207130_arXiv.txt": {
        "abstract": "We review physical  motivations, phenomenological consequences, and open problems of the so-called pre-big bang scenario in superstring cosmology. ",
        "introduction": "\\label{Sec1} \\setcounter{equation}{0} \\setcounter{figure}{0}  During the past thirty years, mainly thanks to accelerator experiments of higher and higher energy and precision, the standard model of particle physics  has established itself as the uncontested winner in the race for a consistent description of electroweak and strong interaction phenomena at distances above $10^{-15}$ cm or so. There are, nonetheless, good reasons  (in particular the increasing evidence for non-vanishing neutrino masses \\cite{KaTo01,SNO,SNOa})  to believe that the standard model is not the end of the story. The surprising validity of this model at energies below $100$ GeV, as well as the (in)famous Higgs mass fine-tuning problem, suggest some supersymmetric extension of the standard model  (for a review see \\cite{Nilles}) as the most likely improved description of non-gravitational phenomena over a few more decades in the ladder of scales. It is however quite likely that other questions that are left unanswered by the standard model, such as the peculiarities of fermionic masses and mixings, the family pattern, C, P, CP, B violation, etc., will only find their answers at --or around-- the  much higher energies at which all gauge interactions appear to unify \\cite{Ama91}.   This energy scale appears to be embarrassingly close (on a logarithmic scale) to the so-called Planck mass, $M_{\\rm P} \\sim 10^{19}$ GeV, the scale at which gravity becomes strong and needs to be quantized.  The situation with gravitational phenomena is completely different. Even the good old Newton law is known to be valid only down to the $1$ mm scale \\cite{Hoyle01}, so that much interest has been  devoted  to the possibility of large modifications of gravity below that distance,  either from new forces mediated by light scalars such as the dilaton of string theory \\cite{Tay88}, or from the existence of large extra dimensions felt exclusively  by gravity  \\cite{Arkani98,RS2}. General relativity is well tested at large scales; nevertheless; present evidence for a (small) vacuum energy density \\cite{Riess98,Perlmutter99} suggests that, even on cosmologically  large distances, the  strict Einstein theory might turn out to be inadequate. Evidently, the construction of a standard model for gravity  and cosmology lags much behind its particle physics counterpart.  The hot big bang model (see for instance \\cite{Weinberg}), originally thought of as another great success of general relativity, was later discovered  to suffer from huge fine-tuning problems. Some of these conceptual  problems are solved by the standard inflationary paradigm (see \\cite{Lin90,KT90} for a review), yet inflation remains a generic idea in search of a theory that will embody it naturally. Furthermore, the classical theory of inflation does not really address the problem of how the initial conditions needed for a successful inflation came about. The answer to  this question is certainly related to even more fundamental issues, such as: How did it all start? What caused the big bang? Has there been a singularity at $t=0$? Unfortunately, these questions lie deeply inside the short-distance, high-curvature regime of gravity where quantum corrections cannot be neglected. Attempts at answering these questions using quantum cosmology based on Einstein's theory has resulted in a lot of heated discussions  \\cite{Linde98,TuHa98}, with no firm conclusions.  It is very likely that both a standard model for gravity and cosmology and a full understanding of the standard model of particle physics will require our understanding of physics down to the shortest scale, the Planck length $\\la_{\\rm P}\\sim 10^{-33}$ cm. Until the Green--Schwarz revolution of 1984 \\cite{GSW87}, the above conclusion would have meant postponing indefinitely those kinds of questions. Since then, however, particle theorists have studied and developed superstring  theory (see \\cite{Pol98} for a   recent review, as well as \\cite{Greene99} for a non-specialized introduction), which appears to represent a consistent framework not only for addressing  (and possibly answering)  those questions, but even for unifying our understanding of gravitational and non-gravitational phenomena, and therefore for relating the two classes of questions.  The so-called ``pre-big bang\" scenario described in this report  has to be seen in the above perspective as a possible example, even just as a toy model, of what cosmology can look like if we assume that the sought for standard model  of gravity and cosmology is based on (some particular version of) superstring theory. Although most string theorists would certainly  agree on the importance of studying the cosmological consequences of string theory, it is a priori far from obvious that the ``state of the art\" in this field can provide an unambiguous answer to this question. Indeed, most of our understanding of superstring theory is still based on  perturbative expansions,  while most of the recent progress in non-perturbative string theory has been achieved in the context of   ``vacua\" (i.e. classical solutions to the field equations) that respect a large number of supersymmetries \\cite{Pol98}. By contrast, our understanding of string theory at large curvatures and couplings, especially in the absence of supersymmetry, is still largely incomplete.  A cosmological background, and a fortiori one that evolves rapidly in time, breaks (albeit spontaneously) supersymmetry. This is why the Planckian regime of cosmology appears to be intractable for the   time being.  It is very fortunate, in this respect, that in the pre-big bang  scenario  the Universe is supposed to emerge from a highly perturbative initial state {\\it preceding} the big bang. Therefore,  early enough before (and late enough after) the big bang, we may presume to know the effective theory to be solved. The difficult part to be dealt with non-perturbatively remains the transition from the pre- to the post-big bang regime, through a high-curvature (and/or possibly a large-coupling) phase. Thus, from a more phenomenological standpoint, the relevant question becomes: Are the predictions of the pre-big bang scenario robust with respect to the details of  the non-perturbative phase?   It is difficult of course to give a clear-cut answer to this question, but an analogy with QCD and the physics of strong interactions may be helpful. Because of asymptotic freedom, QCD can be treated perturbatively at short distance (high momentum transfers). However, even ``hard\" processes such as $e^+ e^- \\rightarrow {\\rm hadrons}$ are not fully within perturbative control.  Some soft non-perturbative physics  always gets mixed in at some level, e.g. when partons eventually turn into hadrons. The reason why certain sufficiently inclusive quantities are believed to be calculable is that  large- and short-distance   physics ``decouple\", so that, for instance, the hadronization process does not   affect certain ``infrared-safe\" quantities, computed at the quark--gluon level.  In the case of string  cosmology the situation should be similar, although somehow reversed \\cite{GVGvsG95}.  For gravity, in fact, the large-distance, small-curvature regime is easy to deal with, while the short-distance, high-curvature is hard. Yet, we shall argue that some consequences of string cosmology, those concerning length scales that were very large with respect to the string scale (or the horizon) in the high-curvature regime, should not be affected (other than by a trivial kinematical redshift)   by the details of the pre- to post-big bang transition. The above reasoning does not imply, of course, that string theorists should not address the hard, non-perturbative questions {\\it now}. On the contrary, the ``easy\" part of the   game will provide precious information about what the relevant hard questions are,  and on how to formulate them.  Finally, possible reservations on a ``top--down\" string cosmology approach may naturally arise from a cosmology community accustomed to a  data-driven, ``bottom--up\" approach. We do  believe ourselves that a good model of cosmology is unlikely to emerge from theoretical considerations alone. Input from the data will be essential in the selection among various theoretical alternatives.  We also believe, however, that a balanced combination of theoretical and experimental imput should be the best guarantee for an eventual success.  Insisting on the soundness of the underlying theory (e.g. on its renormalizability)  was indeed  essential in the progressive construction of the standard model, just as were  the quantity and the quality of  experimental data. Cosmology today  resembles the particle physics of the  sixties: there is no shortage of data, and these are becoming more and more precise but also more and more challenging while, theoretically,  we are still playing with very phenomenological (even if undoubtedly successful) models, lacking a clear connection to other branches of fundamental physics, and therefore remaining largely unconstrained.  \\subsection{Coping with a beginning of time} \\label{Sec1.1}  Both the standard Friedmann--Robertson--Walker (FRW) cosmological scenario \\cite{Weinberg} and  the standard inflationary scenario \\cite{Gut81,KT90,Lin90}  assume that time had a beginning. Many of the problems with the former model simply stem  from the fact that, at the start of the classical era, so little time had elapsed since the beginning. Indeed, in the FRW framework, the  proper size of the (now observable) Universe was about $10^{-3}$ cm across at the start of  the classical era, say at a time of the order of a few Planck times, $t_{\\rm P} \\sim 10^{-43}$ s.  This is of course a very tiny Universe with respect to  its present size ($\\sim 10^{28}~ \\rm{cm}$), yet it is huge with  respect to the horizon  (the distance travelled by light) at that time,  $\\la_{\\rm P}  = c t_{\\rm P} \\sim 10^{-33}$ cm.  In other words, a few Planck times after the big bang, our observable Universe consisted of about $(10^{30})^3 = 10^{90}$ Planckian-size, causally disconnected  regions. Simply not enough time had elapsed since the beginning for the Universe to become homogeneous (e.g. to thermalize) over its entire size. Furthermore, soon after $t=t_{\\rm P}$, the Universe must have been characterized by a huge hierarchy between its Hubble radius, on the one hand, and its spatial-curvature radius,  on the other. The relative factor of (at least) $10^{30}$ appears as an incredible amount of fine-tuning on the initial state of the Universe, corresponding to a huge asymmetry between space and time derivatives, or, in more abstract terms, between intrinsic and extrinsic curvature.   Was this asymmetry really there? And, if so, can it be explained in any, more natural way?  The conventional answer to the difficulties of the standard scenario is to wash out inhomogeneities and spatial curvature by introducing, in the history of the Universe, a long period of accelerated expansion, called inflation \\cite{Gut81,KT90,Lin90}. It has been pointed out, however, that standard inflation cannot be  ``past-eternal\" \\cite{BoGuVil01} (and cannot avoid the initial singularity \\cite{Vil92,BoVil94}), so that the question of what preceded inflation is very relevant. Insisting on the assumption that the Universe (and time itself) started at the big bang leaves only the possibility of having post-big bang  inflation mend an insufficiently smooth and flat Universe arising from the big bang.  Unfortunately, that solution has its own problems, for instance those of fine-tuned initial conditions for the inflaton field and its potential. A consistent quantum cosmology approach giving birth to a Universe in the ``right\" initial state is still  much under debate \\cite{Haw84,Vil84,Linde84,Zeldo84,Rub84}. Furthermore, the inflaton is introduced {\\em ad hoc} and inflation is not part  of  a grander theory of elementary particles and fundamental interactions such as superstring theory.  In spite of its possible importance, and of repeated motivated attempts \\cite{EENQ86,MP86,BG86},  a conventional realization of an inflationary phase in a string theory context is in fact problematic \\cite{CAO90}, in particular because the dilaton  --the fundamental string theory scalar-- cannot be (at least trivially) identified with the inflaton  --the fundamental scalar of the standard inflationary scenario \\cite{BS93}.  Here we shall argue that, instead, superstring theory gives strong hints in favour of a totally different approach to solving the problems of the standard cosmological scenario. This new possibility arises if we  assume that, in string theory, the big bang singularity is fictitious and that it makes therefore sense to ``continue\" time to the   past of the big bang itself.  \\subsection{Inflation before the big bang} \\label{Sec1.2}  If the history of the Universe can be continued backward in time past the big bang, new possibilities arise for a causal evolution to have {\\it produced}  a big bang with the desired characteristics. The actual pre-big bang scenario  presented  in this report is just one possible realization of the above general idea. Since, as we shall see, it is easy to generate a phase of pre-big bang inflation  driven by the kinetic energy of the dilaton (somewhat in analogy with kinetic-inflation ideas \\cite{Levin}), we will discuss, as the simplest possibility,  a minimal cosmological scenario, which avoids making use of standard (i.e. potential-energy-driven)  post-big bang inflation.  This does not, though, that pre- and post-big bang inflation are mutually exclusive or incompatible. Should near-future high-precision experiments definitely indicate that an inflation that is exclusively of the pre-big bang type is disfavoured with respect to conventional, post-big bang, ``slow-roll\" inflation, one should ask whether a pre-big bang phase can naturally lead to ``initial\" conditions suitable for igniting an inflationary epoch of the slow-roll type, rather than a standard, non-inflationary, FRW cosmology.  One model-independent feature of pre-big bang cosmology is clear:  by its very definition, the pre-big bang phase should be  an evolution towards --rather than away from-- a high-curvature regime. As we shall see in Section \\ref{Sec2}, this is precisely what the symmetries of the string cosmology equations suggest, an unconventional realization of the inflationary scenario, in which the phase of accelerated cosmological evolution  occurs while the Universe is approaching --rather than getting away from-- the high-curvature, Planckian regime.  The main difference between the string cosmology and the standard inflationary scenarios can therefore be underlined through  the opposite behaviour of the curvature scale as a function of time, as  shown in Fig. \\ref{f11}. As we go backward in time, instead of a monotonic growth (predicted by the standard scenario), or of a ``de-Sitter-like\" phase of nearly constant curvature (as in the conventional inflationary picture), the curvature grows,  reaches a maximum controlled by the string scale $M_{\\rm s}=\\la_{\\rm s}^{-1}$,  and then starts decreasing towards an asymptotically flat state, the string perturbative vacuum. The big bang singularity is regularized by a ``stringy\" phase of high but finite curvature, occurring at the end of the initial inflationary evolution.  We should warn the reader, from the very beginning of this review, that this scenario is far from being complete and understood in all of its aspects, and that many important problems are to be solved still. Nevertheless, the results obtained up to now have been encouraging, in the sense that it now seems possible to formulate models for the pre-big bang evolution of our Universe that fit consistently in a string theory context, and which are compatible with various phenomenological and theoretical bounds. Not only: the parameter space of such models seems to be accessible to direct observations in a  relativiely near future and, at present, it is already indirectly constrained by various astrophysical, cosmological, and particle physics data.  \\begin{figure}[t] \\centerline{\\epsfig{file=f11.ps,width=72mm}} \\vskip 5mm \\caption{\\sl Qualitative evolution of the curvature scale in the standard cosmological model, in conventional inflationary models and in string-cosmology models.} \\label{f11} \\end{figure}  To close this subsection we should mention, as a historical note, that the idea of a phase of growing curvature preceding that of standard decelerated expansion, is neither   new  in cosmology, nor  peculiar to string theory. Indeed, if the growth of the curvature corresponds to a contraction, it is reminiscent of Tolman's cyclic Universe \\cite{Tol}, in which the birth of our  present Universe is preceded by a phase of gravitational collapse (see also \\cite{DurLau96,BP91,Novello93}). Also, and more conventionally, the growth of the curvature may be implemented as a phase of Kaluza--Klein superinflation \\cite{Sha84,Abb84,Kolb84}, in which the accelerated expansion of our three-dimensional space is sustained by the  contraction of the internal dimensions and/or by some exotic source, with the appropriate equation of state (in particular, strings \\cite{GSV91} and extended objects).  In the context of general relativity, however, the problem is how to avoid the curvature singularity appearing at the end of the phase of growing curvature. This is in general impossible, for both contraction and superinflationary expansion, unless one accepts rather drastic modifications of the classical gravitational theory. In the contracting case, for instance, the damping of the curvature and a smooth transition to the phase of decreasing curvature can be arranged through the introduction of a non-minimal and gauge-non-invariant coupling of gravity to a cosmic vector \\cite{Novello78} or scalar \\cite{SatSi86,Barrow93} field,  with a (phenomenological) modification of the equation of state in the Planckian curvature regime \\cite{Rosen85,Wesson85}, or with the use of a non-metric, Weyl-integrable connection \\cite{Novello93}. In the case of superinflation,  a smooth transition can be arranged through a breaking of the local Lorentz symmetry of general relativity \\cite{Gas85,Gas98a}, a geometric contribution of the spin of the fermionic sources \\cite{Gas86}, or the embedding of the space-time geometry into a more fundamental quantum phase-space dynamics \\cite{Caia91,Gas91}. In the more exotic context of topological transitions, a smooth evolution from contraction to expansion, through a state of minimal size, is also obtained with the adiabatic compression and the dimensional transmutation of the de Sitter vacuum \\cite{GasUm91}.  In the context of string theory, on the contrary, the growth of the curvature is naturally associated to the growth of the dilaton and of the coupling constants (see for instance Section \\ref{Sec2}). This effect, on the one hand,  sustains the phase of superinflationary expansion, with no  need of matter sources or extra dimensions. On the other hand, it necessarily leads the Universe to a regime in which not only the curvature but also the couplings become strong, so that typical ``stringy\" effects   become important and  are expected to smooth out the curvature singularity. This means that there is no need to look for more or less {\\em ad hoc}  modifications of the theory,  as string theory itself is expected to provide the appropriate tools for a complete and self-consistent cosmological scenario.   \\subsection{Pre-big bang inflation and conformal frames} \\label{Sec1.3}   While  postponing to the next section the issue of physical motivations, it is important to classify the various inflationary possibilities just from their kinematical properties. To be more precise, let us consider the so-called flatness problem (the arguments  are similar, and the conclusions the same, for the horizon problem mentioned in Subsection \\ref{Sec1.1}). We shall assume, on the basis of the approximate isotropy observed at large scale, that our present cosmological phase can be correctly described by the ordinary Einstein--Friedmann equations.  In that case, the gravitational part of the equations contains two  contributions from the metric: $k/a^2$, coming from the spatial (or intrinsic) curvature, and $H^{2}$, coming from the gravitational kinetic  energy (or extrinsic curvature). Present observations imply that the spatial curvature term, if not negligible, is at least non-dominant, i.e. \\beq r^2={k\\over a^2H^2} \\laq 1. \\label{11} \\eeq  On the other hand, during a phase of standard, decelerated expansion, the ratio $r$ grows with time. Indeed, if $a \\sim t^\\b$, \\beq r \\sim \\dot a ^{-1} \\sim t^{1-\\b} , \\label{12} \\eeq so that $r$ is growing both in the matter-dominated ($\\b=2/3$) and in the radiation-dominated ($\\b=1/2$) era. Thus, as we go back in time, $r$ becomes smaller and smaller. If, for instance,  we wish to impose initial conditions at the Planck scale,  we must require a fine-tuning suppressing by $30$ orders of magnitude the spatial curvature term with respect to the other terms of the cosmological equations. Even if initial conditions are given at a lower scale (say the GUT scale) the amount of fine-tuning is still nearly as bad.  This problem can be solved by introducing an early phase during which the value of $r$, initially of order $1$, decreases so much in time that its subsequent growth  during FRW evolution keeps it still below $1$ today. It is evident that, on pure kinematic grounds, this requirement can be implemented in two classes of backgrounds.  \\begin{itemize} \\item[{\\bf (I)}] : $a \\sim t^\\b, ~\\b >1,  ~t \\ra +\\infty$. This class of background corresponds to what is conventionally called ``power inflation\" \\cite{Luc85}, describing accelerated expansion and decreasing curvature scale,  $\\dot a >0$, $\\ddot a >0$, $\\dot H <0$. It contains,  as the  limiting case ($\\b \\ra \\infty$), exponential de Sitter inflation,  $a \\sim e^{Ht}$, $\\dot H =0$, describing accelerated expansion with constant curvature. \\item[{\\bf(II)}] : $a \\sim (-t)^\\b,~ \\b <1, ~ t \\ra 0_-$. This case contains two subclasses. \\begin{itemize} \\item[{\\bf(IIa)}] : $\\b<0$, corresponding to ``superinflation\" or ``pole inflation\" \\cite{Sha84,Abb84,Kolb84}, and describing  accelerated expansion with growing curvature scale,  $\\dot a >0$, $\\ddot a >0$, $\\dot H >0$; \\item[{\\bf(IIb)}] : $0<\\b<1$, describing accelerated contraction and growing curvature scale \\cite{GasVe93b},   $\\dot a <0$, $\\ddot a <0$, $\\dot H <0$. \\end{itemize} \\end{itemize} In the first class of backgrounds, corresponding to post-big bang inflation, the Universe is driven {\\em away} from the singularity/high-curvature regime, while in the second class inflation drives the Universe {\\em towards} it, with the typical pre-big bang behaviour illustrated in Fig. \\ref{f11}.  We may thus immediately note  a very important ``phenomenological\" difference between post- and pre-big bang inflation. In the former case the Planck era  lies very far in the past, and its physics  remains screened from present observations, since the scales that probed Planckian physics are still far from re-entering. By contrast, in the pre-big bang case, the Planck/string regimes are closer to us (assuming that no or little inflation occurs after the big bang itself). Scales that probe Planckian physics are now the first to re-enter, and to leave an imprint for our observations (see, for instance, the case of a stochastic background of relic gravitational waves, discussed in Section \\ref{Sec5}).  The inflationary character of {Class IIa} backgrounds is well known, and recognized since the earlier studies of the inflationary scenario \\cite{Luc85}. The inflationary character of {Class IIb}  is more unconventional --a sort of ``inflation without inflation\" \\cite{Gas94a}, if we insist on looking at inflation as accelerated expansion--  and was first pointed out only much later \\cite{GasVe93b}. It is amusing to observe that, in the pre-big bang scenario, both subclasses {IIa} and  {IIb} occur. However, as discussed in detail in Subsection \\ref{Sec2.5}, they do not correspond to different models of pre-big bang inflation, but simply  to different kinematical representations of the  {\\em same} scenario in two {\\em different} conformal frames.  In order to illustrate this point, which is important also for  our subsequent arguments, we shall proceed in two steps. First we will show that, through a field redefinition $g=g(\\ti g,\\ti \\phi)$, $\\phi= \\phi (\\ti \\phi)$, it is always possible to move from the string frame (S-frame), in which the lowest order gravidilaton effective action takes the form \\beq S(g,\\phi)= - \\int d^{d+1}x \\sqrt{|g|}~ e^{-\\phi}\\left[R+ g^{\\mu\\nu}\\pa_\\mu\\phi \\pa_\\nu \\phi\\right], \\label{13} \\eeq to the Einstein frame (E-frame), in which the dilaton  $\\phi$ is minimally coupled to the metric and has a canonical  kinetic term: \\beq S(\\ti g,\\ti \\phi)= - \\int d^{d+1}x \\sqrt{|\\ti g|}~ \\left[\\ti R-{1\\over 2} \\ti g^{\\mu\\nu}\\pa_\\mu\\ti \\phi \\pa_\\nu \\ti \\phi\\right] \\label{14} \\eeq (see Subsection \\ref{Sec1.4} for notations and conventions). Secondly, we will show that, by applying such a redefinition, a superinflationary solution obtained in the S-frame becomes an accelerated contraction in the E-frame, and vice versa.  We shall consider, for simplicity, an isotropic, spatially flat background with $d$ spatial dimensions, and  set: \\beq g_{\\mu\\nu}= {\\rm diag} \\left( N^2, -a^2 \\da_{ij}\\right), ~~~~~~~ \\phi=\\phi(t), \\label{15} \\eeq where $g_{00}=N^2$ is to be fixed by an arbitrary gauge choice.  For this background  the S-frame action (\\ref{13}) becomes,  modulo a total derivative, \\beq S(g,\\phi)= - \\int d^{d+1}x {a^d e^{-\\phi}\\over N}\\left[\\dot \\phi^2-2dH\\dot\\phi +d(d-1)H^2 \\right], \\label{16} \\eeq where, as expected, $N$ has no kinetic term and plays the role of a Lagrange multiplier. In the E-frame the variables are $\\ti N, \\ti a, \\ti \\phi$, and the action(\\ref{14}), after integration by parts, takes the canonical form \\beq S(\\ti g, \\ti \\phi)= - \\int d^{d+1}x {\\ti a^d \\over \\ti N}\\left[-{1\\over 2}\\dot {\\ti \\phi}^2 +d(d-1)H^2\\right]. \\label{17} \\eeq A quick comparison with Eq. (\\ref{16}) finally leads  to the field redefinition (not a coordinate transformation!) connecting the Einstein and String frames: \\beq \\ti a = a e^{-\\phi/(d-1)}, ~~~~~~~~~ \\ti N = N e^{-\\phi/(d-1)}, ~~~~~~~~~ \\ti \\phi =\\phi \\sqrt{2\\over d-1}. \\label {18} \\eeq  Consider now  an isotropic, $d$-dimensional vacuum solution of the action (\\ref{16}), describing a superinflationary, pre-big bang expansion driven by the dilaton (see Section \\ref{Sec2}) \\cite{Mul90,GV91}, of Class IIa, with  $\\dot a >0, \\ddot a >0, \\dot H>0, \\dot \\phi >0$: \\beq a= (-t)^{-1/\\sqrt{d}}, ~~~~~~~~~ e^\\phi=  (-t)^{-(\\sqrt{d}+1)}, ~~~~~~~~~ t<0, ~~~~~~~~~~t \\ra 0_-, \\label{19} \\eeq and look for the corresponding E-frame solution. Since the above solution is valid in the synchronous gauge, $N=1$,  we can choose, for instance, the synchronous gauge also in the E-frame, and  fix $\\ti N$ by the condition: \\beq \\ti N dt \\equiv N e^{-\\phi/(d-1)} dt = d \\ti t, \\label{110} \\eeq which defines the E-frame cosmic time $\\ti t$ as: \\beq d \\ti t = e^{-\\phi/(d-1)} dt . \\label{111} \\eeq After integration \\beq t \\sim \\ti t ^{d-1\\over d +\\sqrt{d}}; \\label{112} \\eeq the transformed solution takes the form: \\beq \\ti a= (-\\ti t)^{1/{d}}, ~~~~~~~~~ e^{\\ti \\phi}=  (-\\ti t)^{-\\sqrt{2(d-1)\\over d}}, ~~~~~~~~~ \\ti t<0, ~~~~~~~~~ \\ti t \\ra 0_- . \\label{113} \\eeq It can easily be checked that this solution describes accelerated contraction of {Class IIb}, with growing dilaton and growing curvature scale: \\beq {d\\ti a\\over d \\ti t}<0,~~~~~~~ {d^2\\ti a\\over d \\ti t^2}<0,~~~~~~~ {d\\ti H\\over d \\ti t}<0,~~~~~~~ {d\\ti \\phi\\over d \\ti t}>0. \\label{114} \\eeq The same result applies if we transform other isotropic solutions from the String to the Einstein frame, for instance the superinflationary solutions with perfect fluid sources \\cite{GasVe93b}, presented in Section \\ref{Sec2}.  To conclude this section, and for later use, let us stress that the main dynamical difference between post-big bang inflation, Class I metrics,   and pre-big bang inflation, {Class II} metrics,  can also be conveniently illustrated in terms of the  proper size of the event horizon, relative to a given comoving observer.  Consider in fact the proper distance $d_e(t)$ of the event horizon from a comoving observer, at rest in an isotropic, conformally flat background \\cite{Rin56}: \\beq d_e(t)= a(t)\\int _t^{t_M} dt' a^{-1} (t'). \\label{115} \\eeq Here $t_M$ is the maximal allowed extension, towards the future, of the cosmic time coordinate for the given background manifold. The above integral converges for all the above classes of accelerated (expanding or contracting) scale factors. In the case of {Class I} metrics we have, in particular, \\beq d_e(t)= t^\\beta\\int _t^{\\infty} dt' t'^{-\\b}= {t\\over \\b-1} = {\\b\\over \\b-1} H^{-1}(t) \\label{116} \\eeq for power-law inflation ($\\b >1, t>0$), and \\beq d_e(t)= e^{Ht}\\int _t^{\\infty} dt' e^{-Ht'}=  H^{-1} \\label{117} \\eeq for de Sitter inflation. For {\\sl Class II} metrics ($\\b <1, t<0$) we have instead \\beq d_e(t)= (-t)^\\beta\\int _t^{0} dt' (-t')^{-\\b}= {(-t)\\over 1-\\b} ={\\b\\over \\b-1} H^{-1}(t). \\label{118} \\eeq In all cases the proper size $d_e(t)$ evolves in time like the so-called Hubble horizon (i.e. the inverse of the modulus of the Hubble parameter), and then like the inverse of the curvature scale. The size of the  horizon is thus constant or growing in standard inflation ({Class I}), decreasing in pre-big bang inflation ({Class II}), both in the S-frame and in the E-frame.  \\begin{figure}[t] \\centerline{\\epsfig{file=f12.ps,width=72mm}} \\vskip 5mm \\caption{\\sl Qualitative evolution of the Hubble horizon (dashed curve) and of the scale factor (solid curve) in the standard, post-big bang inflationary scenario.} \\label{f12} \\end{figure}  \\begin{figure}[t] \\centerline{\\epsfig{file=f13.ps,width=72mm}} \\vskip 5mm \\caption{\\sl Qualitative evolution of the Hubble horizon (dashed curve) and of the scale factor (solid curve) in the pre-big bang inflationary scenario, in the S-frame, $a(t)$, and in the E-frame, $\\ti a(t)$.} \\label{f13} \\end{figure}  Such an important difference is clearly illustrated in Figs.  \\ref{f12} and \\ref{f13}, where the dashed lines represent the evolution of the horizon and the solid curves the evolution of the scale factor.  The shaded area at time $t_0$ represents the portion of Universe inside our present Hubble radius. As we go back in time, according to the standard scenario, the horizon shrinks linearly ($H^{-1} \\sim t$); however, the decrease of the scale factor is slower, so that, at the beginning of the phase of standard evolution ($t=t_1$), we end up with a causal horizon much smaller than the  portion of Universe that we presently observe. This is the ``horizon problem\" already mentioned at the beginning of this section.  In Fig. \\ref{f12} the phase of standard evolution is preceded in  time by a phase of standard, post-big bang (in particular de Sitter) inflation. Going back in time, for $t<t_1$, the scale factor keeps shrinking, and our portion of  Universe ``re-enters\" the Hubble radius during a phase of constant (or slightly growing in time) horizon.  In Fig. \\ref{f13} the standard evolution is preceded in time by a phase of pre-big bang inflation, with growing curvature. The  Universe ``re-enters\"  the Hubble radius during a phase of shrinking  horizon. To emphasize the difference, we have plotted the evolution of the scale factor both as expanding in the S-frame, $a(t)$, and as contracting in the E-frame, $\\ti a (t)$. Unlike in post-big bang inflation, the proper size of the initial portion of the Universe may be very large in strings (or Planck) units, {\\em but not larger than the initial horizon} itself \\cite{Gas00}, as emphasized in the picture, and as will be discussed in a more quantitative way in Section \\ref{Sec2.4}. The initial horizon $H_i^{-1}$ is large because the initial curvature scale is small (in string cosmology, in particular,  $H_i \\ll 1/\\la_{\\rm s}$).  This is a basic consequence of the choice of the initial state; in the context of the string cosmology scenario, this approaches the flat, cold and empty string perturbative vacuum, (see the discussion of Section \\ref{Sec3}). This initial state has to be contrasted with the extremely curved, hot and dense initial state of the standard scenario, characterizing a  Universe that starts inflating at (or soon after) the Planck scale, $H_i \\sim 1/\\la_{\\rm P}$ (see also \\cite{Gas00a} for a more detailed comparison and discussion of pre-big bang  versus post-big bang inflation).  \\subsection{Outline, notations and conventions} \\label{Sec1.4}  We give here a general overlook at the material presented in the various sections of this report. Furthermore, each section will begin with an outline of the content of each of its subsections, and will try to be as self-contained as possible, in order to help the reader interested  only in some particular aspects of this field.  In Section {\\ref{Sec2}, after a very quick reminder of some relevant properties of superstring theory,  we review   the string-theoretic motivations behind the pre-big bang  scenario,   and outline the main ideas. In Section \\ref{Sec3}, after formulating on   the basis of the previous discussion a postulate of `` asymptotic past triviality\", we discuss, within that framework, the problem of initial conditions and fine-tuning.  As a preliminary to the discussion of the observational consequences of the pre-big bang scenario, Section \\ref{Sec4} presents some general results on the evolution of quantum fluctuations. Section \\ref{Sec5} will deal  with the specific case of tensor (metric) perturbations and Section \\ref{Sec6} with  scalar (dilatonic) ones, while Section \\ref{Sec7} will consider gauge (in particular electromagnetic) and axionic perturbations, and their possible physical relevance to galactic magnetic fields and large scale structure, respectively.   The last part of this report is devoted to various open problems, in particular to the  possibility of a smooth transition from the pre- to the post-big bang regime, the so-called ``exit problem\" (Section \\ref{Sec8}).   Several lines of approach to  (and partial solutions of) this problem are presented, including a discussion of heat and entropy production in pre-big bang cosmology.  A possible minisuperspace approach to the quantum regime  will be illustrated in Section \\ref{Sec9}. Finally,  Section \\ref{Sec10} contains the discussion of a possible dilatonic interpretation of the quintessence, a short presentation of other open problems, and  an outlook.  For further study of string cosmology and of the pre-big bang scenario, we  also refer the interested reader to other review papers \\cite{Lidsey00,Gas00a,Bar98,Fen00,Bran01}, as well as to two recent introductory lectures \\cite{Gas99a,Ve99a}. A regularly updated  collection of selected papers on the pre-big bang scenario is also available at \\cite{GasWeb}.  We finally report  here our conventions for the metric and the curvature tensor, together with the definitions of some variables frequently used in the paper.  We shall always use natural units $\\hbar=c=k_B=1$. Unless otherwise stated, the metric signature is fixed by $g_{00}>0$; the Riemann and Ricci tensors are defined by \\beq R_{\\mu\\nu\\a}\\,^\\b= \\pa_\\mu\\Ga_{\\nu\\a}\\,^\\b+ \\Ga_{\\mu\\r}\\,^\\b \\Ga_{\\nu\\a}\\,^\\r - (\\mu \\leftrightarrow \\nu), ~~~~~~ R_{\\nu\\a}= R_{\\mu\\nu\\a}\\,^\\mu. \\label{119} \\eeq In particular, for a Bianchi-I-type metric, and in the synchronous gauge, \\beq g_{\\mu\\nu}= {\\rm diag} \\left( 1, -a^2_i(t) \\da_{ij}\\right), \\label{120} \\eeq our conventions lead to \\bea && R_0\\,^0= - \\sum_i \\left(\\dot H_i +H_i^2\\right), ~~~~~~~~~ R_i\\,^j= - \\dot H_i\\da_i^j -H_i\\da_i^j \\sum_kH_k,\\nonumber \\\\ && R= -\\sum_i\\left( 2 \\dot H_i + H_i^2\\right)- \\left(\\sum_i H_i\\right)^2, \\label{121} \\eea where $H_i= d \\ln a_i/dt$. In a $D=d+1$ space-time manifold, Greek indices run from $0$ to $d$, while Latin indices run from $1$ to $d$.  The duality invariant dilatonic variable, the ``shifted dilaton\"  $\\fb$, is referred to a $d$-dimensional spatial section of finite volume, $(\\int d^dx\\sqrt{|g|})_{t={\\rm const}}<\\infty$, and is defined by \\beq e^{-\\fb}= \\int{ d^{d}x \\over \\la_{\\rm s}^d}\\sqrt{|g_d|} e^{-\\phi}. \\label{122} \\eeq In a Bianchi-I-type metric background, in particular, we shall absorb into $\\phi$ the constant shift $-\\ln(\\la_{\\rm s}^{-d}\\int  d^dx)$ (required to secure the scalar behaviour of $\\fb$ under coordinate reparametrizations), and we shall set \\beq \\fb= \\phi- \\sum_i \\ln a_i . \\label{123} \\eeq  Finally, $\\la_{\\rm s}$  is the fundamental length scale of string theory, related to the string mass $M_{\\rm s}$  and to the string tension $T$ (the mass per unit length) by \\beq \\la_{\\rm s}^2= M_{\\rm s}^{-2}T^{-1}\\equiv 2\\pi \\ap. \\label{124} \\eeq At the tree level (i.e. at lowest order) in the string coupling   $g_{\\rm s}$, the string length is related to the Planck length $\\la_{\\rm P}$,  and to the gravitational constant $G_D$ in $d+1$ dimensions, by \\beq 8\\pi G_D=\\la_{\\rm P}^{d-1}= \\la_{\\rm s} ^{d-1}e^{\\phi}. \\label{125} \\eeq In $d=3$, in particular, the relation between the string and the Planck mass $M_{\\rm P}=\\la_{\\rm P}^{-1}$ reads \\beq (\\la_{\\rm P}/\\la_{\\rm s})^2=(M_{\\rm s}/M_{\\rm P})^2= e^{\\phi}. \\label{126} \\eeq We shall often work in units such that $2 \\la_{\\rm s} ^{d-1}=1$,  i.e. $16 \\pi G_D=1$, in which $\\exp (\\phi)$ parametrizes, in the String frame, the (dimensionless) strength of the gravitational coupling. ",
        "conclusions": "\\label{Sec10} \\setcounter{equation}{0} \\setcounter{figure}{0}  In this last section we will conclude our report by discussing a possible ``late-time\" consequence of the pre-big bang scenario --the large-scale dominance of the dilatonic dark energy density-- and by presenting a list of various problems and aspects of such a scenario that are not included in the previous sections. Finally, we will summarize our personal outlook of string cosmology, together with some speculations about its possible future perspectives.  \\subsection{Towards the future: a dilaton-dominated Universe?} \\label{Sec10.1}  The dilaton is, undoubtedly, one of the most important ingredients  of the pre-big bang scenario illustrated in this paper. Indeed, the dilaton implements the duality symmetry, controls the strength of all interactions, sustains the initial inflationary evolution, contributes to the amplification of the quantum fluctuations (and, in particular, to the production of seeds for the magnetic fields), leading eventually to the formation of a cosmic background of massive scalar particles. All such effects are typical of string cosmology and represent the ``imprint\" of the pre-big bang scenario with respect to other, more conventional, inflationary scenarios.  Not satisfied with all such effects, however, the dilaton could still be hale and hearty, and still in action even today, so that it would affect in a determinant way not only the very early cosmological past, but also the present (and, possibly, future) state of our Universe. The dilaton potential energy, or a mixture of kinetic and potential energy density, could represent in fact the dark component responsible for the cosmic acceleration observed very recently \\cite{Riess98,Perlmutter99}. In this sense, string theory can automatically provide, with the dilaton, a ``non-minimal\" model of quintessence \\cite{Gas01,GPV01} (i.e. a model of quintessence based on a scalar field non-minimally coupled to gravity and to elementary matter fields).  There are, basically, two possible scenarios, depending on the (currently uncertain) shape of the non-perturbative dilaton potential $V(\\phi)$. In the context of superstring models for grand-unified theories (GUTs), the present value of the dilaton should determine in fact the whole set of gravitational and gauge coupling parameters, which are today constant (or, if they are time-dependent, are nevertheless running very slowly on a cosmological time scale). This constraint can be implemented in two ways.  The first possibility is a dilaton almost frozen at a minimum of the potential, in the moderate-coupling regime with $-\\phi$ of order $1$, in such a way that the GUT gauge coupling is fixed to \\cite{Kap85} \\beq \\a_{\\rm GUT} \\sim (M_{\\rm s}/M_{\\rm P})^2 \\sim e^\\phi \\sim 0.1 - 0.01 . \\label{101} \\eeq The second possibility is a dilaton monotonically running towards $+\\infty$, and the couplings saturated at small values in the strong-coupling regime because of the large number $N$ of fields entering the loop corrections, or the large value of the quadratic Casimir $C$ for unification gauge groups like $E_8$. Realistic values of $\\a_{\\rm GUT}$ and $M_{\\rm s}$ can be obtained at $\\phi \\gg 1$, for $N \\sim C \\sim 10^2$, typically as \\cite{Ven98,Ven01}: \\beq \\a_{\\rm GUT} \\sim {e^\\phi \\over 1+ C e^\\phi}, ~~~~~~~~~~~~~~~~~ {M_{\\rm s}\\over M_{\\rm P}} \\sim {e^\\phi \\over 1+ N e^\\phi}. \\label{102} \\eeq  In both cases, if the present value $V_0\\equiv V(\\phi_0)$ of the potential is of the order of the present Hubble scale $H_0^2$ (and, possibly, $\\dot \\phi^2 \\sim V_0$), the dilaton can reproduce the observed ``dark-energy\" effects, playing the role of the so-called ``quintessence\" \\cite{CalDaSte98}. The required fine-tuning of the amplitude of the potential (and the corresponding ultra-light value of the dilaton mass,  $m \\sim H_0$), seems to be unavoidable. In this context, however, it seems possible to alleviate the problem of the ``cosmic coincidence\" \\cite{Stein97}, generated by the observed (approximate) equality  $V_0 \\sim \\r_0$, where $\\r_0$ is the present energy density of the dominant dark-matter component. Let us explain how this could happen, in both  the weak-coupling and the strong-coupling scenarios of dilatonic quintessence.  In the first case we note that, even assuming that the dilaton gets frozen at a value $\\phi_0$ after the transition to the post-big bang regime, with a potential energy that is initially negligible ($V_0 \\ll H_{\\rm eq}^2$), it may keep frozen for the whole duration of the radiation epoch, but it tends to be shifted away from equilibrium as soon as the Universe enters the matter-dominated regime.  Suppose in fact that we  add a potential $V(\\phi)$ to the low-energy gravidilaton action: \\beq S= -{1\\over 2\\la_{\\rm s}^{2}} \\int d^{4}x \\sqrt{|g|}~ e^{-\\phi}\\left[R+ \\left(\\nabla \\phi\\right)^2+V(\\phi) \\right] +S_m, \\label{103} \\eeq where $S_m$ describes, for simplicity, perfect-fluid sources minimally coupled to the background. For a homogeneous and conformally flat metric, the dilaton equation can then be written in the form \\beq \\fpp+3H\\fpu -\\fpu^2 +{1\\over 2} e^\\phi (\\r-3p) +V'+V=0. \\label{104} \\eeq Combining this equation with the standard conservation equation of the matter sources, \\beq \\dot \\r +3H(\\r +p)=0, \\label{105} \\eeq it follows that constant and stable solutions $\\phi=\\phi_0$ (with $\\fpu=0=\\fpp$) are allowed in three cases only: 1) vacuum, $\\r=p=0$, $V+V'=0$;  2)  cosmological constant, $\\r=-p=\\r_0=$ const, $V+V'=-2e^{\\phi_0}\\r_0=$ const;  3) radiation, $\\r=3p$,  $V+V'=0$.  So, even if the dilaton is ``sleeping\" in the radiation era, it ``wakes up\" and starts rolling away from the freezing position determined by $V+V'=0$ just after the equilibrium epoch. We are thus led to the following question: For which values of the potential $V_0$ may the dilaton bounce back to the minimum, and the matter era  be followed by a quintessential, potential-dominated epoch?  If the answer would point {\\em only and precisely} at $V_0 \\sim H_0^2$, then the cosmic coincidence would be explained. If the answer would indicate for $V_0$ a restricted range of values, including $H_0^2$, the coincidence problem would remain, but it would be alleviated.  The answer to the above question depends on the shape of the potential, and on the effective strength of the dilaton coupling to macroscopic matter. To illustrate this point we may consider the E-frame cosmological equations (\\ref{62}), obtained from the action (\\ref{103}) through the conformal transformation \\beq \\ti g_{\\mu\\nu} =g_{\\mu\\nu} e^{-\\phi}, ~~~~~ \\ti \\phi= \\phi, ~~~~~ \\ti V = e^\\phi V, ~~~~~ \\ti \\r = \\r e^{2\\phi}, ~~~~~ \\ti p =p e^{2\\phi}. \\label{106} \\eeq In units of $16 \\pi G =1$ (and omitting the tilde, for simplicity), the dilaton equation becomes \\beq \\ddot {\\phi}+3H \\dot {\\phi} +{1\\over 2} \\alpha(\\phi)(\\r-3p) +{V}'=0 , \\label{107} \\eeq where we have taken into account, through the coupling function $\\a(\\phi)$, a possible loop renormalization of the effective dilaton coupling in the matter action. In the radiation era, $\\r=3p$ and a stable, frozen dilaton ($\\fpu=0$) thus corresponds to an extremum of the E-frame potential, $V'=0$. In the matter era, when $p=0$, there is a dilaton acceleration away from the minimum, $\\fpp =- \\a \\r/2$, possibly contrasted by the restoring force $-V'$.  Consider now a typical, supersymmetry-breaking dilaton potential, instantonically suppressed \\cite{BG86} at $\\phi \\ra -\\infty$, with a non-trivial structure developing a minimum in the region of moderate coupling, and exponentially growing at $\\phi \\ra +\\infty$ because of the conformal transformation to the E-frame. A ``minimal\" example of such a potential can be simply parametrized (in the E-frame) as follows \\cite{KaOl93} \\beq V= m^2 \\left[e^{k_1(\\phi_-\\phi_1)}+ \\beta e^{-k_2(\\phi_-\\phi_1)}\\right] e^{-\\ep \\exp\\left[-\\gamma (\\phi_-\\phi_1)\\right]}, \\label{108} \\eeq where $k_1,k_2, \\phi_1,\\ep ,\\b,\\ga$ are dimensionless numbers of order $1$ For an illustrative purpose we will choose here the particular values $k_1=k_2=\\b=\\ga=1$, $ \\ep=0.1$, $\\phi_1=-3$, in such a way that the minimum is at $\\phi_0 =-3.112...$, and that $ g_{\\rm s}^2 = \\exp (\\phi_0)\\simeq 0.045$, in agreement with Eq. (\\ref{101}).  With the above choice of parameters, the dilaton potential is plotted in Fig. \\ref{f101} for different values of $m$. It is important to note that the extremum $\\phi_0$ {\\em is not} separated from the perturbative regime  ($\\phi \\ra -\\infty$) by an infinite potential barrier, and that the lower  $V_0 \\sim m^2$, the lower  the barrier, the weaker is the restoring force $-V'$ and the easier it is  for the dilaton  to escape from the minimum and run to $-\\infty$.  \\begin{figure}[t] \\centerline{\\epsfig{file=f101.ps,width=72mm}} \\vskip 5mm \\caption{\\sl Plot of $V(\\phi)$ from Eq. (\\ref{108}), with $k_1=k_2=\\b=\\ga=1$, $ \\ep=0.1$, $\\phi_1=-3$. The three curves, from top to bottom, correspond respectively to $m=1/10, 1/12, 1/15$ in units of $M_{\\rm P}^2=2$.} \\label{f101} \\end{figure}  To match present phenomenology,  the value of $V_0$, in such a context, has to be chosen very small in string units, $V_0 \\sim H_0^2 \\sim (10^{-33}~ {\\rm eV})^2$ (see below for a possible justification). It follows that the potential barrier is very low, and the dilaton would certainly escape from the minimum at the beginning of the matter era, {\\em unless} the coupling to matter ($\\a\\r/2$) is also correspondingly small. Such a coupling, on the other hand, {\\em has} to be very small, because the corresponding scalar force has a very long range, $V''(\\phi_0) \\sim m^2\\sim V_0\\sim H_0^2$, and the dilaton must be strongly decoupled ($\\a \\ll1$, at least today) to avoid unacceptable violations of the equivalence principle \\cite{DP94a,DP94b}.  It can be shown, as a consequence, that the present values of $\\a$ allowed by the gravitational phenomenology ($\\a_0 \\laq 10^{-3}$) are compatible with a late cosmological phase dominated by the dilaton potential only for a restricted range of $V_0$, which depends on the value of $\\a$ at the equilibrium epoch, $\\a_{\\rm eq}$ \\cite{Gas01}. Expanding $\\a(\\phi)$ around the minimum, and starting from $\\a_{\\rm eq}=0.1$, for instance, a numerical analysis shows that the dilaton, after a small shift at the equilibrium epoch, bounces back to the minimum provided $10^{-7} H_{\\rm eq} \\laq m \\laq H_{\\rm eq}$, which includes the realistic case $m \\sim V_0^{1/2} \\sim H_0\\sim 10^{-6} H_{\\rm eq}$. In such a context, therefore, the coincidence problem is not  strictly solved, but possibly alleviated (see \\cite{Gas01} for a more detailed discussion).  An alternative scenario for a dilatonic interpretation of the observed ``quintessential\" effects, is based on the assumption that the dilaton never gets  trapped in a minimum, and evolves monotonically (and boundlessly) from negative (pre-big bang) to positive (post-big bang) values, with a potential smoothly approaching zero as $\\phi \\ra +\\infty$ (see Fig. \\ref{f21}). The non-perturbative potential, in this context, is typically characterized by a bell-like shape,  which can be parametrized (in the S-frame) as \\cite{GPV01}: \\beq V=m^2 \\left[ e^{-{1\\over \\b_1}\\exp (-\\phi)}- e^{-{1\\over \\b_2}\\exp (-\\phi)}\\right], \\label{109} \\eeq where $\\b_1>\\b_2>0$ are dimensionless numbers of order unity and, again, the mass scale $m$ satisfies $m \\ll H_{\\rm eq}$, to avoid that the dilaton becomes dominant too early.  As before, the potential  is instantonically suppressed at $\\phi \\ra -\\infty$, but now is also esponentially suppressed in the $\\phi \\ra +\\infty$ limit. The mass scale $m$, in the context of a non-perturbative potential, is naturally related to the string scale by $m \\sim \\exp\\left[-\\b \\a^{-1}_{\\rm GUT}\\right] M_{\\rm s}$, where $\\b$ is a model-dependent parameter of order $1$. This may explain the smallness of the potential in string units, but it should not hide the fact that the constant parameter $\\b$ has to be precisely adjusted  {\\em ad hoc} if we want to start the accelerated dilaton-dominated phase not much  earlier than at redshifts of order one \\cite{Turner01}, and not later than today.  In the context of a running-dilaton scenario, the gauge couplings and the ratio $M_{\\rm s}/M_{\\rm P}$ are to be stabilized, in the $\\phi \\ra +\\infty$ limit, by the loop corrections. Such corrections, on the other hand, also play a fundamental role in determining the effective dilatonic charge of the elementary fields appearing in the matter action. Thanks to this interplay, a dilatonic charge that switches on at late times, in the (non-baryonic) dark-matter sector, seems able to provide a possible dynamical explanation of the cosmic coincidence \\cite{GPV01} (see \\cite{DaPiaVe02,DaPiaVe02a} for possible related violations of free-fall universality, and time variation of the natural constants, induced in this context).  Such a dynamical approach to the coincidence problem is based, in particular, on the following loop-corrected action \\beq S = -{1\\over 2\\la_{\\rm s}^2} \\int d^{4}x \\sqrt{-  g}~ \\left[e^{-\\psi(\\phi)} R+ Z(\\phi) \\left(\\nabla \\phi\\right)^2 + {2\\la_{\\rm s}^2} V(\\phi)\\right] +S_m(\\phi, {\\rm matter}), \\label{1010} \\eeq where $V$ is the potential of Eq. (\\ref{109}), and the loop form factors $\\psi$ and $Z$, in a minimal ``induced-gravity\" scenario, are assumed to be parametrized at large $\\phi$ as follows \\cite{Ven98,Ven01}: \\beq e^{-\\psi(\\phi)}\\, = \\,  e^{-\\phi} + c_1^2 \\, , ~~~~~~~~~~~ Z(\\phi)\\, = \\, e^{-\\phi} - c_2^2 , \\label{1011} \\eeq where $c_1^2 \\sim c_2^2 \\sim 10^2$, in agreement with Eq. (\\ref{102}). We also assume that the action $S_m$ contains non-baryonic dark-matter whose dilatonic charge density per unit of gravitational mass, $\\sg_d/\\r_d$, switches on as $\\phi \\ra +\\infty$, and can be parametrized as \\cite{GPV01}: \\beq q(\\phi)= {\\sg_d\\over \\r_d}= -{2\\over \\sqrt{-g} \\r_d} {\\da S_m\\over \\da \\phi} =q_0 {e^{q_0 \\phi}\\over c^2+ e^{q_0\\phi}}, \\label{1012} \\eeq where, again, $c^2 \\sim 10^2$. Note that we are using here a definition of  dilatonic charge density different from the one previously introduced in Eqs. (\\ref{616}) and (\\ref{840}).  The analytical and numerical study of such a model shows that, in a way that is largely independent from the initial conditions, the Universe is eventually driven to an asymptotic accelerated regime characterized by a frozen ratio between the dark-matter and the dilatonic energy density, $\\r_d/\\r_\\phi=$ const. The present approximate equality of $\\r_d$ and $V_0$ is no longer a coincidence, in this context, but a dynamical property of the asymptotic configuration, as in models of ``coupled quintessence\" \\cite{Wett00,AmenToc01}.  For a better illustration of the ``late-time\" post-big bang scenario, described by Eqs. (\\ref{1010})--(\\ref{1012}), it is convenient to consider the E-frame cosmological equations for the tilted variables defined through the conformal transformation \\beq a = c_1 \\ti a e^{\\psi/2}, ~~~~~ d t = c_1 d \\ti t e^{\\psi/2}, ~~~~~ \\ti \\r = c_1^2  \\r e^{2\\psi} , ~~~~~ \\ti p = c_1^2  p e^{2\\psi} , ~~~~~ \\ti \\sg = c_1^2 \\sg  e^{2\\psi} \\label{1013} \\eeq (we are considering, as before, perfect-fluid sources evolving in a conformally flat metric background). Note that  $c_1^2$, according to Eq. (\\ref{1011}), controls the asymptotic value of the ratio between the string and the Planck scale, $c_1^2\\la_{\\rm s}^{-2}=\\la_{\\rm P}^{-2}$. The E-frame equations thus become (omitting the tilde, and in units $2 \\la_{\\rm P}^2=16\\pi G=1$) \\cite{GPV01} \\bea && 6H^2 = \\r +\\r_\\phi, \\label{1014} \\\\ && 4 \\dot H + 6H^2 =-p -p_\\phi, \\label{1015} \\\\ && k^2(\\phi) \\left(\\ddot{\\phi}+3 H \\dot{\\phi}\\right) + k(\\phi)\\, k'(\\phi)\\, \\dot{\\phi}^2 + \\hat{V}'(\\phi) + \\frac{1}{2}\\left[{\\psi'(\\phi)} (\\rho - 3 p) + q\\r \\right] = 0, \\label{1016} \\eea where a dot denotes differentiation with respect to the E-frame cosmic time, and we have defined: \\bea && k^2(\\phi) = 3 \\psi^{\\prime 2} - 2 {e^\\psi} Z , ~~~~~~~~~~~~~~~~\\hat V = c_1^4 e^{2\\psi} V , \\label{1017}\\\\ && \\r_\\phi= {1\\over2} k^2(\\phi) \\dot \\phi^2 +\\hat V(\\phi), ~~~~~~~~~~~~ p_\\phi= {1\\over2} k^2(\\phi) \\dot \\phi^2 -\\hat V(\\phi). \\label {1018} \\eea The source terms $\\r,p,\\sg=q\\r$ generically include the contribution of radiation, baryonic and non-baryonic matter components, but only for the non-baryonic dark  component $\\r_d$ is the dilaton charge  assumed to be significantly different from zero, and parametrized as in Eq. (\\ref{1012}). By analysing the time evolution of the (total) dilaton energy density, $\\r_\\phi$, we then find that the post-big bang solutions are generally characterized by three dynamical phases \\cite{GPV01}.  Suppose in fact that the radiation-dominated, post-big bang regime starts with with a negligible dilaton potential,  a negligible dark-matter energy density $\\r_d \\sim a^{-3}$, and  a fully kinetic dilaton energy density ,which is rapidly diluted as $\\r_\\phi \\sim a^{-6}$. When $\\r_\\phi$ falls below $\\r_d$ one finds that the Universe enters a first ``focusing\" phase with $\\r_\\phi \\sim a^{-2}$, so that the dilaton kinetic energy starts growing with respect to radiation, and converges towards the other energy components.  When the matter becomes dominant,  the Universe enters a subsequent ``dragging\" phase, where $\\r_\\phi$ and $\\r_d$ evolve in time with the same behaviour, together with the baryonic matter density (uncoupled to the dilaton), $\\r_b \\sim a^{-3}$.  Actually,  $\\r_\\phi$ and $\\r_d$ are diluted a little bit faster (like $ a^{-3 -\\ep^2}$) than $\\r_b$, but the deviation is controlled by a parameter, which is constant and very small during the dragging phase, \\beq \\left[\\ep(\\phi)\\right]_{\\rm drag} \\equiv \\left[\\psi' + q\\over k \\right]_{\\rm drag} \\sim {q_0 c_1\\over c^2 c_2}\\sim {1\\over c^2} \\ll1. \\label{1019} \\eeq  At late times (but not later than today) the potential $V(\\phi)$ and the dilaton charge $q(\\phi)$ of dark-matter come eventually into play, and the Universe enters a ``freezing\" phase characterized by a constant ratio $\\r_\\phi/\\r_d$, and by accelerated evolution. In this asymptotic regime, $\\phi \\ra + \\infty$, the parameters of the model approach the constant  values \\beq k^2(\\phi) = 2c_2^2/c_1^2 \\equiv 2/\\lambda^2,  \\quad \\sg = \\sg_d,  \\quad \\r  = \\r_d,  \\quad q(\\phi) = q_0, \\label{1020} \\eeq and the E-frame equations can be written as \\bea && \\label{dark1} \\dot \\r_d +3H\\r_d -\\frac{q_0}{2}\\, \\r_d \\dot \\phi=0, \\quad \\quad \\quad \\dot \\r_\\phi +6H\\r_k +\\frac{q_0}{2}\\, \\r_d\\dot \\phi=0, \\label{1021}\\\\ && 1=\\Om_d+\\Om_k+\\Om_V, \\quad \\quad \\quad \\quad ~~~ 1+{2\\dot H\\over 3 H^2}=\\Om_V-\\Om_k, \\label{1022} \\eea where \\bea && \\r_d= 6H^2 \\Om_d, ~~~~~~~~~~~~~~~~~~~~~~\\r_\\phi=\\r_k+\\r_V,\\label{1023}\\\\ && \\r_k= 6H^2 \\Om_k= \\dot \\phi^2/ \\lambda^2, ~~~~~~~~~~ \\r_V= 6H^2 \\Om_V=\\hat V. ~~~~~~~~ \\label {1024} \\eea They can be solved by an asymptotic configuration with constant $\\Om_k$, $\\Om_V$ and $\\Om_d$ (acting as late-time attractor \\cite{Amen00}), where the dilaton fraction of critical energy density, $\\Om_k+\\Om_V$, the dilaton equation of state, $w_\\phi= (\\Om_k-\\Om_V)/(\\Om_k+\\Om_V)$, and the asymptotic acceleration parameter, are given by \\cite{GPV01} \\bea && \\label{1025} \\Omega_\\phi\\, =\\, \\frac{12 +  q_0(q_0+2)\\lambda ^2}{ (q_0+2)^2\\lambda ^2}, \\quad \\quad \\quad w_\\phi\\, = \\, - \\frac{ q_0 (q_0+2)\\lambda^2}{12 +  q_0(q_0+2)\\lambda^2}, \\\\ && {\\ddot a \\over aH^2}= 1+{\\dot H\\over H^2}= {q_0-1\\over q_0+2}. \\label{1026} \\eea  Such an asymptotic configuration is illustrated in Fig. \\ref{f102}, where we have plotted various curves at $\\Om_\\phi=$ const and $w_\\phi=$ const in the $\\{q_0,\\la\\}$ plane. As shown in the figure, a positive acceleration ($q_0>1$) is perfectly compatible with the range of $\\Om_\\phi$ and $w_\\phi$ allowed by present phenomenological constraints \\cite{Wang00,Balbi01}, namely $0.6 \\laq \\Om_\\phi \\laq 0.7$ and $-1 \\leq w_\\phi \\laq -0.4$.  \\begin{figure}[t] \\centerline{\\epsfig{file=f102.ps,width=72mm}} \\vskip 5mm \\caption{\\sl Asymptotic configurations in the plane  $\\{q_0,\\la\\}$, according to Eq. (\\ref{1025}). The full curves correspond to the constant values  $\\Om_5=0.5$, $\\Om_6=0.6$, $\\Om_7=0.7$, $\\Om_8=0.8$, and the dashed curves to the constant values of $w_\\phi$ indicated in the figure.} \\label{f102} \\end{figure}  We note, finally, that a simple integration of Eq. (\\ref{1026}) gives the late-time evolution of the dilaton and dark-matter energy density, \\beq \\r_\\phi \\sim \\r_d \\sim H^2 \\sim a^{-6/(2+q_0)}. \\label{1027} \\eeq The baryon dark matter, uncoupled to the dilaton ($q=0, \\r_b \\sim a^{-3}$), is rapidly diluted with respect to $\\r_d$, and we thus obtain, in the context of such a dilatonic scenario, also a possible explanation of the present smallness of the ratio $\\r_b/\\r_d$ \\cite{GPV01} (see also \\cite{AmenToc01a}).  \\subsection{Other open problems} \\label{Sec10.2}  In this report we have focalized our attention on many important aspects of the pre-big bang scenario (in particular, on those mainly studied in the past years), without pretending, of course, to have presented an exhaustive discussion of all possible problems and aspects of such scenario. There are important topics that we have left almost untouched, also because too few results ae available on these in the present literature. Thus, it seems to us appropriate to conclude with a list (and a brief discussion) of such topics, which we hope will be the object of more intensive studies in the next years.  In particular, we have given very little room to a discussion, in the pre-big bang context, of a realistic model of (dynamical?) dimensional reduction of the superstring manifolds, the (possible) compactification and the associated stabilization of the internal dimensions. The best approach to this problem, at present, is probably the model of \\cite{BraVa89}, generalized to a brane-gas in \\cite{ABE00,BEK01}, which assumes, however, a toroidal structure of the space-time manifold from the very beginning (as in standard Kaluza--Klein compactification). There are no results, at present, concerning the possible evolution of the pre-big bang phase into a multi-dimensional non-compact structure, eventually characterized by a warped geometry able to confine long-range interactions on a four-dimensional brane \\cite{RS2}.  Also, we have not  discussed in detail the matching of the pre-big bang regime to the subsequent Friedmann--Robertson--Walker (FRW) phase, the possibility in this context of background oscillations (of the metric and of the dilaton), with related ``preheating\" and reheating effects,  and the associated particle production (for instance gravitino production, which is a potential problem  \\cite{Wei82} for any inflationary model based on a supersymmetric theory). All these problems, including dimensional reduction, are (more or less directly) related to the fact that there is at present no compelling model of a complete and realistic exit transition (see Section \\ref{Sec8}).  It is worth mentioning, in this respect, that the problem of reheating and of possible dangerous relics (moduli and gravitinos) in pre-big bang cosmology has recently been faced in \\cite{Buo00}. The main difference from the reheating scenario of standard inflationary models is that, for pre-big bang inflation, the matching to the FRW phase may occur in principle at higher-curvature scales (typically, $H_1 \\sim M_{\\rm s} \\sim 10^{17}$ GeV), where gravitinos and moduli fields should also be copiously produced in scattering processes. Another important difference is that the particles present at the beginning of the FRW phase are produced (with various spectra) directly from the amplification of the quantum fluctuations of the background fields during inflation (as stressed in Subsection \\ref{Sec8.4}), with typical densities $\\r_i(t) \\sim N_i H_1^4 (a_1/a)^4$, where $N_i$ is the number of helicity states in species $i$.  In such a context, the thermalization scale of scalar and vector particles, charged under the group of the observable sector, has been estimated in \\cite{Buo00} by assuming that the dilaton and the internal dimensions are already frozen at the beginning of the FRW phase, when $H=H_1$ and $g=g_1 \\equiv M_{\\rm s} /M_{\\rm P}(t_1)$. The charged particles, with typical energy $\\om\\sim H_1 a_1/a$ and number density $n_r \\sim \\r_r/\\om$, where \\beq \\r_r = \\Om_r \\r_c \\sim {N_r\\over N_{\\rm tot}} {H_1^2 M_{\\rm P}^2 a_1^4\\over a^4} = {N_r\\over N_{\\rm tot}} \\left(H_1M_{\\rm s}\\over g_1\\right)^2 \\left( a_1\\over a\\right)^4 \\label{1028} \\eeq ($ N_{\\rm tot}$ is the total number of degrees of freedom present at $t=t_1$), interact with cross section $\\sg \\sim \\a^2 /\\om^2$, where $\\a \\sim g_1^2$. The thermalization scale $H_{\\rm th} \\sim n_r \\sg$ is then given by \\beq H_{\\rm th} \\sim g_1^4 \\left(N_r \\over  N_{\\rm tot}\\right)^2 {M_{\\rm s}^4 \\over H_1^3}. \\label{1029} \\eeq  If the charged particles are already dominat at $t=t_1$, i.e. $N_r \\sim N_{\\rm tot}$, $\\Om_r \\sim 1$, and if $H_1 \\sim M_{\\rm s}$ as in minimal pre-big bang models, then the post-big bang Universe thermalizes at a scale $H_{\\rm th} \\sim g_1^4 H_1\\sim (10^{-4}$ -- $10^{-8})~ M_{\\rm s}$, with a corresponding reheating temperature $T_{\\rm r} \\sim (H_{\\rm th} M_{\\rm P})^{1/2} \\sim g_1^{5/2} M_{\\rm P}\\sim g_1^{3/2}M_{\\rm s}$. If, on the contrary, $\\Om_r <1$, then reheating is only achieved once the gauge-singlet fields, which carry the remainder of the total energy density, have decayed into radiation.  Such a reheating process generates an entropy density $s \\sim \\r_r/T_r$, which is not sufficient, however, to dilute the unwanted relic particles produced gravitationally during inflation, and also in scatterings of the thermal bath for $t>t_1$. For the moduli, in particular, the number-density-to-entropy-density ratio $Y_m =n_m/s_m$ turns out to be \\cite{Buo00} $Y_m \\sim 0.3 /g_\\ast$, where $g_\\ast$ is the number of degrees of freedom in the radiation bath after thermalization.  For $g_\\ast \\sim 10^2$ -- $10^3$ this is well above the limit imposed by nucleosynthesis \\cite{Wei82}, which implies $Y \\laq 10^{-13}$.  The same is true for gravitinos. Indeed, even if they are not gravitationally produced during pre-big bang inflation because they keep effectively massless \\cite{BruHa00}, gravitinos are produced in scatterings of the high-temperature thermal bath, and their number to entropy ratio can be estimated (for $H_1 \\sim M_{\\rm s}$) as \\cite{Buo00} $Y_{3/2} \\sim g_1^{5/2} g_\\ast ^{-7/4}$, which is still in excess of the nucleosynthesis limit (see also \\cite{Lemo99,Tsu01} for gravitino production in a string cosmology context).  The dilution of the unwanted relics possibly present at the beginning of the post-big bang phase thus requires a strong entropy release, at the level of about ten orders of magnitude for gravitationally produced relics in the context of minimal pre-big bang models (so as to suppress $Y$ from $10^{-3}$ to $10^{-13}$). Such an entropy production may be due to a secondary reheating phase, as already discussed in Subsection \\ref{Sec6.3}, and a typical possibility corresponds to the oscillations and decay of a scalar (or pseudoscalar) field, which gets a mass in the post-big bang epoch through a symmetry- (or supersymmetry-) breaking mechanism.  A natural string-cosmology candidate for this effect is the axion (as discussed in Subsection \\ref{Sec7.5}), or even the dilaton, if it is heavy enough. In particular, if the initial value of the oscillating field is  at least of order $M_{\\rm s}$, it is possible to dilute in this way moduli and gravitinos, but not monopoles (possibly produced by GUT symmetry breaking), while an initial value of order $M_{\\rm P}$ is marginally sufficient also for monopoles \\cite{Buo00}.  It is possible, however, that the monopole problem be independently solved by an efficient monopole--antimonopole annihilation \\cite{Cop88}. We should keep in mind, also, that in more complicated, non-minimal pre-big bang backgrounds (see Subsection \\ref{Sec5.3}), the energy distribution among the produced particles may be drastically altered, with a resulting easier dilution of the relics component, and a smaller required amount of entropy production (see \\cite{Buo00} for a detailed discussion).  We note, finally, that the entropy-production process is independently interesting by itself, in such a context, as it can naturally accommodate baryogenesis if the oscillating scalar is identified with the Affleck--Dine condensate, made of squarks and sleptons, whose decay generates the baryon asymmetry \\cite{AD85}. In particular, a ``mixed\" reheating phase with two oscillating fields (a modulus and a condensate), seems to be able to reproduce both the required dilution of relics and the right baryon asymmetry  in the context of pre-big bang cosmology \\cite{Buo00}. The oscillations of the Affleck--Dine condensate could be sufficient, by themselves, to solve the moduli and gravitinos problem in the context of non-minimal pre-big bang models, provided the resulting baryon asymmetry were kept small enough by a very small CP-violation parameter, or by a very efficient electroweak mechanism of baryon-number erasure \\cite{Gailla95}.  \\subsection{Outlook} \\label{Sec10.3}  Although it is always difficult to make forecasts, particularly in theoretical physics, we are at least confident of one thing: research on the implications of string/M-theory on fundamental cosmological questions is not just a momentary fashion. It is going to continue at an increasing rate as long as new cosmological data come in, and put more and more puzzles in front of theoretical cosmology. Indeed, the last few years have witnessed a sudden jump of interest in this field,  triggered partly by the new data on CMB anisotropy and on evidence for an accelerating Universe, partly by the new theoretical developments related to large extra dimensions, the possible lowering of the quantum-gravity scale, and the brane-world ideas.  One line of research has been devoted to the modifications of the standard cosmological equations for observers confined to a brane immersed in  a higher-dimensional space-time \\cite{Bine00}. Such modifications are in principle sufficiently drastic to make us worry about how to preserve the successes of the standard set-up, on the one hand, and  about how to use the deviations in order to improve the situation where the standard description may face difficulties  (e.g. during inflation \\cite{CoLiLi01} and early cosmology) on the other hand.  A second development, which is even more relevant to this report, is the proposal of new cosmologies \\cite{KOST1,KOSST} that share many properties with the pre-big bang scenario. As in the latter, they assume that the big bang singularity is spurious, and that the Universe (and time) had a long existence prior to it. Also, the initial state is assumed to be very perturbative, although not as generic as in our approach. The most  important difference, however, is that such models are constructed around the brane-world idea (see Subsection \\ref{Sec8.5}). Thus, although some of the pictures look very similar, their meaning is drastically different. As an example,  the plane-wave collision described  in Subsection \\ref{Sec3.4}  superficially resembles the brane collision of the ekpyrotic scenario; but while in the former our Universe lies in the bulk, in the model of \\cite{KOST1,KOSST} it lies  on one of the two colliding branes. Also, the big bang is the moment of collision in \\cite{KOST1,KOSST} while, in the pre-big bang scenario, it emerges sometime after the collision,  as a consequence of gravitational collapse.  The model in \\cite{KOSST} is even closer to the pre-big bang idea, in the sense that, there, the extra dimension represents  the dilaton  in the strong-coupling regime, according to the M-theory conjecture \\cite{Hor96}. It is thus possible to draw the bouncing-Universe scenario of \\cite{KOSST} in the same phase-space plane as was used in Fig. \\ref{f81}. The difference between the two scenarios is very simple: while the pre-big bang Universe starts its evolution in the upper-right quadrant of the diagram --and thus inflates, in the  S-frame, as it evolves from weak to strong coupling-- in the scenario of \\cite{KOSST} the Universe starts in the lower-right quadrant and thus contracts, even in the S-frame, while it evolves from strong to weak coupling.  We have seen  in Subsections \\ref{Sec1.3} and \\ref{Sec2.4} that an accelerated contraction is as  good as an accelerated expansion for solving the horizon and flatness problems, hence the model of \\cite{KOSST} appears a priori as good as the pre-big bang model. It is important, however,  to stress some differences, especially regarding the exit problem. If the exit is to  occur at weak coupling, it should be possible to study it within the tree-level effective action (although non-perturbatively in $\\alpha'$). This was precisely the framework described in Subsection \\ref{Sec8.2}, where it was shown that, at least order by order in the $\\alpha'$ expansion, a conserved quantity forbids a transition from contraction to expansion in the S-frame. The possibility remains, however,  of  non-perturbative (world-sheet instanton) effects for achieving the bounce.  Still at the theoretical level, both the ekpyrotic and the pre-big bang scenario suffer from the criticism of  \\cite{DH00/1} about the generic occurrence of a phase of chaotic oscillations. This is even more crucial for the models of \\cite{KOST1,KOSST}, since they need a highly fine-tuned initial state \\cite{KKL,KKLT} and, again,  cannot invoke string-loop corrections for stopping BKL oscillations.  At the more phenomenological level the main difference between the two classes of models concerns the spectrum of (adiabatic) scalar metric perturbations. While these have  been claimed to have naturally a blue spectrum in a pre-big bang context \\cite{BruMu95} (see also \\cite{Hwang98,Tsu02}), the authors of  \\cite{KOST1} have claimed to get naturally (although through a certain hypothesis on the shape of a potential) a scale-invariant spectrum, as favoured by observations. Without taking sides, we will mention that this claim has been challenged by several authors \\cite{Ly02,BF01} (see however \\cite{DuVe02}). On the other hand, the curvaton mechanism of \\cite{LyWa02} could still generate scale-invariant adiabatic perturbations even if these were not immediately generated during the pre-big bang (or pre-bounce) phase.  Since string/M-theory is proposed by its fans as a candidate unified theory of all phenomena, it simply cannot allow itself to avoid tackling the fundamental questions that classical and quantum gravity are proposing: What is the fate of the  classical singularities --that are so ubiquitous in general relativity-- in the context of a consistent quantum theory of gravity?  In this report the pre-big bang scenario was described in such a way that the reader's attention was drawn  to the big bang singularity of general relativity and to how   its  avoidance, in string theory, could open up new  possibilities  for solving the long-standing problems of classical, hot big bang cosmology. Amusingly, in doing so, we were led to connect this question to the other big puzzle of general relativity: the fate of black-hole singularities, with all its ramifications into the information paradox, the loss of quantum coherence, and the like.  To conclude, let us try to draw some generic lessons from the particular attempt at a new, string-based cosmology that we have illustrated in the previous sections.  \\begin{itemize} \\item[(1)] Neither our Universe, nor space and time themselves, have to emerge from a singularity: the singularities of classical gravitational theories should signal the lack of finite-size and quantum corrections, or the need for new degrees of freedom in the description of physics at very short distances. Also, the Universe did not have to start from the very beginning as a hot and dense ``soup\" of particles and radiation: a hot Universe can emerge from a cold one, thanks to the parametric amplification of the quantum fluctuations of the vacuum, during inflation.  \\item[(2)] Inflation does not need a scalar potential (or, more generally, an effective cosmological constant), and can  naturally appear as the result of the underlying duality symmetries of the cosmological background. Also, inflation can be represented as a contraction (in the appropriate frame), and the study of generic initial conditions can be related, mathematically, to the study of gravitational collapse in general relativity.  \\item[(3)] The Planck scale does not provide any fundamental impenetrable barrier, which would  limit our direct experimental information on what happened before. On the contrary, the phenomenological imprint of the pre-Planckian epoch can be encoded into a rich spectrum of observable relics, reaching us today directly out of the pre-big bang Universe. Pre-big bang physics can thus be the object of dedicated experimental searches, which will be able to tell us, hopefully in a not too far future, whether or not there are chances for present string-cosmology models to provide a successful description of our primordial Universe.  \\end{itemize}    \\vskip 5mm"
    },
    "0207/hep-ph0207342_arXiv.txt": {
        "abstract": "I review some theoretical aspects of neutrino oscillations in the case when more than two neutrino flavours are involved. These include: approximate analytic solutions for 3-flavour (3f) oscillations in matter; matter effects in $\\nu_\\mu \\leftrightarrow \\nu_\\tau$ oscillations; 3f effects in oscillations of solar, atmospheric, reactor and supernova neutrinos and in accelerator long-baseline experiments; CP and T violation in neutrino oscillations in vacuum and in matter; the problem of $U_{e3}$; 4f oscillations.  \\vspace{1pc} ",
        "introduction": "Explanation of the solar and atmospheric neutrino data in terms of neutrino oscillations requires at least three neutrino species, and in fact three neutrino species are known to exist -- $\\nu_e$, $\\nu_\\mu$ and $\\nu_\\tau$. If the LSND experiment is correct, then probably a fourth neutrino type should exist, a light sterile neutrino $\\nu_s$. However, until relatively recently most of the studies of neutrino oscillations were performed in the 2-flavour framework. There were essentially two reasons for that: (1) simplicity -- there are much fewer parameters in the 2-flavour case than in the 3-flavour one, and  the expressions for the transition probabilities are much simpler and by far more tractable, and (2) the hierarchy of $\\Delta m^2$ values, which allows to effectively decouple different oscillation channels. The 2-flavour approach proved to be a good first approximation, which is a consequence of the hierarchy $\\Delta m_\\odot \\ll \\Delta m_{\\rm atm}$ and of the smallness of the leptonic mixing parameter $|U_{e3}|$.  However, the increased accuracy of the available and especially forthcoming neutrino data makes it very important to take into account even relatively small effects in neutrino oscillations. In addition, the experimentally favoured solution of the solar neutrino problem is at present the LMA MSW one, which requires the hierarchy between $\\Delta m_\\odot$ and $\\Delta m_{\\rm atm}$ to be relatively mild. Also, effects specific to $\\ge 3$ flavour neutrino oscillations, such as CP and T violation, are now being very widely discussed. All this makes 3-flavour (or 4-flavour) analyses of neutrino oscillations mandatory.  In my talk I review some theoretical issues pertaining to neutrino oscillations in the case when more than two neutrino species are involved. I mainly concentrate on 3-flavour (3f) oscillations and only very briefly consider the 4f case. The topics that are discussed include: approximate analytic solutions for 3f oscillations in matter; matter effects in $\\nu_\\mu \\leftrightarrow \\nu_\\tau$ oscillations; 3f effects in oscillations of solar, atmospheric, reactor and supernova neutrinos and in accelerator long-baseline experiments; CP and T violation in neutrino oscillations in vacuum and in matter; the problem of $U_{e3}$; 4f oscillations. ",
        "conclusions": "3f effects in solar, atmospheric, reactor and supernova neutrino oscillations and in LBL accelerator neutrino experiments may be quite important. They can lead to up to $\\sim 10$\\% corrections to the oscillation probabilities and also to specific effects, absent in the 2f case. The manifestations of $\\ge 3$ flavours in neutrino oscillations include fundamental CP violation and T violation, matter-induced T violation, matter effects in $\\nu_\\mu \\leftrightarrow \\nu_\\tau$ oscillations, and specific CP- and T-conserving interference terms in oscillation probabilities. The leptonic mixing parameter $U_{e3}$ plays a very special role and its study is of great interest.  In the 4f case, large CP violation and (both fundamental and matter-induced) T violation effects are possible. However, 4f scenarios are strongly disfavoured by the data.  \\newpage"
    },
    "0207/astro-ph0207507_arXiv.txt": {
        "abstract": "We have conducted a systematic search for high proper motion stars in the Digitized Sky Survey, in the area of the sky north of -2.8 degrees in declination and within 25 degrees of the galactic plane. Using the SUPERBLINK software, a powerful automated blink comparator developed by us, we have identified 601 stars in the magnitude range $9<r<20$ with proper motions in the range $0.5<\\mu<2.0$ arcsec yr$^{-1}$ which have not been observed with the HIPPARCOS satellite. Among those, we recovered 460 stars previously listed in Luyten's proper motion catalogs (LHS, NLTT), and discovered 141 new high proper motion stars, ranging in magnitude from R=13.0 to R=19.8. Only 9 stars from the Luyten catalogs that were not observed by HIPPARCOS could not be recovered with SUPERBLINK: 3 had proper motions larger than the search limit of the code ($\\mu>2.0\\arcsec yr^{-1}$), and 5 were missed because they were either too bright for SUPERBLINK to handle or they are in the immediate proximity of very bright stars. Only one of Luyten's stars (LHS1657) could not be recovered at all, even by visual inspection of the POSS plates, and is now suspected to be bogus. The very high success rate in the recovery by SUPERBLINK of faint Luyten stars suggests that our new survey of high proper motion stars is at least 99\\% complete for stars with proper motions $0.5<\\mu<2.0$ arcsec yr$^{-1}$ down to R=19. This paper includes a list of positions, proper motions, magnitudes, and finder charts for all the new high proper motion stars. ",
        "introduction": "The study of stars with large proper motions forms the basis of our knowledge of the stellar content and dynamical structure of the Galaxy. It is largely through surveys of high proper motion (HPM) stars that the census of stars in the solar vicinity has been established. Since the proper motion of a star is inversely proportional to its distance from the observer, a high proper motion is a strong selection criterion for proximity. The power of proper motion as a selection tool is well illustrated when we consider that among the billions of stars brighter than about 20th magnitude that fill the sky, those with proper motions $\\mu$ larger than half of a second of arc per year ($\\mu>0.5\\arcsec yr^{-1}$) are numbered in the thousands.  The stellar proper motion is also proportional to the velocity of the star, projected on the plane of the sky. This is the main drawback for samples of nearby stars derived from HPM surveys: stars which are moving toward or away from the Sun are systematically missed. Another effect is that samples of HPM stars are always over-represented with stars having large transverse velocities. Essentially, the high velocity stars are sampled over a larger volume than the low velocity stars. While this may bias statistical studies of HPM samples, it is actually extremely useful for finding old disk and halo stars, which are rare in the solar neighborhood, and crucial to our understanding of the structure and stellar content of the Galaxy as a whole.  Of course, direct parallax measurements are ultimately the best method to measure the distances to the stars. However, even the most complete survey to date, carried out by the Hipparcos astrometric satellite, is limited to the brightest hundred thousand stars in the sky (brighter than about 12th magnitude, see the {\\it Hipparcos and Tycho catalogs} Perryman 1997). Next-generation astrometric mission are being planned (the GAIA mission, see Lindegren \\& Perryman 1996; the DIVA mission, see R\\\"oser 1999) which will eventually obtain parallaxes for hundreds of millions more stars. But the identification of nearby stars in the 12th-20th magnitude range still depends largely on surveys of high proper motion stars.  To date, the most extensive catalogs of high proper motion stars remain the two catalogs by W. H. Luyten, published over 20 years ago. One, the {\\it LHS catalogue: a catalogue of stars with proper motions exceeding 0.5$\\arcsec$ annually} (Luyten 1979) lists 3602 objects with estimated proper motions $\\mu\\geq0.500\\arcsec$ yr$^{-1}$ and 867 other stars with estimated proper motions $0.235\\leq\\mu<0.500\\arcsec$ yr$^{-1}$. The {\\it LHS catalogue} was compiled from previous lists of known HPM stars and from a list of new HPM stars detected in the 1950's Palomar Sky Survey by hand-blink or with an automated blink-machine. The complete catalog is available at the ViZier service of the {\\it Centre de Donn\\'ees Astronomiques de Strasbourg} (http://vizier.u-strasbg.fr/) under catalog number I/87B. The {\\it New Luyten Catalogue of stars with proper motions larger than two tenths of an arcsecond} (NLTT, Luyten 1979) lists 58845 HPM stars, the majority of which have estimated proper motions $\\mu\\geq0.18\\arcsec$ yr$^{-1}$. The NLTT catalogue is essentially an extension of the LHS catalogs to stars with smaller proper motions, and is also available at the ViZier service under catalog number I/98A.  Both the LHS and NLTT catalogs were first published in a printed version based on a typewriter copy of the catalog, and both are thus subject to occasional typos and misprints, only a few of which are obvious. The entire LHS catalog has been recently investigated by Bakos, Sahu, \\& N\\'emeth (2002), who searched for each and every one of the LHS stars in the Digitized Sky Survey. While they were able to recover the majority of the LHS stars, they found substantial errors in the coordinates quoted by Luyten. A small number of LHS stars could not even be recovered at all, raising serious doubts as to the reliability of the Luyten catalogs. Furthermore, the LHS and NLTT are also clearly incomplete in some parts of the sky, especially in the southern hemisphere and at low galactic latitudes. The incompleteness of the Luyten catalogs, and the large inaccuracies it sometimes contains, make it a less than completely reliable tool in an epoch where data mining of large electronic databases is becoming an important aspect of astronomical research.  While other surveys of high proper motion stars have been conducted since the publication of the LHS and NLTT catalogs, these surveys have been more limited in magnitude range or in survey area. The most notable is the survey conducted with the Hipparcos satellite (Hipparcos and Tycho catalog, ESA 1997; Tycho-2 catalog, Hog {\\it et al.} 2000) which covers the whole sky but is complete only for relatively bright ($V\\lesssim10$) stars. Deeper large proper motion surveys include the UCAC survey conducted by the US. Naval Observatory, which reaches $V\\approx16$ and covers the whole sky south of $\\delta=+24$, and the SuperCOSMOS Sky Survey (Hambly {\\it et al.} 2001) which so far covers 5000 square degrees in the southern galactic cap (15\\% of the sky) down to $R\\approx19$. Other deep surveys covering smaller areas include the search for very high proper motion stars in 1400 square degrees of second epoch Palomar Sky Survey plates by Monet {\\it et al.} (2000), the Calar-ESO proper motion survey (Ruiz {\\it et al.} 2001), and the southern survey by Wroblewski \\& Costa (2001). Other proper motion surveys have been conducted using fields monitored by gravitational microlensing experiments: the EROS 2 proper motion survey (EROS collaboration 1999), and a search of the MACHO data archive (Alcock {\\it et al.} 2000).  For the past few years, we have been conducting our own systematic, automated survey for high proper motion stars using the Digitized Sky Survey (DSS). We are searching for high proper motion stars in the DSS with SUPERBLINK, a powerful software developed by one of us (SL). SUPERBLINK is an extremely efficient, fully automated, blink-comparator which can be used to identify variable and moving objects in any pair of FITS images of a field observed at two different epochs. We have been feeding DSS images to our SUPERBLINK software and assembling a new catalog of HPM stars, complete with ``blinkable'' double-epoch finder charts. This first paper of a series presents some of the first results of our efforts: an updated census of HPM stars with $\\mu>0.5\\arcsec$ yr$^{-1}$ down to a magnitude R=20.0 north of -2.8$^{\\circ}$ in declination and at low galactic latitudes ($|b|<25^{\\circ}$). We report the recovery of 640 Luyten stars and the discovery of 141 new HPM stars within this area. A second series of papers (L\\'epine et al. 2002, in preparation) will present the results of our parallel campaign of spectroscopic observations of those new high proper motion stars.  In \\S2, we describe the data retrieval, image handling, candidate identification, and photometry. \\S3 and \\S4 contrast our results with the Tycho and Luyten catalogs. Finder charts for all the new high proper motion stars are printed in the appendix. ",
        "conclusions": "The current census of northern hemisphere stars with proper motion $0.5<\\mu<2.0 \\arcsec$ yr$^{-1}$ and within 25$^{\\circ}$ of the galactic plane is now essentially complete down to $r=19$. There are at most a handful of such objects which remain to be identified.  The Luyten catalogs of high proper motion stars (LHS, NLTT) are significantly incomplete for faint ($R>15$) stars at low galactic latitudes. The Luyten catalogs also contain numerous errors in the positions of the stars, and should be replaced with new a new catalog of high proper motion stars.  We are currently expanding our search for high proper motion stars at higher galactic latitude in the northern sky. Results will be published in the next paper of this series."
    },
    "0207/astro-ph0207394_arXiv.txt": {
        "abstract": "The connection between helically isotropic MHD turbulence and mean-field dynamo theory is reviewed. The nonlinearity in the mean-field theory is not yet well established, but detailed comparison with simulations begin to help select viable forms of the nonlinearity. The crucial discriminant is the magnetic helicity, which is known to evolve only on a slow resistive time scale in the limit of large magnetic Reynolds number. Particular emphasis is put on the possibility of memory effects, which means that an additional explicitly time-dependent equation for the nonlinearity is solved simultaneously with the mean-field equations. This approach leads to better agreement with the simulations, while it would also produce more favorable agreement between models and stellar dynamos. ",
        "introduction": "In an early paper Parker \\cite{Par55} identified cyclonic convection as a key process for converting large scale toroidal magnetic field into poloidal fields that have coherence over about half a hemisphere. This process is now generally referred to as the $\\alpha$-effect, although it may arise not only from thermal buoyancy \\cite{SKR66}, but also from magnetic buoyancy \\cite{Lei69,FMSS94,BS98}, the magneto-rotational instability \\cite{BNST95,BD97}, or some other magnetic instability \\cite{GilFox97}. In each case the effect of the Coriolis force together with some kind of radial stratification is crucial for making the motions helical \\cite{KR80}. Upward moving fluid expands, and the Coriolis force makes it rotate retrograde, causing negative (positive) kinetic helicity in the northern (southern) hemisphere. Downward moving fluid contracts, rotates in the prograde direction and contributes in the same sense to negative (positive) kinetic helicity on the northern (southern) hemisphere. This causes a positive $\\alpha$-effect in the northern hemisphere, but if magnetic stresses and shear become strong (for example in accretion discs) the sign may reverse \\cite{Bra98,RP00}.  When combined with differential rotation, the main outcome of $\\alpha$-effect models is the possibility of cyclic magnetic fields associated with latitudinal migration. The first global (two-dimensional) models were presented by Steenbeck \\& Krause \\cite{SK69}, but similar models, with different physics, are still being studied today \\cite{RB95,DC99,KRS01}. The migratory behavior is best seen in contours of the longitudinally averaged mean magnetic field versus latitude and time, which should show tilted structures converging to the equator. Such plots can be compared with the solar butterfly diagram of sunspot numbers (so called because the structures resemble a sequence of butterflies).  A key property of all these models is that not only the motions are helical, but the large scale magnetic field itself is also helical. Of course, not all dynamos require helicity, but nonhelical dynamos tend to generate preferentially small-scale fields \\cite{Cat99}. In a recent attempt, Vishniac \\& Cho \\cite{VC01} proposed a mechanism relevant in particular to accretion discs where shear is strong. Their mechanism would not lead to the production of net magnetic helicity, but numerical simulations \\cite{AB01} failed so far to show convincingly {\\it large scale} dynamo action based on the proposed mechanism. Shear does produce large scale fields, but only in the toroidal direction. It does not explain the latitudinal coherence of the field over several tens of degrees (corresponding to several hundred megameters). On the other hand, there is direct observational evidence that the solar magnetic field is indeed helical. (We shall return to observations in \\Sec{Ssun}.)  The trouble with helical fields is that magnetic helicity is conserved by the induction equation in the ideal limit and can only change on a resistive time scale, provided there is no significant loss through boundaries (at the surface or the equator, for example). This approximate magnetic helicity conservation leads to magnetic field saturation on a resistive time scale \\cite{B01}. Depending on how effectively the boundaries transmit magnetic energy and helicity, the final saturation amplitude will be lowered if losses occur preferentially on large scales while the (linear) growth rate of magnetic energy (past initial saturation) remains {\\it unchanged} \\cite{BD01}. In this sense final saturation can be achieved earlier. The above results are particularly clear when the flow is nearly fully helical, i.e.\\ when the normalized kinetic helicity, $\\epsilon_{\\rm f}\\equiv\\bra{\\oo\\cdot\\uu}/(\\omega_{\\rm rms}u_{\\rm rms})$, where $\\oo=\\nab\\times\\uu$ is the vorticity, is large. In the sun, and probably in all other celestial bodies with rotating turbulence, the relative kinetic helicity is small; $\\epsilon_{\\rm f}\\sim5\\%$. It is however this small helical fraction of the turbulence that is responsible for the a finite but small $\\alpha$-effect, so a proper understanding of its dynamics is crucial if one wants to build models based on the $\\alpha$-effect. Below we shall also discuss the alternative that astrophysical dynamos may shed preferentially small scale helical fields through the boundaries. This could theoretically enhance large scale dynamo action \\cite{BF00,KMRS00,BDS02}. ",
        "conclusions": "In this review we have outlined some of the main results of isotropic MHD simulations in the presence of helicity. We have focussed on the connection with the $\\alpha$-effect in mean-field dynamo theory. We should emphasize that in the case where the magnetic energy density is uniform in space, the agreement between simulations and theory is now well established. In all other cases, things are immediately more complicated. Moreover, dynamical quenching cannot readily be generalized to the case where $\\alpha_{\\rm M}$ varies in space. In that case the equation for the magnetic helicity density would not be gauge-invariant. Another problem arises when $\\alpha_{\\rm K}$ varies in space and if it changes sign across the equator, for example. These are very important aspects requiring clarification. It is quite possible that significant improvement in the theory will soon be possible. Without a corresponding generalization of dynamic $\\alpha$-quenching, if it is ever possible, it would be difficult to use dynamo theory  in astrophysically interesting circumstances."
    },
    "0207/astro-ph0207488_arXiv.txt": {
        "abstract": "The \\LB stars are Population\\,I, late B to early F-type stars, with moderate to extreme (up to a factor 100) surface underabundances of most Fe-peak elements and solar abundances of lighter elements (C, N, O, and S). To put constraints on the various existing theories that try to explain these peculiar stars, we investigate the observational properties of \\LB stars compared to a reference sample of normal stars. Using various photometric systems and Hipparcos data, we analyze the validity of standard photometric calibrations, elemental abundances, and Galactic space motions.  There crystallizes a clear picture of a homogeneous group of Population\\,I objects found at all stages of their main-sequence evolution, with a peak at about 1~Gyr. No correlation of astrophysical parameters such as the projected rotational velocities or elemental abundances with age is found, suggesting that the a-priori unknown mechanism, which creates \\LB stars, works continuously for late B to early F-type stars in all stages of main-sequence evolution. Surprisingly, the sodium abundances seem to indicate an interaction between the stars and their local environment. ",
        "introduction": "Knowledge of the evolutionary status of the members of the \\LB group is essential to put tight constraints on the astrophysical processes behind this phenomenon. We have chosen the following working definition as group characteristics: late B- to early F-type, Population\\,I stars with apparently solar abundances of the light elements (C, N, O and S) and moderate to strong underabundances of Fe-peak elements (see Faraggiana \\& Gerbaldi 1998 for a critical summary of the various definitions). Only a maximum of about 2\\% of all objects in the relevant spectral domain are believed to be \\LB\\hskip-2pt-type stars (Paunzen 2001). That already suggests either that the mechanism responsible for the phenomenon works on a very short time-scale (10$^{6}$\\,yrs) or else the general conditions for the development are very strict.  We know already of a few \\LB stars in the Orion OB1 association and one candidate in NGC~2264 (Paunzen 2001), for both of which log\\,$t$\\,$\\approx$\\,7.0. The evolutionary status for two \\LB\\hskip-2pt-type spectroscopic-binary systems (HD~84948 and HD~171948) has been determined as very close to the Zero Age Main Sequence (ZAMS hereafter) for HD~171948 and to the Terminal Age Main Sequence (TAMS hereafter) for HD~84948 (Iliev et al. 2002). The results for the other Galactic field stars are not clear.  In 1995, Iliev \\& Barzova summarized the evolutionary status of 20 \\LB stars (and Vega) using Str\\\"omgren $uvby\\beta$ photometric data. They concluded that most of the stars studied are in the middle of their main-sequence evolution, with only a few objects near the ZAMS.  Paunzen (1997) investigated the parallaxes measured by the Hipparcos satellite for a sample of \\LB\\hskip-2pt-type stars in order to derive luminosities, masses and ages for 18 objects in common with Iliev \\& Barzova (1995). He found no systematic influence of the distance, effective temperature, metallicity and rotational velocity on the difference between photometrically calibrated absolute magnitudes and those derived from Hipparcos parallaxes. Six objects were found to be very close to the ZAMS and a hypothesis was proposed that all other stars are in their Pre-Main-Sequence (PMS hereafter) phase.  Later on, Bohlender, Gonzalez \\& Matthews (1999) and Faraggiana \\& Bonifacio (1999) challenged that hypothesis with plausible arguments such as the unusually vigorous star-forming activity that it implied in the solar neighbourhood and a statistical analysis of normal-type stars.  Already Gray \\& Corbally (1998) stated that the \\LB phenomenon can be found from very early stages to well into the main-sequence life of A-type stars. That conclusion was based on the incidence of \\LB stars among very young A-type stars, which is not very different from the incidence among Galactic field stars.  In this paper we present a much more extensive investigation, including the data of the Hipparcos satellite. With the help of photometric data of the Johnson {\\it UBV}, Str\\\"omgren $uvby\\beta$ and Geneva 7-colour systems, different calibrations of the absolute magnitude, effective temperature and surface gravity are applied and compared.  From evolutionary models (Claret 1995), masses and ages are estimated. They are compared with those derived by Iliev \\& Barzova (1995) and Paunzen (1997).  As a further step, the proper motions of \\LB stars are used to calculate space velocities. That very important information should help further to sharpen the group properties and to sort out probably misclassified stars.  Another point investigated is the question as to whether there exists a typical abundance pattern for the \\LB group. Heiter (2002) and Heiter et al.~(2002) tried to shed more light on the abundance pattern in the context of the proposed theories. We have searched for correlations of the individual abundances with, especially, mass and age.  However, we will not comment on our results in the context of the developed theories and models, because they still depend on too many free parameters. The aim of this paper is not to promote one of the suggested theories but rather to find strict observational constraints which should be incorporated into future theoretical investigations. \\def\\b{\\hskip0.5em}  \\def\\bb{\\hskip1em}  \\def\\d{\\hskip-0.7em} \\begin{table*} \\caption[]{Photometric data, stellar parameters and calibrated values for the program stars. In parenthesis are the errors in the final digits of the corresponding quantity.} \\label{lb_basic} \\begin{center} \\begin{tabular}{lllccccrllll} \\hline \\bb HD\t&\t\\b HR & \\b HIP\t&\t$V$\t&\t$B-V$\t&\t$b-y$\t&\t$A_{\\rm V}$\t&\t\\multicolumn{1}{c}{$v$\\,sin\\,$i$} & \\multicolumn{1}{c}{$T_{\\rm eff}$}\t\t&\t\\multicolumn{1}{c}{\\d log\\,$g$(phot)}\t&\t\\multicolumn{1}{c}{$M_{\\rm V}$}\t\t& \\multicolumn{1}{c}{\\d log\\,L$_{\\ast}$/L$_{\\odot}$}\t\\\\ & & & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[kms$^{-1}$]} & \\multicolumn{1}{c}{[K]} & \\multicolumn{1}{c}{\\d [dex]} & \\multicolumn{1}{c}{[mag]} \\\\ \\noalign{\\vskip6pt} \\bb\\b319\t&\\bb\t12&\\bb\\b636\t&\t5.934\t&\t0.141\t&\t0.079\t&\t0.004\t&{60\\bb}\t&\t8020(135)\t&\t3.74(8)\t&\t1.27(19)\t&\t1.45(8)\t\\\\ \\bb6870\t&\t\t&\t\\bb5321\t&\t7.494\t&\t0.246\t&\t0.164\t&\t0.000\t&{165\\bb}\t&\t7330(102)\t&\t3.84(11)\t&\t2.29(42)\t&\t1.02(17)\t\\\\ \\bb7908\t&\t\t&\t\\bb6108\t&\t7.288\t&\t0.272\t&\t0.192\t&\t0.000\t&\t\t&\t7145(87)\t&\t4.10(12)\t&\t2.60(18)\t&\t0.90(7)\t\\\\ \\b11413\t&\t \\b 541\t&\t\\bb8593\t&\t5.940\t&\t0.147\t&\t0.105\t&\t0.004\t&{125\\bb}\t&\t7925(124)\t&\t3.91(21)\t&\t1.49(10)\t&\t1.35(4)\t\\\\ \\b13755\t&\t\t&\t\\b10304\t&\t7.844\t&\t0.318\t&\t0.181\t&\t0.000\t&\t\t&\t7080(161)\t&\t3.26(10)\t&\t0.93(10)\t&\t1.57(4)\t\\\\ \\b15165\t&\t\t&\t\\b11390\t&\t6.705\t&\t0.333\t&\t0.191\t&\t0.010\t&{90\\bb}\t&\t7010(167)\t&\t3.23(10)\t&\t1.12(16)\t&\t1.50(6)\t\\\\ \\b23392\t&\t\t&\t\\b17462\t&\t8.260\t&\t0.020\t&\t0.014\t&\t0.094\t&\t\t&\t9805(281)\t&\t4.35(9)\t&\t1.43(30)\t&\t1.45(12)\t\\\\ \\b24472\t&\t\t&\t\\b18153\t&\t7.092\t&\t0.304\t&\t0.214\t&\t0.003\t&\t\t&\t6945(131)\t&\t3.81(16)\t&\t2.14(11)\t&\t1.09(5)\t\\\\ \\b30422\t&\t1525\t&\t\\b22192\t&\t6.186\t&\t0.190\t&\t0.101\t&\t0.014\t&{135\\bb}\t&\t7865(108)\t&\t4.00(20)\t&\t2.35(1)\t&\t1.01(1)\t\\\\ \\b31295\t&\t1570\t&\t\\b22845\t&\t4.648\t&\t0.085\t&\t0.044\t&\t0.063\t&{115\\bb}\t&\t8920(177)\t&\t4.20(1)\t&\t1.66(22)\t&\t1.32(9)\t\\\\ \\b35242\t&\t1777\t&\t\\b25205\t&\t6.348\t&\t0.122\t&\t0.068\t&\t0.042\t&{90\\bb}\t&\t8250(103)\t&\t3.90(14)\t&\t1.75(22)\t&\t1.26(9)\t\\\\ \\b54272\t&\t\t&\t\t&\t8.800\t&\t0.261\t&\t0.214\t&\t0.000\t&\t\t&\t7010(217)\t&\t3.83(10)\t&\t2.33(30)\t&\t1.01(12)\t\\\\ \\b74873\t&\t3481\t&\t\\b43121\t&\t5.890\t&\t0.115\t&\t0.064\t&\t0.078\t&{130\\bb}\t&\t8700(245)\t&\t4.21(11)\t&\t1.82(1)\t&\t1.24(1)\t\\\\ \\b75654\t&\t3517\t&\t\\b43354\t&\t6.384\t&\t0.242\t&\t0.161\t&\t0.012\t&{45\\bb}\t&\t7350(104)\t&\t3.77(11)\t&\t1.83(12)\t&\t1.20(5)\t\\\\ \\b81290\t&\t\t&\t\\b46011\t&\t8.866\t&\t0.332\t&\t0.254\t&\t0.124\t&{55\\bb}\t&\t6895(214)\t&\t3.82(28)\t&\t1.85(30)\t&\t1.20(12)\t\\\\ \\b83041\t&\t\t&\t\\b47018\t&\t8.927\t&\t0.294\t&\t0.223\t&\t0.161\t&{95\\bb}\t&\t7120(208)\t&\t3.76(20)\t&\t1.70(30)\t&\t1.26(12)\t\\\\ \\b83277\t&\t\t&\t\\b47155\t&\t8.304\t&\t0.311\t&\t0.226\t&\t0.131\t&\t\t&\t7000(189)\t&\t3.67(18)\t&\t1.49(29)\t&\t1.35(12)\t\\\\ \\b84123\t&\t\t&\t\\b47752\t&\t6.840\t&\t0.297\t&\t0.235\t&\t0.040\t&{20\\bb}\t&\t7025(175)\t&\t3.73(17)\t&\t1.58(15)\t&\t1.31(6)\t\\\\ \\b87271\t&\t\t&\t\\b49328\t&\t7.120\t&\t0.172\t&\t0.151\t&\t0.008\t&\t\t&\t7515(232)\t&\t3.43(10)\t&\t1.02(8)\t&\t1.53(3)\t\\\\ \\b90821\t&\t\t&\t\t&\t9.470\t&\t0.105\t&\t0.068\t&\t0.013\t&{150\\bb}\t&\t8190(79)\t&\t3.73(10)\t&\t0.74(30)\t&\t1.66(12)\t\\\\ \\b91130\t&\t4124\t&\t\\b51556\t&\t5.902\t&\t0.109\t&\t0.073\t&\t0.000\t&{135\\bb}\t&\t8135(98)\t&\t3.78(10)\t&\t1.36(26)\t&\t1.42(11)\t\\\\ 101108\t&\t\t&\t\\b56768\t&\t8.880\t&\t0.179\t&\t0.114\t&\t0.006\t&{90\\bb}\t&\t7810(80)\t&\t3.90(18)\t&\t1.33(30)\t&\t1.42(12)\t\\\\ 102541\t&\t\t&\t\\b57567\t&\t7.939\t&\t0.230\t&\t0.163\t&\t0.097\t&\t\t&\t7665(168)\t&\t4.22(16)\t&\t2.34(21)\t&\t1.01(9)\t\\\\ 105058\t&\t\t&\t\\b58992\t&\t8.900\t&\t0.183\t&\t0.129\t&\t0.009\t&{140\\bb}\t&\t7740(171)\t&\t3.77(30)\t&\t0.86(30)\t&\t1.60(12)\t\\\\ 105759\t&\t\t&\t\\b59346\t&\t6.550\t&\t0.218\t&\t0.142\t&\t0.000\t&{120\\bb}\t&\t7485(102)\t&\t3.65(10)\t&\t1.35(21)\t&\t1.40(8)\t\\\\ 106223\t&\t\t&\t\\b59594\t&\t7.431\t&\t0.288\t&\t0.228\t&\t0.015\t&{90\\bb}\t&\t6855(247)\t&\t3.49(18)\t&\t1.83(45)\t&\t1.22(18)\t\\\\ 107233\t&\t\t&\t\\b60134\t&\t7.353\t&\t0.255\t&       0.192\t&\t0.048\t&{95\\bb}  &\t7265(143)\t&\t4.03(10)\t&\t2.64(13)\t&\t0.88(5)\t\\\\ 109738\t&\t\t&\t\t&\t8.277\t&\t0.198\t&\t0.161\t&\t0.073\t&\t\t&\t7610(145)\t&\t3.90(13)\t&\t1.85(30)\t&\t1.20(12)\t\\\\ 110377\t&\t4824\t&\t\\b61937\t&\t6.228\t&\t0.195\t&\t0.120\t&\t0.000\t&{170\\bb}\t&\t7720(89)\t&\t3.97(14)\t&\t1.96(11)\t&\t1.16(5)\t\\\\ 110411\t&\t4828\t&\t\\b61960\t&\t4.881\t&\t0.077\t&\t0.040\t&\t0.045\t&{165\\bb}\t&\t8930(206)\t&\t4.14(14)\t&\t1.90(28)\t&\t1.22(11)\t\\\\ 111005\t&\t\t&\t\\b62318\t&\t7.959\t&\t0.376\t&\t0.224\t&\t0.009\t&\t\t&\t6860(66)\t&\t3.72(10)\t&\t1.76(53)\t&\t1.24(21)\t\\\\ 111604\t&\t4875\t&\t\\b62641\t&\t5.886\t&\t0.160\t&\t0.112\t&\t0.000\t&{180\\bb}\t&\t7760(149)\t&\t3.61(25)\t&\t0.48(7)\t&\t1.75(3)\t\\\\ 120500\t&\t\t&\t\\b67481\t&\t6.600\t&\t0.131\t&\t0.068\t&\t0.017\t&{125\\bb}\t&\t8220(70)\t&\t3.86(10)\t&\t0.85(34)\t&\t1.62(13)\t\\\\ 120896\t&\t\t&\t\\b67705\t&\t8.495\t&\t0.296\t&\t0.166\t&\t0.000\t&\t\t&\t7260(89)\t&\t3.76(10)\t&\t1.90(30)\t&\t1.18(12)\t\\\\ 125162\t&\t5351\t&\t\\b69732\t&\t4.186\t&\t0.084\t&\t0.051\t&\t0.039\t&{115\\bb}\t&\t8720(156)\t&\t4.07(9)\t&\t1.71(23)\t&\t1.28(9)\t\\\\ 125889\t&\t\t&\t\t&\t9.849\t&\t0.241\t&\t0.206\t&\t0.108\t&\t\t&\t7275(175)\t&\t3.88(9)\t&\t2.32(30)\t&\t1.01(12)\t\\\\ 130767\t&\t\t&\t\\b72505\t&\t6.905\t&\t0.042\t&\t0.002\t&\t0.000\t&\t\t&\t9195(220)\t&\t4.10(8)\t&\t1.27(1)\t&\t1.48(1)\t\\\\ 142703\t&\t5930\t&\t\\b78078\t&\t6.120\t&\t0.240\t&\t0.182\t&\t0.021\t&{100\\bb}\t&\t7265(150)\t&\t3.93(12)\t&\t2.41(12)\t&\t0.97(5)\t\\\\ 142944\t&\t\t&\t\t&\t7.179\t&\t0.293\t&\t0.201\t&\t0.013\t&{180\\bb}\t&\t7000(125)\t&\t3.19(4)\t&\t0.80(30)\t&\t1.62(12)\t\\\\ 149130\t&\t\t&\t\\b81329\t&\t8.498\t&\t0.342\t&\t0.233\t&\t0.109\t&\t\t&\t6945(103)\t&\t3.49(5)\t&\t1.51(30)\t&\t1.34(12)\t\\\\ 153747\t&\t\t&\t\\b83410\t&\t7.420\t&\t0.140\t&\t0.098\t&\t0.128\t&\t\t&\t8205(90)\t&\t3.70(24)\t&\t1.24(30)\t&\t1.46(12)\t\\\\ 154153\t&\t6338\t&\t\\b83650\t&\t6.185\t&\t0.284\t&\t0.199\t&\t0.020\t&\t\t&\t7055(120)\t&\t3.56(6)\t&\t1.86(29)\t&\t1.19(11)\t\\\\ 156954\t&\t\t&\t\\b84895\t&\t7.679\t&\t0.294\t&\t0.200\t&\t0.050\t&{50\\bb}\t&\t7130(93)\t&\t4.04(13)\t&\t2.81(33)\t&\t0.82(13)\t\\\\ 168740\t&\t6871\t&\t\\b90304\t&\t6.138\t&\t0.201\t&\t0.136\t&\t0.035\t&{145\\bb}\t&\t7630(81)\t&\t3.88(14)\t&\t1.82(2)\t&\t1.21(1)\t\\\\ 168947\t&\t\t&\t\t&\t8.123\t&\t0.196\t&\t0.172\t&\t0.116\t&\t\t&\t7555(185)\t&\t3.67(10)\t&\t1.28(30)\t&\t1.43(12)\t\\\\ 170680\t&\t6944\t&\t\\b90806\t&\t5.132\t&\t0.008\t&\t0.008\t&\t0.091\t&{205\\bb}\t&\t9840(248)\t&\t4.15(6)\t&\t0.83(23)\t&\t1.70(9)\t\\\\ 175445\t&\t\t&\t\\b92884\t&\t7.792\t&\t0.110\t&\t0.055\t&\t0.108\t&\t\t&\t8520(198)\t&\t3.96(10)\t&\t1.08(27)\t&\t1.53(11)\t\\\\ 183324\t&\t7400\t&\t\\b95793\t&\t5.795\t&\t0.084\t&\t0.051\t&\t0.083\t&{90\\bb}\t&\t8950(204)\t&\t4.13(4)\t&\t1.64(42)\t&\t1.32(17)\t\\\\ 184779\t&\t\t&\t\t&\t8.940\t&\t0.224\t&\t0.187\t&\t0.000\t&\t\t&\t7210(173)\t&\t3.63(21)\t&\t1.26(30)\t&\t1.43(12)\t\\\\ 192640\t&\t7736\t&\t\\b99770\t&\t4.934\t&\t0.154\t&\t0.099\t&\t0.016\t&{80\\bb}\t&\t7940(96)\t&\t3.95(18)\t&\t1.84(1)\t&\t1.22(1)\t\\\\ 193256\t&\t7764C\t&\t100286\t&\t7.721\t&\t0.213\t&\t0.116\t&\t0.063\t&{250\\bb}\t&\t7740(94)\t&\t3.69(17)\t&\t1.08(30)\t&\t1.51(12)\t\\\\ 193281\t&\t7764A\t&\t100288\t&\t6.557\t&\t0.190\t&\t0.098\t&\t0.111\t&{95\\bb}\t&\t8035(115)\t&\t3.54(4)\t&\t0.41(30)\t&\t1.79(12)\t\\\\ 198160\t&\t7959\t&\t102962\t&\t5.663\t&\t0.189\t&\t0.108\t&\t0.022\t&{200\\bb}\t&\t7870(129)\t&\t3.99(9)\t&\t1.47(41)\t&\t1.36(16)\t\\\\ 204041\t&\t8203\t&\t105819\t&\t6.456\t&\t0.161\t&\t0.092\t&\t0.026\t&{65\\bb}\t&\t7980(97)\t&\t3.97(8)\t&\t1.75(18)\t&\t1.25(7)\t\\\\ 210111\t&\t8437\t&\t109306\t&\t6.377\t&\t0.203\t&\t0.136\t&\t0.000\t&{55\\bb}\t&\t7550(123)\t&\t3.84(15)\t&\t1.76(15)\t&\t1.23(6)\t\\\\ 216847\t&\t\t&\t113351\t&\t7.060\t&\t0.242\t&\t0.155\t&\t0.000\t&\t\t&\t7355(78)\t&\t3.47(14)\t&\t0.93(24)\t&\t1.56(10)\t\\\\ 221756\t&\t8947\t&\t116354\t&\t5.576\t&\t0.095\t&\t0.056\t&\t0.043\t&{105\\bb}\t&\t8510(188)\t&\t3.90(3)\t&\t1.16(16)\t&\t1.50(6)\t\\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We have used all currently available photometric data as well as Hipparcos data to determine astrophysical parameters such as the effective temperatures, surface gravities and luminosities. As a next step, masses and ages were calibrated within appropriate post-MS evolutionary models. Furthermore, Galactic space motions were calculated with the help of radial velocities from the literature. The comparison with already published results shows good agreement of the derived parameters. All results were compared with those of a test sample of normal-type objects in the same spectral range chosen in order to match the $(B-V)_0$ distribution of the \\LB group. From a comprehensive statistical analysis we conclude: \\begin{itemize} \\item The standard photometric calibrations within the Johnson {\\it UBV}, Str\\\"omgren $uvby\\beta$ and Geneva 7-colour systems are valid for this group of chemically peculiar stars. \\item The group of \\LB stars consists of true Population\\,I objects which can be found over the whole area of the MS with a peak at a rather evolved stage ($\\approx$\\,1\\,Gyr). That is in line with the distribution of the test sample. \\item The \\LB type group is not significantly distinct from normal stars except, possibly, by having slightly lower temperatures and masses for $t_{\\rm rel}$\\,$>$\\,0.8. The $v$\\,sin\\,$i$ range is rather narrow throughout the MS with a mean value of about 120\\,kms$^{-1}$. \\item There seems to exist a non-uniform distribution of effective temperatures for group members with a large proportion of objects (more than 70\\%) cooler than 8000\\,K. \\item It seems that objects with peculiar hydrogen-line profiles are preferentially found among later stages of stellar evolution. \\item No correlation of age with elemental abundance or projected rotational velocity has been detected. \\item A comparison of the stellar Na abundances with nearby IS sight lines hints at an interaction between the \\LB stars and the ISM. \\item There is one single mechanism responsible for the observed phenomenon which produces moderate to strong underabundances working continuously from very early (10\\,Myr) to very late evolutionary stages (2.5\\,Gyr). It produces the same absolute abundances throughout the MS lifetime for 2\\% of all luminosity class V objects with effective temperatures from 10500\\,K to 6500\\,K. \\item The current list of stars seem to define a very homogeneous group, validating the proposed membership criteria in the optical and UV region. \\end{itemize} These rather strict observational results for a significant number of \\LB stars will need to be taken into account in future work on theories and models trying to explain the phenomenon. The constraints presented here will help considerably to reduce the number of free parameters in the models and finally to provide a critical test for them. \\\\ \\\\ {\\noindent \\footnotesize {\\bf Acknowledgements.} This work was partly supported by the Fonds zur F\\\"orderung der wissenschaftlichen Forschung, project P14984 and it is dedicated to G.~Derka who died during its preparation. Use was made of the SIMBAD database, operated at CDS, Strasbourg, France and the GCPD database, operated at the Institute of Astronomy of the University of Lausanne. }"
    },
    "0207/astro-ph0207441_arXiv.txt": {
        "abstract": "Different scenarios of baryogenesis are briefly reviewed from the point of view of possibility of generation of cosmologically interesting amount of antimatter. It is argued that creation of antimatter is possible and natural in many models. In some models not only anti-helium may be produced but also a heavier anti-elements and future observations of the latter would be critical for discovery or establishing stronger upper limits on existence of antimatter. Incidentally a recent observation of iron-rich quasar may present a support to one special model of antimatter creation. ",
        "introduction": "} Our region of the universe is certainly dominated by matter, protons, electrons, and nuclei consisting of protons and neutrons. No antimatter in any significant amount is observed. A little of antiprotons and positrons in cosmic rays can be explained by their secondary origin in collisions of protons, electrons or photons with usual matter. Cosmological excess of matter over antimatter is described by the ratio \\be \\beta = {N_B - N_{\\bar B} \\over N_\\gamma} \\approx 6\\cdot 10^{-10} \\label{beta} \\ee where $N_{B,\\bar B,\\gamma}$ are respectively the cosmic number densities of baryons, antibaryons, and photons in microwave background radiation (CMBR). At the present day $N_\\gamma = 412/ {\\rm cm}^3$ and $N_B \\gg N_{\\bar B}$ (at least in our  neighborhood.). It is believed that in the early universe, at high temperatures, $T> 100$ MeV, the number densities of baryons and antibaryons (at these temperatures they were in quark state) were almost equal with relative accuracy of the order of $\\beta$.  According to simple models of baryogenesis, pioneered by Sakharov\\cite{sakharov67} in 1967, baryon asymmetry is homogeneous, i.e. $\\beta$ does not depend on space points, and the total baryonic charge of the universe is non-zero: \\be B_{tot} = \\int \\beta\\,d^3 x \\neq 0 \\label{Btot} \\ee Still it is not excluded neither theoretically nor observationally that this may be not so and we face the following big questions:  \\begin{enumerate}  \\item{} Is $\\beta = const$ or it could be a function of space point, $\\beta = \\beta (x)$?  \\item{} If $\\beta = \\beta (x)$ what is the characteristic scale $L_B$ of its variation? Especially interesting is if $L_B$ may be smaller than the present-day horizon $L_{hor} \\sim 3 $ Gpc, or it is possible that $L_B < L_{hor}$?  \\item{} If $\\beta $ indeed varies, may it be that in some astronomically sizable regions $\\beta <0$, that is some parts of the universe are antimatter dominated?  \\item{} If $\\beta <0$ is allowed what is the global baryonic charge of the universe? Is $B_{tot} \\neq 0$, so the universe is globally charge asymmetric or $B_{tot} =0$ and the universe is globally charge symmetric.  \\end{enumerate}  In this talk I will discuss the present observational bounds on existence of antimatter and scenarios of baryogenesis that might lead to astronomically interesting antimatter domains or antimatter objects. ",
        "conclusions": "It seems that creation of astronomically interesting amount of antimatter is not only possible but quite natural in many scenarios of baryogenesis. In particular, the conditions for creation of antimatter are especially favorable if a condensate of a complex scalar field existed in the early universe which was a source of temporary charge parity breaking.  Theory allows globally charge symmetric universe which, however may have already problems with the existing observation. On the other hand, a dominance of matter or antimatter looks equally natural and a ratio, $\\epsilon$, of the amount of cosmic antimatter to matter is an unknown parameter of the theory. At the present day we cannot exclude neither small not large  $\\epsilon$ and it even possible that we live in a relatively small matter domain in antimatter dominated universe. Still more detailed and accurate theoretical calculations in concrete models are necessary since they could already provide sensible bounds on the possible abundance and the distance to the nearest antimatter domains or astronomical objects. As we discussed in Section~\\ref{sec-anti} antimatter may even live in the halo of our Galaxy.   A search for cosmic antimatter is a very exciting, though extremely difficult, challenge for the future observations. Its discovery would indicate to an unusual mechanism of charge symmetry breaking in cosmology and might be an additional proof of inflationary scenario. At the moment it seems that the most promising way to search for cosmic antimatter is to look for anti-nuclei in cosmic rays. According to the discussed above models those could be not only anti-helium but heavier ones up to anti-iron."
    },
    "0207/astro-ph0207327_arXiv.txt": {
        "abstract": "If core collapse leads to the formation of a rapidly rotating bar-unstable proto-neutron star surrounded by fall-back material, then we might expect it to cool and fragment to form a double (proto)-neutron star binary into a super-close orbit.  The lighter star should survive for awhile, until tidal mass loss propels it toward the minimum stable mass of a (proto)-neutron star, whereupon it explodes.  Imshennik \\& Popov have shown that the explosion of the unstable, cold star can result in a large recoil velocity of the remaining neutron star. Here, we consider several factors that mitigate the effect and broaden the range of final recoil speeds, in particular the finite velocity and gravitational deflection of the ejecta, a range of original masses for the low mass companion and its cooling history, rotational phase averaging of the momentum impulse from non-instantaneous mass loss, and the possibility of a common envelope phase. In spite of these mitigating factors, we argue that this mechanism can still lead to substantial neutron star recoil speeds, close to, or even above, $1000\\kms$. ",
        "introduction": "It is recognized that radio pulsars have peculiar space velocities between $\\approx 30 \\kms$ and $\\approx 1600\\kms$, significantly greater than those of their progenitor stars. The highest speed ever recorded is from the Guitar Nebula pulsar (Cordes et al. 1993; Cordes \\& Chernoff 1998), while the lowest has been recently measured for B2016+28 by Brisken  et al. (2002) from very accurate VLBA  pulsar parallaxes. Statistical studies, aimed at inferring the peculiar velocity at birth from the observed speed, give mean three-dimensional velocities of 100 -- 500 $\\kms$ for the isolated pulsars (Lyne \\& Lorimer 1994; Lorimer, Bailes, \\& Harrison 1997; Hansen \\& Phinney 1997; Cordes \\& Chernoff 1998; Arzoumanian, Chernoff, \\& Cordes 2002).  An early explanation for the large space velocities called for {\\it recoil} in a close binary that becomes unbound at the time of (symmetric) supernova explosion (Blaauw 1961; Iben \\& Tutukov 1996). Now, a number of observations  hint at a {\\it natal} origin of these high space velocities, or at a combination of orbital disruption (when the progenitor lives in a binary) and internal kick reaction. Evidence for a kick at the time of neutron star birth is now found in a variety of systems: in runaway O/B associations (Leonard \\& Dewey 1992), in highly eccentric Be/NS binaries (van den Heuvel \\& Rappaport 1986; Portgies-Zwart \\& Verbunt 1996), in the binary pulsar J0045-7319 (Kaspi et al. 1996; to explain its current spin-orbit configuration), and in double neutron star binaries such as B1913+16 (Bailes 1988; Weisberg, Romani, \\& Taylor 1989; Cordes, Wasserman, \\& Blaskiewicz 1990; Kramer 1998; Wex, Kalogera \\& Kramer 2000) where misalignment between the spin and orbital angular momentum axes indicates velocity asymmetry in the last supernova. All these observations support the view that the formation of neutron stars is accompanied by anisotropic explosion. This notion is bolstered by evolutionary studies of binary populations (Dewey \\& Cordes 1987; Fryer \\& Kalogera 1997; Fryer, Burrows, \\& Benz 1998) and by studies (e.g. Cordes \\& Wasserman  1984) on the survival of binary systems into their  late evolutionary stages after supernova explosion.  Among the physical processes that have been proposed to account for the kicks are large-scale density asymmetries seeded in the pre-supernova core (leading to anisotropic shock propagation), asymmetric  neutrino emission in presence of ultra-strong magnetic fields (see Lai, Chernoff, \\& Cordes 2001 for a review), or off-centered electromagnetic dipole emission from the young pulsar (Harrison \\& Tademaru 1975). However, none of these  mechanisms can explain  kicks as large as $\\sim 1600 \\kms$.   In this paper we reconsider the idea first put forward by Imshennik \\& Popov (1998) that, in the collapse of a rotating core, one or more self-gravitating lumps of neutronized matter may form in close orbit around the central nascent neutron star, transfer mass in the short lived binary, and ultimately explode, causing the remaining, massive neutron star to acquire a substantial kick velocity, as high as the highest observed. The light member explodes as mass transfer drives it below the minimum stable mass for a neutron stars. In the light star, stability is lost upon decompression by the $\\beta-$decaying  neutrons and nuclear fissions by radio-active neutron-rich nuclei (Colpi, Shapiro \\& Teukolsky 1989, 1991; Blinnikov et al. 1990), that deposit energy driving matter into rapid expansion (Colpi, Shapiro \\& Teukolsky 1993; Sumiyoshi et al. 1998). The kick tries its origin from the orbital motion of this evanescent super-close binary, that forms in the collapse of a rapidly rotating (isolated) iron core.   We will study several effects that may modify the magnitude of the kick, such as gravitational bending of the exploding debris, rotational averaging of the momentum impulse, orbit decay, and delayed neutron star cooling.  Formation of a proto-neutron star companion around the main neutron star has never been verified in numerical simulations due to computational limitations. For this reason we elaborate on a study of Bonnell (1994) on the formation of binary/multiple systems in collapsing gas cloud cores and its extension to the stellar core collapse in the aftermath of a supernova explosion (Bonnell \\& Pringle 1995) to motivate our working hypothesis. The formation of such exotic binaries has been conjectured to occur in many works (Ruffini \\& Wheeler 1971; Clark \\& Eardley 1977; Blinnikov, Novikov, Perevodchikova, \\& Polnarev 1984; Nakamura \\& Fukugita 1989; Stella \\& Treves 1987).     \\section {Light fragments around proto-neutron stars}  \\subsection {The Scenario}  Formation of a light companion around a main body implies breaking of spherical and axial symmetry during collapse and following core bounce.  During dynamical collapse, unstable bar modes ($m=2$) can grow in a fluid (even non rotating) that may end with fragmentation. However this is known to occur only if the cloud core contracts almost isothermally, as in the case of star's formation from unstable cold gas clouds (Bonnell 1994).  Core collapse in type II supernovae is far from isothermal (it is described by a effective polytropic index $\\gamma\\simeq 1.3$) so that instabilities of this type do not have time to grow (Lai 2000) \\footnote{Large-scale asymmetries imprinted in the iron core prior to collapse may lead to anisotropic explosions that produce kicks as indicated by Goldreich, Lai \\& Sahrling (1996). Whether they lead to fragmentation is unknown.}, and simulations of non-axisymmetric rotating core collapse confirm this trend (Rampp, Muller \\& Ruffer 1998; Centrella, New, Lowe, \\& Brown, 2001). Can fragmentation/fission be excited after core bounce ?  Rapid rotation in equilibrium bodies is known to excite non-axisymmetric dynamical instabilities and these instabilities may grow in the proto-neutron star core. Interestingly, core collapse simulations of unstable rotating iron cores (Heger, Langer \\& Woosley 2000; Fryer \\& Heger 2000) or polytropes (Zwerger \\& Muller 1997) indicate that proto-neutron stars, soon after formation, can rotate differentially above the dynamical stability limit set when the rotational to gravitational potential energy ratio $T_{\\rm rot}/\\vert W\\vert$ is larger than the value $\\beta_{\\rm dyn}=0.25-0.26$ (Saijo et al. 2001). Strong non-linear growth of the dominant bar-like deformation ($m=2$) is seen in these cores (described as polytropes by Rampp, Muller, \\& Ruffert 1998). However there is no sign of fission into separate condensations.  The bar evolves, producing two spiral arms that transport the core's excess angular momentum outwards (see also Shibata, Baumgarte \\& Shapiro 2000). This reduces the bar's angular momentum such that a single central body is formed. How can a body develop local condensations when bar-unstable ?   According to Bonnell's picture, the evolution of the bar instability is more complex, in reality.  If the rapidly spinning proto-neutron star core goes bar unstable when surrounded by a fall-back disk, then matter present in the bar-driven spiral arms interacts with this material.  The sweeping of a spiral arm into fall-back gas can gather sufficient matter to condense into a fragment of neutronized matter. This occurs because the $m=1$ mode grows during the development of the $m=2$ mode. The $m=1$ mode causes the displacement of the unpinned (free-to-move) core and creates an off-center spiral arm that sweeps up more material on one side than the other during \"continuing\" accretion. The condensation eventually collapses into a low mass neutron star.  Detailed simulations confirming or dismissing the occurrence of such instability are still lacking, so further considerations of the lump masses, temperatures and entropy contents are necessarily speculative.     \\subsection {Cooling scenarios and the minimum mass}    Here we wish to explore the possibility that a light (proto-)neutron star forms around the main central body, lives for a while, and later explodes, imprinting a kick to the  neutron star  that remains (due to linear momentum conservation) before gravitational waves or hydrodynamical effects can drive it toward the central star.  Can a light (proto)-neutron star form from the condensation of material accumulated in the off-centered spiral arm ? What limits can be imposed on its mass ?  Cooling plays a key role in addressing these questions. Goussard, Haensel \\& Zdunik (1998), Strobel, Schaab \\& Weigel (1999), and  Strobel \\& Weigel (2001) have shown that the value of the minimum stable mass $m_{\\mmc}$ for a neutron star (located at the turning point in the mass-radius relation of equilibria) is a function of the temperature: it  varies from $\\simgreat 1 \\msun$ at 50-100 milliseconds after core bounce (setting the actual value of the mass of the central neutron star), to $\\sim 0.7\\msun$ after $\\sim 1$ second, down to $\\sim 0.3\\msun$ after 30 seconds, reaching the value of $\\sim 0.0925\\msun$ (Baym, Pethick \\& Sutherland 1971) known for cold catalyzed matter ($T<1$ MeV) several hundred seconds later. Thus, the lump of nuclear matter gathered in the spiral arm by the instability may become self-bound if its mass $m$  is above the minimum corresponding to that particular temperature. Once formed it is stabilized against expansion by cooling. The actual value of $m$ is thus determined by the overall dynamics of collapse after core bounce and varies between 0.0925 and $\\sim 1\\msun:$ It depends on the time at which the instability sets in, on the amount of fall-back material (potential reservoir of matter in the lump) and on the cooling history.  The value of $m$ and  of the mass ratio $q=m/M$ in the binary (with $M$ the heavier of the two stars) remains  unpredictable to us at this level,  so we may just depict three possible scenarios for the formation and evolution of this evanescent binary: case (A) when the instability sets after few hundreds of milliseconds or a second after core bounce and binary formation occurs so that $m\\simgreat 0.7\\msun$ (implying a pre-existing iron core of  large mass if the primary is as massive as 1.4$\\msun$); case (B) when a lighter star forms around a main body with $m\\simeq 0.2-0.4\\msun$ after several tens of seconds from core bounce; case (C) when cooling is sufficiently advanced that the minimum mass approaches its asymptotic value and the binary  can have $m\\simless 0.2\\msun$.    The magnitude of a natal kick in (A) is difficult to estimate and we must await for realistic simulations of core collapse.  Values of $m/M$ above a given threshold in a binary are known to lead to unstable mass transfer and thus to final coalescence.  We note that the evolution in this case might be similar to that described in coalescing neutron star binaries with slightly unequal masses (see Rosswog et al. 2000) where it has been shown that large kicks can be acquired as a consequence of mass loss via a wind. Case (B) and (C) will be explored in this paper. Imshennik \\& Popov (1998) estimated neutron star recoil speeds of order $1500-2000\\kms$ for scenario (C).  In cases (B) and (C), one factor that could lower the final neutron star speed is the finite velocity of the ejecta.  Colpi, Shapiro \\& Teukolsky (1993) have shown that in the dynamical phase of the explosion the ejecta can attain speeds varying from 10,000 to $\\sim 50,000 \\kms$ (the upper bound being related to the value of the binding energy of the star at the minimum mass relative to a dispersed state of iron). These speeds are close to the escape velocity from the binary and therefore the final kick imparted to the remaining neutron star may be influenced substantially by gravitational deflection of the ejecta, and velocity phase averaging during the explosion. A further threat would be a rapid decay of the orbital separation due to unstable mass transfer that may lead to final coalescence, and emission of gravitational waves.  In $\\S 3$ we will describe the gravitational bending of the ejecta, while in $\\S 4$ we will study initial conditions and, subsequently, orbit decay and mass exchange, for cases (B) and (C).  We will then give an estimate of the kick speed including phase-velocity averaging. ",
        "conclusions": ""
    },
    "0207/astro-ph0207111_arXiv.txt": {
        "abstract": "A semi-analytic model of cluster cooling flows is presented. The model assumes that episodic nuclear activity followed by radiative cooling without mass-dropout cycles the cluster gas between a relatively homogeneous, nearly isothermal post-outburst state and a cuspy configuration in which a cooling catastrophe initiates the next nuclear outburst. Fitting the model to {\\it Chandra\\/} data for the Hydra cluster, a lower limit of $284\\Myr$ until the next outburst of Hydra A is derived. Density, temperature and emission-measure profiles at several times prior to the cooling catastrophe are presented. It proves possible to fit the mass $M(\\sigma)$ with entropy index $P\\rho^{-\\gamma}$ less than $\\sigma$ to a simple power-law form, which is almost invariant as the cluster cools. We show that radiative cooling automatically establishes this power-law form if the entropy index was constant throughout the cluster gas at some early epoch or after an AGN activity cycle. To high precision, the central value of $\\sigma$ decreases linearly in time. The fraction of clusters in a magnitude-limited sample that have gas cooler than $T$ is calculated, and is shown to be small for $T=2\\keV$. Similarly, only 1 percent of clusters in such a sample contain gas with $P\\rho^{-\\gamma} < 2\\keV\\cm^2$. Entropy production in shocks is shown to be small. The entropy that is radiated from the cluster can be replaced if a few percent of the cluster gas passes through bubbles heated during an outburst of the AGN. ",
        "introduction": "Observations from the {\\it Chandra\\/} and {\\it XMM-Newton\\/} missions have shown that the intergalactic media of cooling-flow clusters is not a multiphase medium of the type required by the theory of cooling flows that has been widely accepted for over a decade \\citep[e.g.][]{af94}. Moreover, when these data are taken together with earlier radio and optical data, it is now clear that over the lifetimes of these systems very little gas has cooled to temperatures much below the virial temperature.  Consequently, these systems cannot be the scenes of a quiescent steady inflow regulated by distributed `mass dropout' as that theory postulated. In the light of the new observations \\citep{bbk01,mp01,mwn01,mkp02,mbf02}, increasing numbers of workers \\citep[e.g.][and references therein]{bmc02} are accepting the view that catastrophic cooling of intergalactic gas to very low temperatures takes place only at very small radii, and that on larger scales the intergalactic medium is periodically reheated by an active galactic nucleus (AGN) that is fed by central mass dropout. This general picture has been argued for several times over the last decade \\citep{tb93,bt95,jb96,jb99,co97,jb01}, but the physics involved is complex because the problem is inherently three-dimensional and unsteady, and many details remain to be filled in now that clearer observations are becoming available.  Recently, a number of authors have re-investigated the possibility that heat transported from the cluster outskirts to the centre by thermal conduction could contribute significantly to re-plenishing radiated energy \\citep{vsf02,fvm02,zn02,rb02}. The classical argument against thermal conduction remains that there is a narrow range of conductivity in which conduction is numerically significant but fails to eliminate the cooling region entirely \\citep{BinneyC,MeiksCf}. Moreover, conductivities that predict significant conductive heat input to dense regions now, are probably incompatible with the existence of cooling-flow clusters since they would have caused cluster gas to evaporate early in the universe \\citep{al02}. In this paper we model the effects of AGN alone, and neglect thermal conduction.  Although the mechanism by which the AGN heats the IGM is unclear -- possibilities include the impact on ambient thermal plasma of collimated outflows from the AGN \\citep{hrb98,ka98b,jb99} and inverse Compton scattering by thermal plasma of hard photons from the AGN \\citep{co97,co01} -- features have been observed near the centres of several clusters that are almost certainly bubbles of hot, low-density plasma that have been created by the AGN.  Simulations \\citep{cbkbf00,ssb01,bk01,qbb01,bkc01} suggest that these rise through the cluster's gravitational potential well on a dynamical timescale. The details of how a rising bubble mixes with surrounding gas and disperses its energy around the intracluster medium cannot be securely deduced from simulations, because they depend on what happens on very small scales at the edge of a bubble. The general idea that heating by AGN results in a kind of convection in the ICM is probably correct, however. The purpose of this note is to present a simple semi-analytical model of the life-cycle of a typical cluster that this process drives.  Section 2 defines the model, fits  its initial condition to {\\it Chandra\\/} data for the Hydra cluster, and shows its observable characteristics at a number of later epochs. Section 3 discusses the creation of entropy during an outburst of the AGN and estimates the fraction of the cluster gas that an outburst processes  through bubbles. Section 4 calculates the distribution of cluster-centre temperatures in a magnitude-limited sample of clusters and shows that temperatures below $\\sim2\\keV$ will very rarely be detected. Section 5 similarly calculates the distribution of the minimum values of the specific entropy that will be detected in a survey. Section 6 sums up. Throughout we assume a flat cosmology with $\\Omega_\\Lambda=0.7$ and $H_0=65\\kms\\Mpc^{-1}$.  \\begin{figure} \\centerline{\\psfig{file=MC667_fig1.ps,width=\\hsize}} \\caption{The function $M(\\sigma)$ for the Hydra cluster estimated from the data of \\citet{dnm01}.  The units of $\\sigma$ are $ 10^{31} \\cm^4\\g^{-2/3}\\s^{-2}$ and in these units $\\sigma_0=17$. The dotted line shows equation (\\ref{mass}) with $\\epsilon =1.48$.\\label{hydras}} \\end{figure} ",
        "conclusions": "In the Hydra cluster we find that the distribution of mass in entropy index $\\sigma$ has a simple power-law form. We find that this form is to a good approximation preserved when the cluster gas cools in hydrostatic equilibrium. Furthermore, the activity of an AGN at the cluster centre drives convective flows in the cluster gas. Continued convetion would eventually lead to a cluster with gas of a constant entropy throughout its entire volume. We show that starting from such a constant entropy distribution, the power-law form observed in the Hydra cluster arises naturally as the cluster gas cools. Thus the discovery that $M(\\sigma)$ for Hydra fits a power-law is no accident, but would have been discovered if we observed Hydra at an earlier or a later epoch.  We find that the central value of the entropy index, $\\sigma_0$, is a linear function of time until right up to the final cooling catastrophe that provokes an outburst of the central galactic nucleus. These results enable one to calculate the dynamics of a cooling flow with remarkable ease.  We estimate that in the absence of heat sources Hydra has about $280\\Myr$ to run before there is a central cooling catastrophe. This result places a lower bound on the time between nuclear outbursts. The actual inter-outburst time is likely to be  longer both because there is residual heating by supernovae/galactic winds and a low-level AGN, and because we have no reason to suppose that Hydra is in its immediate post-outburst state; it has probably been cooling for some time.  In trying to understand a cooling flow it is useful to focus on the system's entropy at least as much as on its energy. Between outbursts the story as regards entropy is simple: it is being radiated at a rate that can be readily calculated from the X-ray brightness profile. The story regarding energy is much more complicated: the gas losses energy  radiatively but recovers some through the work done by both the surrounding IGM and the intracluster gravitational field. Our premise is that the nuclear outburst that is provoked by each cooling catastrophe restores the cluster gas to essentially the same state that it had immediately after the previous outburst. This premise requires that the entropy created by irreversible processes during each outburst is equal to the entropy radiated between outbursts.  We assume that the restructuring of the cluster gas during an outburst is effected by bubbles of plasma. We calculate both the entropy created when a bubble is inflated, and that created as it mixes in with, and heats, the bulk of the intracluster gas. The relative sizes of these two entropy sources is a weak function of the mass fraction that passes through bubbles during an outburst, but is typically of order 3--5. If a few percent of the cluster gas passes through bubbles during an outburst, the time between outbursts is a few hundred Myr, as the observations suggest.  When gas cools from temperatures $\\gta3\\keV$ at which bremsstrahlung is the dominant cooling process, it spends only a very small fraction of the total cooling time at $T<1\\keV$. In the case of cluster gas this effect is magnified by the large density contrast between the cluster centre and the half-mass radius. Consequently, not only does gas at $T<1\\keV$ exist only for a very small fraction of the inter-outburst time, but it is confined to an extremely small fraction of the total volume. We have used our simple models of cooling flows to calculate the fraction of clusters in a magnitude-limited sample in which gas cooler than temperature $T$ would be detectable. We find that without pointed observations gas cooler than $2\\keV$ would be found in only $\\sim10^{-3}$ of clusters; with pointed observations ten times more sensitive than the discovery survey, such cool gas would be detected in $\\sim4$ percent of clusters. Thus, we should not be surprised to find that gas cooler than $\\sim1\\keV$ has not been detected in clusters observed by {\\it XMM\\/} and {\\it Chandra}.  Similarly, we predict that pointed observations of a complete sample of clusters would detect gas at entropies below $\\sigma\\sim2\\keV\\cm^2$ in only $\\sim1$ percent of clusters because low-entropy gas appears only fleetingly and in small volumes. This lower limit on the entropy is at a value of $\\sigma$ that is a factor $\\sim50$ lower than the `entropy floor' claimed for intergalactic gas from X-rays studies of low-mass clusters at radii well outside the region where radiative cooling is important.  However, with the spectral and spatial resolution of {\\it XMM\\/} and {\\it Chandra} the lower entropy limit should be detectable in samples of galaxy clusters."
    },
    "0207/astro-ph0207261_arXiv.txt": {
        "abstract": "Even though the technology of adaptive optics (AO) is rapidly maturing, calibration of the resulting images remains a major challenge. The AO point-spread function (PSF) changes quickly both in time and position on the sky. In a typical observation the star used for guiding will be separated from the scientific target by 10{\\arcsec} to 30{\\arcsec}. This is sufficient separation to render images of the guide star by themselves nearly useless in characterizing the PSF at the off-axis target position. A semi-empirical technique is described that improves the determination of the AO off-axis PSF. The method uses calibration images of dense star fields to determine the change in PSF with field position. It then uses this information to correct contemporaneous images of the guide star to produce a PSF that is more accurate for both the target position and the time of a scientific observation. We report on tests of the method using natural-guide-star AO systems on the Canada-France-Hawaii Telescope and Lick Observatory Shane Telescope, augmented by simple atmospheric computer simulations. At 25{\\arcsec} off-axis, predicting the PSF full width at half maximum using only information about the guide star results in an error of 60\\%. Using an image of a dense star field lowers this error to 33\\%, and our method, which also folds in information about the on-axis PSF, further decreases the error to 19{\\%}. ",
        "introduction": "A key ingredient to successful analysis of images produced with the aid of an adaptive optics (AO) system is the determination of the point-spread function (PSF). The application dictates how precisely the PSF must be determined. For example, limiting the uncertainty in PSF-fitting photometry in a crowded star field to only a few percent will demand very high accuracy in knowledge of the Strehl ratio. Another example, but a case where a good reference star is less likely to be found, is the subtraction of the point-like core from the image of a quasar-host or radio galaxy \\citep{Hutchings1998, Hutchings1999, Steinbring2002}. In this case Strehl ratio may be low, and the AO PSF approximately Gaussian. For a fixed volume, the peak of a 2-dimensional circular Gaussian is inversely proportional to the square of its full-width at half maximum (FWHM). Therefore, reducing the FWHM by, say, 25\\% will drive the peak up by 78\\%. A 25\\% underestimate of the QSO core FWHM will have a detrimental effect on the quality of fit. In practice, accuracy of only 10{\\%} to 20{\\%} in PSF FWHM is sufficient for detection of the host galaxy.  The task of determining the PSF is made difficult by three factors. First, the performance of the AO system depends on the brightness of the star which it uses to measure correction - the ``guide star.'' Second, the delivered correction depends directly on quickly changing seeing conditions. Third, the AO PSF is strongly field dependent (anisoplanatic). The first factor requires that scientific observations and calibration measurements employ either the same guide star or ones of similar brightness. Choosing a fainter guide star for calibration may require, to preserve $S/N$, that the AO system obtain longer integration times per measurement. Phase corrugations change continuously as the wind moves turbulence across the telescope pupil, and therefore the lower bandwidth will result in poorer correction. The latter two factors are related to the degree of turbulence along the line of sight. The seeing scales linearly with the inverse of the coherence length of phase distortions, $r_0$ (Fried's parameter), which for observations at zenith angle $\\gamma$, are related to the atmospheric refractive-index structure ``constant'', $C_n^2$, and wavelength, $\\lambda$, by $$r_0(\\lambda, \\gamma)=0.185\\lambda^{6/5}\\cos^{3/5}{\\gamma} \\Big(\\int{C_n^2(h) dh}\\Big)^{-3/5}, \\eqno (1)$$ where $h$ is height in the atmosphere (for review papers see Roddier 1981, 1999; Beckers 1983). Thus a measurement of the PSF needs to be simultaneous with the target observation or at least sample the same variation in $r_0$. Differences in phase distortion along different lines of sight, exacerbated by a $C_n^2(h)$ profile skewed to high altitude, will dramatically degrade both Strehl ratio and FWHM with increasing telescope offset from the guide star. A separation of as little as 10{\\arcsec} is sufficient to render images of the guide star unusable as a PSF estimate for the target. The effect of tilt correlation also causes the off-axis PSF to be anisotropic.  For any significant offset the amount of beam overlap between guide and PSF star in the tangential direction (perpendicular to the offset) will be greater than for the sagittal direction (along the offset). The lower correlation of aberrations in the sagittal direction causes the PSF to be elongated towards the guide star (see e.g. McClure et al. 1991,  Voitsekhovich \\& Bara 1999).  Limited AO system bandwidth may also elongate the PSF in the direction of the prevailing wind. Thus the PSF can depend on position angle in the frame, and this requires that the PSF measurement be made near the target position as well. Figure~\\ref{figure_star_onaxis_offaxis} illustrates the dramatic spatial dependence of the PSF. It shows natural-guide-star AO observations of two well-separated stars, taken simultaneously  (these data are discussed later, in Section \\ref{cfht_pueo}). The right-hand panel is an image of the guide star and the left-hand panel is an image of a star at an offset of 20{\\arcsec}. Both images have a field of view (FOV) of roughly 2{\\arcsec} $\\times$ 2{\\arcsec}, and the direction to the guide star is indicated by the arrow in the left-hand panel. The on-axis image has a high Strehl ratio, greater than 30\\%, but the off-axis Strehl ratio is closer to 10\\%. Note the degradation of the off-axis PSF and its elongation towards the guide star.  If a sufficiently bright star were within a few arcseconds of the target, the task of PSF measurement would be greatly simplified. The observer could use this ``PSF star'' as a nearby and simultaneous measurement of the PSF during the target observation. Unfortunately, the probability that a suitably bright PSF reference will be available is very small. Furthermore, some AO systems operate with imagers originally designed for uncompensated seeing, and their new optics have sacrificed field of view in order to obtain Nyquist sampling. For example, the small FOV such as the 4{\\arcsec} $\\times$ 4{\\arcsec} provided by the $256\\times256$ pixel Keck AO Near-Infrared Spectrometer (NIRSPEC) imager makes a fortuitous observation of both target and PSF star even more unlikely.  When faint extended objects are observed with natural guide-star AO, they will always be offset from the guide star. For a large FOV imager such as the Canada-France-Hawaii Telescope (CFHT) $1024\\times1024$ pixel infrared camera (KIR), the guide star is often within the 36{\\arcsec} $\\times$ 36{\\arcsec} FOV of the detector. This large FOV has the advantage of ensuring that a bright star will be available to characterize the on-axis PSF. Unfortunately, target integrations are typically long enough to leave a saturated image of the star. Still, several very short calibration exposures could be interleaved with those observations. An average of these would provide an unsaturated image of the star to characterize the on-axis PSF.  But even without images of the guide star, one can still determine the on-axis PSF. Consider a wave-front with corrugations consistent with a Kolmogorov spectrum. Now, consider perfect correction of the wave-front phase over an unobstructed pupil. An analytic expression can be found that gives the improvement in Strehl ratio due to successive removal of an increasing number of Zernike polynomials from the aberrated wave-front \\citep{Noll1976}. For example, removing the first 3 Zernike modes (excluding piston) corresponds to a reduction in the residual mean square wave-front distortion by a factor of 5. In a real AO system, these residual errors can in principle be determined from wave-front sensor measurements and thus used to predict the delivered on-axis PSF. The wavelengths of sensing and correction may also be different. The light from the guide star is usually split; optical light is sent to the wave-front sensor while correction is applied to the instrument light path in the near-infrared. \\citet{Veran1997} have exploited this and developed software that uses wave-front information to predict the on-axis PSF in the near-infrared for the CFHT AO system.  Knowing the on-axis PSF will not give the PSF at the target position. However, an estimate of the change in the PSF with increasing offset can be determined theoretically. For the case of a single thin turbulence layer at height $h$, infinite telescope aperture, and perfect correction of an infinite number of modes, the off-axis Strehl ratio, $S$, falls as \\citep{Roddier1999} $$S\\simeq\\exp[-(\\theta/\\theta_0)^{5/3}] \\eqno (2)$$ where $\\theta$ is the offset angle on the sky, and $$\\theta_0=0.314{{r_0 \\cos{\\gamma}}\\over{h}} \\eqno(3)$$ is the isoplanatic angle. For $r_0(0.5 \\mu{\\rm m})=10$ cm and $h=2$ km this gives $\\theta_0=19${\\arcsec} for $K$-band observations at zenith. That is, the PSF at an offset of 19{\\arcsec} will have a Strehl ratio degraded by a factor of $e$ from the on-axis value.  Determining the real AO off-axis PSF is more complicated. A real $C_n^2(h)$ profile will contain significant power at low altitude (dome seeing, ground-layer), but the more important contribution to anisoplanicity occurs at higher altitude. For example, measurements using SCIntillation Detection And Ranging (SCIDAR) show that a single thin high-altitude layer typically dominates the nighttime free-atmosphere $C_n^2(h)$ profile above Mauna Kea \\citep{Racine1995}. These observations predate the use of so-called generalized SCIDAR, however, and thus do not give a true picture of the conditions at the ground layer. But, evidence from generalized SCIDAR measurements at other observatories suggest that dominance by at most a few high layers may occur elsewhere as well (cf. Klueckers et al. 1998). Equation 3 suggests that the isoplanatic angle is inversely proportional to the height of the dominant layer. In reality, the off-axis Strehl ratio will depend on the mean height of the atmospheric turbulence, $\\bar{h}$, and on the equivalent number of corrected Zernike modes, $N$. Higher order modes will decorrelate more quickly than low order ones, resulting in a smaller isoplanatic angle (see e.g. Roddier 1999). That is, equation 2 should be replaced by \\citep{Flicker2002} $$S\\simeq\\exp[-(\\theta/\\theta_N)^{5/3}], \\eqno (4)$$ with the exponent tending towards 2 for very partial compensation \\citep{Roddier1993}. The ``instrumental'' isoplanatic angle is then given by $$\\theta_N=0.314{{r_0 \\cos{\\gamma}}\\over{\\bar{h}}}, \\eqno(5)$$ where \\citep{Fried1982} $$\\bar{h}=\\sec{\\gamma}\\Big(\\int{C_n^2(h)h^{5/3}dh}\\Big/{\\int{C_n^2(h)dh}}\\Big)^{-3/5}. \\eqno(6)$$ Now appropriately modified, equations 4 and 5 will give an estimate of the isoplanatic angle and layer height from measurements of the off-axis Strehl ratio. In practice, these Strehl ratio measurements could be made with a calibration observation of a suitably dense star field. Unfortunately, the Strehl ratio, both on-axis and off-axis, will vary with the seeing. In fact, since the off-axis PSF depends on not just the quickly varying $r_0$, but also on an evolving $C_n^2$ profile, it will vary even more rapidly, and with greater amplitude \\citep{Rigaut1999}. Thus, if the image is to be obtained as a mosaic, it must be completed quickly, before changes in $r_0$ or $\\bar{h}$ can introduce objectionable inhomogeneities. This calibration measurement, plus knowledge of the on-axis PSF, could be used to predict the off-axis PSF as long as there is no significant variation in the height of the dominant turbulence layer(s) or change in the equivalent number of compensated modes.  It should be stressed that the off-axis calibration image contains degenerate information about $\\bar{h}$ and $N$. Furthermore, if one does not know $r_0$ during the on-axis PSF measurement, its value of Strehl ratio is also degenerate, but in $r_0$ and $N$.  That is, as $r_0$ changes so does $N$, which in turn affects the isoplanatic angle. Although the value of $r_0$ can be determined from the AO system wavefront-sensor signals, thus relieving the on-axis PSF ambiguity, an independent measure of $\\bar{h}$ is more difficult to obtain.  Taking SCIDAR measurements during the AO observations could provide this information, but at substantial cost and effort.  However, even with the degeneracy one might obtain useful results. If a uniform and low value of $N$ were maintained, the off-axis PSF would change very gradually with offset. Figure~\\ref{figure_star_onaxis_offaxis} illustrates the case where Strehl ratios are high, and small changes in $r_0$ will dramatically affect not only the on-axis Strehl ratio, but also the spatial variation in the PSF.  A ``first-order'' method of determining the off-axis PSF variation - one that does not account for the variation in $N$ - may be less successful here.  In response to a need to obtain reliable off-axis PSFs for the processing of faint galaxies, one of us (ES) developed a means of providing PSF information at various field positions using calibration observations of dense star fields \\citep{Hutchings1998, Hutchings1999, Steinbring1998, Steinbring2001, Steinbring2002}. In this method, the observer obtains images of a dense star field once or at most a few times each night; the core of a globular cluster works best. The AO system is guided on a star of similar brightness to the one used in the target observations. Frequently, the required PSF is at an off-axis position outside the FOV of the camera, so a mosaic of the star field is generated by repositioning the telescope at successively greater off-axis distances. The goal is to obtain an unsaturated image of the guide star as well as a roughly uniform distribution of stars at many off-axis positions. Now, although the guide star in the star-field calibration image is chosen to be similar in brightness to the one used in the target observation, it is very unlikely that $r_0$ and, thus, the delivered PSF is identical. Therefore, the full PSF reconstruction technique must fold in information about the on-axis PSF {\\it during} the scientific observation with the star-field mosaic data.  Our method to deal with time dependence is to deconvolve the image of the dense star field with the image of the guide star in that field. We assume that the effects of field-dependent aberrations in the camera optics are small compared to anisoplanicity. We would expect, if the full image is deconvolved with a sub-image of just the guide star until the entire flux is contained within one pixel, that re-convolution with that same sub-image will reconstitute the original image. That is, the PSF at each of the off-axis star positions will be restored. Therefore, we should be able to take the deconvolved image of the star field and convolve it {\\it with the image of the guide star from a target observation}. The result would be the image of the star field as if it were observed under the natural seeing conditions at the time of the target observation. In this way, the image of the deconvolved star field contains the differences between the on-axis and off-axis PSFs. It is a map of the convolution kernel necessary to restore the off-axis PSFs over the field and will be referred to here as the PSF kernel map. Equations 4 and 5 suggest that the accuracy of predictions with the kernel map will decrease if the values of $\\bar{h}$ and $N$ during construction of the mosaic differ significantly from those during the target observation. To restate this another way, the off-axis PSF cannot be exactly separated into an on-axis PSF which depends only on $r_0$, and an off-axis kernel which depends only on $\\bar{h}$. Even so, we propose that the PSF kernel map is a useful approximation and that it provides a plausible method for predicting the off-axis PSF. It is not a mathematically exact method, and thus the results may depend on the deconvolution technique, number of iterations, convergence crition, among other things. For example, deconvolution of the guide-star to a spike of one pixel in width may lead to a spatial-frequency finite-support problem. One way to remedy this would be to convolve with the updated on-axis PSF first, and then deconvolve with the guide star from the calibration field. However, we will show that variation in the atmospheric conditions during observation of the star field can account for most of the uncertainty in the results, and thus for our demonstration the details of the deconvolution are not as important.  Observation of the calibration field is difficult and time-consuming, and must be a consideration in the application of the method. If only a small number of targets are observed, i.e. at only a few offset angles, a more practical method may be to measure the kernel with binaries of the appropriate separations. This may also allow for the interleaving of calibration and scientific observations, which the scarcity of globular clusters makes more difficult.  We will show, however, that even a single calibration observation taken during the night is helpful.  It may be that the problem is sufficiently well constrained that if one knows the on-axis Strehl and the Strehl degradation at one point, a model dependent on $N$ and $\\bar{h}$ will be sufficient for reconstruction of the kernel map, but we leave that work for a later paper.  A theoretical expression for PSF anisoplanatism has been derived by \\citet{Fusco2000}, but needs an independent measurement of $C_n^2$ (in their example case, from thermosonde data) to provide the kernel. In this paper we focus on what results can be obtained from just a calibration image. A benefit of a semi-empirical approach is that it will apply to all AO systems.  Conjugation of the AO correction to high altitude layers, such as employed in Gemini Altair, will reduce variation in the off-axis PSF, and should improve the results from our method. Multi-conjugate AO (MCAO) should lessen PSF variation even more. In MCAO, PSF uniformity is more dependent on system geometry than atmospheric conditions, and a direct measurement of the PSF kernel map may prove to be a very successful tool. This may allow field dependent deconvolution, for example.  Broad-band AO observations of dense star fields are frequently made by other observers, but almost always for scientific observations of the stars in those fields and not for calibration purposes. They are generally unusable for our analysis for one or more of the following reasons: 1. The image of the guide star is saturated. 2. The FOV of the mosaic is too small or contains too few stars. 3. Typically, a long time elapses between the observation of the guide star and the other parts of the mosaic. Variations in seeing thus render any measurement of anisoplanicity unreliable.  The work discussed here involves the observation and analysis of dense star fields for the sole purpose of calibrating the AO off-axis PSF. The observations and data reductions are outlined in Section~\\ref{observations}.  Computer simulations were used to further characterize the data and the PSF prediction method was applied to them. These analyses are found in Section~\\ref{analysis} along with a discussion of the uncertainties in the predictions. Our conclusions follow in Section~\\ref{conclusions}. ",
        "conclusions": "We have discussed observations of dense star fields for the purpose of calibrating AO off-axis PSFs at CFHT and Lick Observatory. The results of our simulations suggest that a simple model with a single thin turbulence layer is a useful model for our data. For both the \\objectname{M 5} and \\objectname{M 15} observations, a dominant layer at between 2 and 5 km above the telescope and a low number of corrected modes reproduces the anisoplanicity seen in the data.  A semi-empirical technique was applied that utilizes these findings to improve the prediction of AO off-axis PSFs. The results show this simple method reduces error in the prediction of the basic radial variation in the PSF. We have insufficient data to determine if it can predict the suspected variation in anisoplanicity due to position angle about the guide star. The observations of \\objectname{M 15} suggest that a single calibration observation of a dense stellar field is a useful prediction of anisoplanicity over the course of a night, or even for a different night, at least for the atmospheric conditions present at Lick Observatory. Errors in prediction of FWHM with this method are approximately 19\\% compared to 33\\% without any on-axis information, and 60\\% without any anisoplanatic correction at all. For maximum accuracy the number of corrected modes should be similar, or larger, in the calibration image.  The accuracy of determining anisoplanicity from a calibration image is limited because the effects of $r_0$, $\\bar{h}$, and $N$ are intertwined. Better results could be realized by separating these effects through independent measurement of $C_n^2$.  The simulations help us understand why the method is successful, if only modestly. The dominant-layer height evidently does not change during a night or from night to night significantly enough to dramatically alter the isoplanatic patch delivered by the system. Obtaining a calibration image that measures this anisoplanicity is, however, a difficult task. A mosaiced image must be completed quickly, before variations in delivered Strehl ratio ruin the measurement. Each of our mosaics was constructed in under 10 minutes, but variation in Strehl ratio by 15\\% to 50\\% during that time ultimately limited the accuracy of the results.  In September 2001 we again used the Lick AO system but this time employed its laser beacon to acquire more calibration observations of dense star fields. Instead of narrow strips, coverage of a large box-shaped region of sky was obtained. This will allow us to characterize anisoplanatic behavior both radially away from and azimuthally about the guide star. Future observations would also benefit from confirmatory SCIDAR measurements of the $C_n^2$ profile."
    },
    "0207/astro-ph0207057_arXiv.txt": {
        "abstract": "{ Intraday Variability of compact extragalactic radio sources can be interpreted as quenched scintillation due to turbulent density fluctuations of the nearby ionized interstellar medium. We demonstrate that the statistical analysis of IDV time series contains both information about sub-structure of the source on the scale of several 10 $\\mu$as and about the turbulent state of the ISM. The source structure and ISM properties cannot be disentangled using IDV observations alone. A comparison with the known morphology of the `local bubble' and the turbulent ISM from pulsar observations constrains possible source models. We further argue that earth orbit synthesis fails for  non-stationary relativistic sources and no reliable 2D-Fourier reconstruction is possible. } ",
        "introduction": "Intraday Variability (IDV) (see Wagner \\& Witzel (\\cite{Wagner}) for a review) is a common phenomenon of flat-spectrum radio cores in quasars and BL Lacs. From light travel time arguments the observed brightness temperature of IDV sources are in the range of $10^{16 \\ldots 21}$~K and far in excess of the inverse Compton limit. The suggested intrinsic explanation require either extreme Doppler boosting or special source geometries. Gravitational micro-lensing and scintillation in the ISM of our galaxy are also discussed as possible propagation effects causing variability.  Two classes of IDV sources\\footnote{Besides normal IDV two extremely fast sources PKS 0405-385 (Kedziora-Chudczer et al. \\cite{Kedziora-Chudczer}) and J1819+3845 (Dennett-Thorpe \\& de Bruyn \\cite{Dennett-Thorpe}) have been found with very short time-scales of less than $1$ hour.} are distinguished according to their structure functions ($S\\!F$). One shows a continuous increase of $S\\!F$ with time-lag, while the other reaches a `plateau' at a well defined time-scale $t_\\mathrm{IDV}$ (see Fig.~1).   In this paper we will discuss the scintillation hypothesis for the `plateau' class. Scintillation is caused by turbulence in the ISM, which is a stochastic process and the resulting time series are best analysed in terms of first order structure functions. Following Simonetti, Cordes \\& Heeschen (\\cite{Simonetti}) we define the discrete structure function for time series $f(t_i)$ by \\begin{equation} S\\!F(\\tau_j) = N^{-1}_{ij}\\sum_{i=1}^{n} w(i) w(i+j) \\left[f(t_i) - f(t_i + \\tau_j)\\right]^2 \\,. \\end{equation} Here $N_{ij}$ is the normalisation; $w(i)$ are weighting functions, so that $w(i) w(i+j) >0$, if a measurement at $t_i$ and another measurement at $t_i + \\tau_j$ were obtained. The weighting function also accounts for bining of unevenly sampled data. The structure function is related to the autocorrelation $\\rho(\\tau)$ by $S\\!F(\\tau) = 2[\\rho(0) - \\rho(\\tau)]$.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{paper_fig1.eps} \\caption{Structure function for IDV in total intensity for the quasar 0917+624. The $S\\!F \\propto \\tau^{\\alpha}$ has a slope of $\\alpha \\approx 1$, a 'plateau' with $S\\!F = 2 m^2$ for long time-lags and a characteristic timescale for decorrelation $t_\\mathrm{IDV}$ at the first maximum. \\label{fig:0917_sf} } \\end{figure} ",
        "conclusions": "Quenched scintillation gives a consistent picture of Intraday Variability for most IDV sources. The time-scale of about 1~day fits with scattering in the ionized gas of the wall of the local bubble of a moderately boosted synchrotron source with brightness temperatures below the inverse Compton limit.  From the analytic expressions for modulation index $m$ and structure function at small time-lags presented here, it is possible to determine two of the four unknowns: distance to the scattering medium, scattering measure, velocity of the medium, and size of the source. In the slab model of the scattering medium together with a constant surface brightness of the source it is possible to determine the slope $\\beta \\approx 4$ of turbulent power spectrum from the structure function.  The structure of IDV sources beyond  the size of the most compact component cannot be uniquely determined from existing experiments. Because most IDV sources are intrinsically variable on time-scales of several months earth orbit synthesis cannot be reliably realized.  Extreme scattering events (Fiedler et al. \\cite{Fiedler}) show that the ISM is structured on scales of several AU down to 0.01 AU (Cim\\`o et al. \\cite{Cimo}) well into the scintillation domain discussed here. Turbulence in the ISM is therefore not homogeneous on scales of $v$ (3 months) = 1~AU. Changes of the long term variability pattern like the sudden end of IDV in 0917+624 (Fuhrmann et al. \\cite{Fuhrmann}) can be due to changes intrinsic to the source or in the scattering measure of the ISM."
    },
    "0207/astro-ph0207668_arXiv.txt": {
        "abstract": "We have performed 3D hydrodynamical simulations of FR-II radio sources in $\\beta$-profile cooling-flow clusters.  The effects of cooling of the cluster gas were incorporated into a modified version of the \\textsc{zeus-mp} code.  The simulations followed not only the active phase of the radio source, but also the long term behaviour for up to 2 Gyr after the jets of the radio source were switched off.  We find as expected that the radio source has a significant effect on the cooling flow while it is active, however we also find that the effects of the radio source on the cluster are long-lived.  A buoyancy driven convective flow is established as the remnants of the radio source rise through the cluster dragging material from the cluster core.  Although the central Mpc of the cluster reverts to having a cooling flow, this asymmetric convective flow is able to remove the cool gas accumulating at the cluster core and indeed there is a net outflow persisting for timescales of about an order of magnitude longer than the time for which the source is active or longer. The convective flow may also provide a mechanism to enhance the metallicity of the cluster gas at large cluster radii. ",
        "introduction": "There has been significant recent interest in the interaction between radio sources and clusters.  This has been driven by observations with both the \\textit{ROSAT} and \\textit{CHANDRA} satellites which have demonstrated clearly that the presence of a powerful FR-II radio source very significantly affects the ICM/IGM in the vicinity of the source.  For example, for the best studied cases of Cygnus A \\citep*{Carilli et al,Smith et al} and 3C84 in Perseus (\\citealt{Fabian et al2000}; \\citealt{Fabian et al2002}), cavities in the X-ray emitting gas coincident with the radio cocoon are clearly observed and for Cygnus A the effects of a strong bow shock can be clearly seen.  Recent reviews of the X-ray observational data on radio-source/cluster interactions are given in \\citet{McNamara2002} and \\citet{McNamara et al2000}.  This interaction between the radio source and the ICM has led a number of authors to consider whether radio sources could solve the so-called ``cooling-flow problem'' (e.g. \\citealt*{BinneyTabor, McNamara et al2000, Reynolds et al, Churazov et al, Quilis et al, Bohringer et al, RuszkowskiBegelman, BrighentiMathews, RHB}, hereafter RHB), in which the cool gas expected to accumulate at the cluster centre as a result of cooling is not detected.  A number of authors have discussed the heating of the ICM by radio sources in the context of the overall energy budget for the cluster (e.g. \\citealt{BinneyTabor, KaiserAlexander1999}; RHB; \\citealt{Alexander} and references therein).  \\citet{Churazov et al2002} discuss the energy balance between the cooling flow and the AGN on the basis that there is some feedback between the cooling flow (which results in matter accreting onto a supermassive black hole) and the resulting mechanical power of the AGN. This could also be responsible for the intermittency of the radio source as inferred in some observations \\citep{McNamara et al2000_1, Fabian et al2000, OwenEilek}. The energy input from powerful FR-II radio sources is certainly sufficient to have a significant effect on the cluster, however the lifetimes of FR-II radio sources are of order 100 Myr (e.g. \\citealt{AlexanderLeahy}), they are therefore transient events in the lifetime of a cooling flow \\citep{Fabian}.  An important question is therefore whether a radio source can have a long-term effect on the cluster; for this reason it is necessary to consider the evolution of the remnants of dead radio sources (i.e. after the jets cease energy injection into the lobes).  During these latter stages of evolution the dynamics of the remnants of the radio source will be dominated by buoyancy forces \\citep{GullNorthover}.  A number of studies have considered the evolution of buoyant plasma bubbles within a cluster.  For example \\citet*{SarazinBaumOdea} and \\citet{Churazov et al} developed models for 2A 0335+096 and M87 respectively where the observed radio and X-ray structures are modelled by buoyant bubbles of radio-emitting plasma which drag colder material that had been deposited by the cooling flow outwards from the cluster core.  These bubbles are Rayleigh-Taylor unstable and therefore form a mushroom cloud type structure \\citep{BruggenKaiser}. \\citet*{Saxton et al} have modelled the northern middle radio lobe of Centaurus A as a buoyant bubble, and simulations of buoyant gas in a cluster environment have been presented in 3D by \\citet{Bruggen et al} and in high resolution 2D by \\cite{BruggenKaiser2002}. In all of these studies gas was injected continuously near to the centre of the cluster with zero velocity and in pressure balance with the surrounding gas and the system allowed to evolve.  Recently RHB have used a more realistic model which involved setting up the initial radio plasma by simulating a jet; these conditions were then used as the starting point of a new simulation in which the jet activity was turned off.  Although the basic physical processes have been established by these studies, the details of the hydrodynamical evolution of a cooling flow cluster containing a radio source needs to be fully determined if we are to answer the question about the long-term effects of radio sources.  In this and a future paper we extend these studies.  We follow RHB in simulating a radio source evolving in a cluster so that the initial conditions are as realistic as is possible, however we use a fully 3D simulation and also include cooling of the ICM.  The radio source evolves into a cluster environment with an established cooling flow, it is then turned off after a simulation time corresponding to approximately 50 Myr and we then follow the evolution of the system for a further 2 Gyr. ",
        "conclusions": "\\label{sec:discussion}  As has been noted by many authors (e.g. \\citealt{Alexander} and references therein; RHB), radio sources are certainly sufficiently energetic to provide substantial heating to a cluster.  For example Perseus has an X-ray luminosity of approximately $10^{37}$ W and powerful radio sources have jet kinetic powers of order $10^{39}$ W, therefore if the radio source lasts for a few $10^7$ yr the energy input from a single powerful radio source is sufficient to provide all the energy to power the X-ray emission for a substantial fraction of a Hubble time.  However radio sources are short-lived events and it is necessary to consider more carefully the balance between heating rate and cooling rate.  Our results go some way to understanding the possible role of radio sources in cooling-flow clusters.  Our simulations have shown that the effect of a powerful radio source on its host cluster is to modify the dynamics of the cluster gas on timescales very much longer than the lifetime of the active phase of the radio source.  The buoyant remnants of the radio source drive a large scale convective flow which is efficient at removing gas from the cluster core on timescales comparable to a Hubble time. By this mechanism the coolest gas will be moved to larger cluster radii.  Such a flow would also provide a mechanism for metal enrichment of the cluster gas on large scales.  Observational support for these processes comes from the \\textit{CHANDRA} study of A133 \\citep{Fujita et al}; this cooling-flow cluster contains a ``radio relic'' which is connected to the host CD galaxy by a tongue of cooler X-ray emitting gas.  The authors consider the possibility that the X-ray structure is explained by uplifting by the buoyant radio relic and demonstrate that this is feasible.  They reject this possibility since the morphology of the cool X-ray emission does not agree with the model predictions of \\citet{Churazov et al} or \\citet{Bruggen et al}.  However, our simulations predict a structure in the cooler X-ray gas (Figure 2 a\\&c) which is remarkably similar to the \\textit{CHANDRA} results.  The structure in the X-ray emitting gas depends strongly on the relative power of the radio source and this will be explored in a forthcoming paper.  The simulations presented here although addressing the same problem as RHB provide important new insight into the evolution of radio-source containing clusters.  Firstly the buoyant remnants in the current simulations show less mixing than in the simulations by RHB; the result is that more of the available energy goes into large-scale motion establishing the buoyancy-driven convective flow we have described.  This is due in part to our simulations being 3D which permit non-axisymmetric Kelvin Helmholtz instabilities to develop which reduces the amount of mixing although we do see large-scale instabilities.  Unfortunately the price we pay for going to 3D is to loose resolution which will suppress high spatial frequency instabilities which would be efficient at mixing.  However it is possible that such modes may be suppressed in the real astrophysical system by the magnetic fields which must be present within the radio lobes.  The low grid resolution used here means that numerical diffusion could be a problem. Here we argue that such numerical effects are small in comparison with the overall behaviour observed in our simulations. Firstly, we consider the relative importance of numerical diffusion versus cooling by comparing the results of run 2 with a simulation employing the same parameters, but without cooling (run 2a). For run 2a, the cumulative mass flow through a spherical surface at $300~\\rm{kpc}$ radius was $2.3\\times 10^{12}~\\text{M}_{\\sun}$ compared to $1.0\\times 10^{12}~\\text{M}_{\\sun}$ for run 2. This value increases with time due to the convective effects of the rising remnants for run 2a, compared with run 2 where this value is decreasing with time as the cooling flow begins to dominate the behaviour. We are therefore confident that the effects observed in the simulation due to cooling are real, and are not being strongly affected out by numerical diffusion of energy out of the computational domain. A further check on the numerical accuracy is to compare results at differing resolutions. In a forthcoming paper we will present a 2D parameter space investigation, and also consider in detail the effects of varying resolutions on the simulations. These results show that the rate of energy loss through numerical diffusion is essentially independent of the resolution used for resolutions in the range $128\\times64\\rightarrow512\\times512$ and overall energy conservation in our simulations at all resolutions is as good as that of RHB."
    },
    "0207/astro-ph0207342_arXiv.txt": {
        "abstract": "In the present study, we analyze the effects of a flux of Alfv\\'en waves acting together with radiation pressure on grains as an acceleration mechanism of the wind of late-type stars. In the wind model we simulate the presence of grains through a strong damping of the waves, we use a non-isothermal profile for temperature coherent with grain formation theories. We examine the changes in the velocity profile of the wind and we show that if the grains are created in the region $1.1 < r/r_0 < 2.0$ their presence will affect the mass loss and terminal velocity. The model is applied to a K5 supergiant star and for Betelgeuse ($\\alpha \\, Ori$). ",
        "introduction": "Mass loss from late-type stars is evidenced by the blueshifted circumstellar absorption lines present in their spectra. Giant and supergiant stars eject mass at a rate of $\\sim 10^{-10} - 10^{-5} \\; M_{\\odot} \\, yr^{-1}$  with terminal velocities $u_{\\infty}  \\sim 10 - 50 \\; km \\, s^{-1}$, which is lower than the surface escape velocity ($v_{e0}$) (\\cite{deutsch56}; \\cite{weymann62}; \\cite{habing96}; \\cite{lamerscassinelli99}; \\cite{carpenter99}). An outward-directed flux of Alfv\\'en waves has been proposed for driving these winds (\\cite{belcherolbert75}; \\cite{haisch80}; \\cite{hartmannmacgregor80}; \\cite{jpo89a}; \\cite{macgregorcharbonneau94}; \\cite{charbonneau95}). For late-type giant and supergiant stars, this nonthermal energy flux must be deposited in the region of subsonic flow to explain the observed $u_\\infty \\; < \\; v_{e0}$ (\\cite{hartmannmacgregor80}; \\cite{leerholzer80}).  We have direct evidence for Alfv\\'en waves in the solar wind. Waves are observed with large perturbations in the magnetic field, with negligible density perturbations, indicating the presence of Alfv\\'en waves. Since the early observation of MHD waves in the solar wind (\\cite{belcherdavis71}), various authors have suggested that Alfv\\'en waves could be important in transfer momentum to the wind. We know from the solar coronal holes observations that the density is likely to fall off faster than $r^{-2}$ (e.g. \\cite{hollweg81}; \\cite{doschek97}), while the magnetic field strengths is proportional to $r^{-2}$ (\\cite{linsky87}). In our simulations we show that this behavior exists up to $r < 3r_0$. The energy flux per unit area transported by the wave, $\\phi_A$, is $\\phi_A \\simeq \\epsilon v_A \\simeq (1/2)\\rho_{0}{\\langle \\delta v^2 \\rangle} v_A$, where $\\epsilon$ is the energy density of the waves. This energy flux is constant, when there is no damping, or decreases due to damping, such that $\\epsilon$ decreases with $r$. Since the pressure associated with the wave is proportional to $\\epsilon$, then this pressure also decreases with $r$. The result of this is a pressure gradient that accelerates the gas (\\cite{castor81}).  Alfv\\'en waves have been used to explain the heating of stellar corona and the production of stellar winds in many regions of the Hertzsprung-Russell (HR) diagram. They are used to explain the solar wind acceleration (\\cite{jpo89b}; \\cite{jpoy94}), cool supergiant stars (\\cite{charbonneau95}; \\cite{airapetian00}) and winds in hot stars (\\cite{cass79}; \\cite{underhill83}; \\cite{sjpo93}). These waves can also be important in T Tauri variables (\\cite{hartmann82}; \\cite{jpo89c}; \\cite{vjpo00}). The origin of Alfv\\'en waves was suggested by \\cite{op86} to be the annihilation of twisted magnetic fields near the surface. These annihilation sites that could also explain the variable hot-spots observed in late-type stars (Gray 2000, 2001; \\cite{young00}). The waves can also be generated by convective motions inner in the envelope, that create perturbations in the magnetic fields.   On the other hand, in cool supergiant stars, the presence of grains are well know determined by infrared excesses (\\cite{woolf69}). However, the region of nucleation is not precisely determined. These grains are supposed to be the main responsible for radiation absorption, and also are supposed to contribute in the acceleration of the wind (e.g. \\cite{jura86}; \\cite{elitzur01}). The temperature profile of the wind for these stars is also not very well known, but by nucleation theories (\\cite{gail87}) the temperature range for grain formation is $400 \\, K < T < 2500 \\, K$.  Besides, the high luminosity of these stars have lead many authors to include the radiation pressure as a mechanism to explain the mass loss. MacGregor and Stencel (1992) studied the interaction between gas and dust in the curcumstellar envelope of late-type stars. In their analysis the grains and gas are well coupled if the grain radius is $\\sim 10^{-5} \\; cm$ while for small grains $< 5 \\times 10^{-6} \\; cm$, the collision momentum transfer is insufficient to induce expansion of the atmosphere as a whole.  Jatenco-Pereira and Opher (1989a) (hereafter JPO), presented a model for mass loss in a cool supergiant K5 star with $M = 16 M_{\\odot}$, where a flux of Alfv\\'en waves flux ($\\phi_{A}$) is the main mechanism for accelerating the wind. The model solves the equations along the magnetic field axis, and assumes (i) a divergent geometry for magnetic field lines, (ii) a pure plasma (without dust) and (iii) isothermal wind. Holzer, Fla and Leer (1983) argued that wind driven by Alfv\\'en waves from late-type stars only would be possible in an average ion-friction damping length $L_f \\sim 0.85 - 1.0 r_0$. Since $L_f$ is proportional to $P^2$ (where $P$ is the average Alfv\\'en period), this requires that nature fine-tune $P \\sim 1.77 \\times 10^4 \\; s$. For long wave period ($P > 10^4 \\; s$), the wave damping length is large and the ion-neutral friction damping is not important (\\cite{depontieu01}). On the other hand, Davila (1985) emphasized that on the surface of the Sun one might expect that boundary effects may be important, and a WKB approximation may not be valid whenever the MHD wavelength ($\\lambda$) becomes greater than $d \\sim 10^{10} \\; cm \\; \\sim 0.14 r_0$, the characteristic transverse dimension of a solar coronal hole, in despite of Usmanov et al. (2000) that used a WKB approximation for the solar wind, in comparison with Ulysses data. We assume an Alfv\\'en wave flux with an averaged wavelength $\\lambda < 0.5 r_0$ such that the WKB approximation is valid. JPO used a wave period $P > 10^4 \\; s$ and discussed other possible damping mechanism for Alfv\\'en waves such as nonlinear damping, resonant surface damping and turbulent damping. They show that these damping processes can produce acceptable winds in late-type giant stars. We extend the JPO model including a simulation of a strong damping of Alfv\\'en waves, possibly due to a grain interaction, and the effect of radiation pressure in the wind acceleration. The changes in the velocity profile of the wind are compared with the pure plasma and isothermal model.  In section 2, we briefly describe the JPO model. In section 3, we present the inclusion of grain interaction with Alfv\\'{e}n waves in a non-isothermal temperature profile and the radiation pressure on grains. In section 4, we apply this magnetic-radiation model in a K5 supergiant star and $\\alpha \\, Ori$ and discuss the results. In section 5 we summarize our conclusions. ",
        "conclusions": "In this work we present a model to study mass loss processes in late-type stars, using an Alfv\\'{e}n waves flux added to radiation pressure as accelerating mechanisms. We simulate the presence of grains in the sense of a strong damping of Alfv\\'en waves in the region $r > GFP$ (which is about $1.1r_0 < GFP < 2r_0$). In this region we added the radiation pressure as an acceleration mechanism of the wind. In accordance with grain formation theories and in agreement with observations, we took into account a temperature profile for the wind going from $\\sim 2000 \\; K$ above the photosphere, with a rapidly grown to $\\sim 10^{4} K$ at $r = 2r_{0}$ and a decreasing profile for $r > 2r_0$. The models show that as closer as to the surface the grain is formed, lower is the asymptotic velocity and that the wind velocity is not changed if the region of grain formation is located above $2r_0$. Although a number of hypotheses have been made this hybrid model can lead to a more realistic wind and gives more constraint in the physical parameters of the wind. The model was applied to a K5 supergiant star and Betelgeuse, resulting in values that agrees with observations. Further studies will include the physical parameters of the grains as well as their distribution in the wind.  \\smallskip"
    },
    "0207/astro-ph0207174_arXiv.txt": {
        "abstract": "Although quintessence models have many attractive cosmological features, they face two major difficulties. First, it has not yet been possible to find one which convincingly realizes the goal of explaining present-day cosmic acceleration generically using only attractor solutions. Second, quintessence has proven difficult to obtain within realistic microscopic theories, largely due to two major obstructions. Both of these difficulties are summarized in this article, together with a recent proposal for circumventing the second of them within a brane-world context. It is shown that this proposal leads to a broader class of dynamics for the quintessence field, in which its couplings slowly run (or: `walk') over cosmological time scales. The walking of the quintessence couplings opens up new possibilities for solving the first problem: that of obtaining acceptable transitions between attractor solutions. ",
        "introduction": "The discovery by cosmologists that the Universe is currently dominated by two distinct types of unknown forms of matter is a development with truly Copernican implications for our picture of the Universe as a whole. We have known for some time that visible matter likes to cluster on large scales into galaxies and galaxy clusters, with 90\\% or more of the mass of these objects consisting of an unknown nonbaryonic `dark matter' \\cite{DMCitations}. The more recent surprise was the discovery that this clustered dark matter itself makes up no more than 30\\% of the overall energy density of the Universe, with the remaining 70\\% apparently consisting of a different kind of unknown substance, sometimes called `dark energy' \\cite{DECitations}.  It is absolutely breathtaking that so little is known about these two most abundant forms of matter. What is known is usually expressed in terms of their equations of state, through the ratio of pressure to energy density, $w = p/\\rho$. The formation of galaxies and galaxy clusters by gravitational attraction appears to require the dark matter to be `cold', with $w$ close to zero. Similarly, the current acceleration which the universal expansion is undergoing indicates $w \\lsim -0.3$ for the dark energy.  Perhaps the simplest explanation for the dark energy is that it is simply the energy density of the quantum vacuum, since this satisfies $w = -1$ and is generically nonzero in realistic quantum field theories. The problem with this explanation is that the predicted vacuum energy is typically at least $10^{56}$ times larger than what is observed. The Cosmological Constant Problem \\cite{CCProblem} is the recognition that at present no way is known to naturally obtain a vacuum energy anywhere near the required size within a realistic microscopic theory.  Quintessence models \\cite{Quintessence} take a different tack to explain the dark energy. In these models the dark energy is attributed to the dynamics of a scalar field, $\\phi$, which is currently evolving in a cosmologically interesting way. Since $w = (K-V)/(K+V)$, where $K= \\frac12 \\dot\\phi^2$ and $V$ are the scalar's kinetic and potential energies, the condition that $w$ be negative requires the scalar's evolution at present must be slow in the sense that $K \\lsim V$. (The vacuum energy is a special case of this kind of evolution, where $K=0$.)  These models do not solve the Cosmological Constant Problem, since they do not provide natural reasons for $V$ to be currently so small, but they can (potentially) explain why an evolving scalar field could naturally have an energy density which is now so similar to that of other forms of matter, like the dark matter. They can do so, firstly because their equations of motion often admit `tracking' solutions, within which the scalar energy density closely follows (or tracks) the dominant energy density of the Universe as it evolves. Furthermore, the late-time evolution of the scalar field is often drawn to these `tracking' solutions for wide choices of initial conditions, because these solutions are also `attractors' for the scalar equations of motion.  Unfortunately, since these tracker solutions typically require the scalar itself not to be the dominant energy density, in order to become the dark energy the scalar must eventually leave the tracker solution. Although one might hope that this could also be naturally achieved -- such as being perhaps due to a transient behaviour due to the crossover from radiation to matter domination -- so far it has proven difficult to make a completely convincing cosmology along these lines.  The purpose of this article is to describe a new category of quintessence model, which could be called `Walking Quintessence' \\cite{ABRS1,ABRS2}, that may offer new ways to accomplish this crossover. They may do so because within these models the couplings of the scalar field slowly run (or walk) as the Universe evolves, and this walking may help facilitate the crossover between tracking solutions. What is remarkable is that these models were developed in an attempt to address a completely different set of (very serious) problems which arise once one tries to obtain quintessence from a realistic model of microscopic physics.  Since they play such an important part in the motivation of the models, the bulk of the article is devoted to these serious microphysical problems. The problems themselves are first summarized in the next section, followed by a description how they are addressed within the attractive brane-world picture which has emerged as a potential low-energy consequence of string theory. The results of a sample cosmology built on this model are then briefly presented, intended as first step towards a more systematic exploration of the cosmologies which are suggested by this class of models. ",
        "conclusions": ""
    },
    "0207/astro-ph0207497_arXiv.txt": {
        "abstract": "Using the technique of PSF-matched image subtraction, we have analyzed archival {\\em HST}/WFPC2 data to reveal details of at least two light-echo structures, including some unknown before now, around SN~1993J in the galaxy M81. In particular, we see one partial sheet of material 81~pc in front of the SN and tilted $\\sim60^\\circ$ relative to the disk plane of M81, and another 220~pc in front of the SN, roughly parallel to the disk. The inferred echoing material is consistent with the \\ion{H}{1} surface density detected in this region of M81's disk; however, these data imply a fragmented covering factor for the echoing structures. We discuss prospects for future (roughly annual) visits by {\\em HST} to image these and yet undiscovered echoes in the interstellar and circumstellar environment of SN~1993J. ",
        "introduction": "Supernova (SN) 1993J was a type II SN % \\citep{rip93} in the nearby galaxy NGC 3031 (M81), the closest SN seen in the past decade. It's spectrum quickly lost most of its hydrogen emission lines, indicating that most of the progenitor's hydrogen envelope had been stripped, and justifying classification of the SN as type IIb \\citep{nom93}. The lost material probably formed a circumstellar shell, as evidenced by emission in X-rays \\citep{zim94}, radio \\citep{bar94}, and narrow optical lines \\citep{bnt94,mat00}.  Since the SN exploded near a spiral arm in M81, it is expected to illuminate interstellar and circumstellar material in the form of light echoes. Such echoes have been reported for e.g.,\\ SN~1998bu \\citep{cap01}, SN~1991T \\citep{sch94,spa99} and SN~1987A \\citep{cro88a}. Light echoes from SN~1987A have revealed structures on interstellar and circumstellar scales \\citep[and references therein]{cro95,cro01}, which have been used to map the surrounding material in three dimensions and tie it to kinematic information \\citep{xu95,xu99,cro00}, offering unique insights into the history of the associated stars and gas.  Recently, \\citet{liu02} discovered a light echo around SN~1993J based on data taken on 2001 June 4 on the WF4 detector of the Wide Field Planetary Camera 2 (WFPC2) aboard the {\\em Hubble Space Telescope} ({\\em HST}). We have reanalyzed these now publicly available data, and others (\\S2), and have found yet another distinct echo. Using image subtraction and PSF-fitting techniques we have also produced detailed analyses (\\S3) of the structure and reflectivity of both echoing structures. These results already reveal intriguing details about the interstellar medium in M81, and are likely to continue to do so in the future. ",
        "conclusions": "Figure 1a shows the F555W image from 2001, and Figure 1b shows its difference from the F555W integration of 1995.  The 450W image from 2001, with stellar sources removed (\\S2), is shown in Figure 1c, and the difference image in F814W (with obviously lower signal-to-noise ratio) corresponding to panel (b) is show in Figure 1d.  These clearly show two echoes from 2001 (lighter shade) as well as a confusion due to poor subtraction within $\\sim 0\\farcs4$ of the SN.  The outermost echo is seen at radii extending from $\\theta=1\\farcs84$ to $1\\farcs95$ from the SN, at position angles 170$<$PA$<$290, with the largest radii near PA 225.  These correspond to distances from the SN sightline $r = \\theta D = 32.4$ to 34.3~pc ($D = 3.63$~Mpc: \\citealt{fre94}).  Since it is a light echo, one can compute the foreground distance (along the SN sightline) $z = r^2/2ct - ct/2 = 209$ to 235 pc, implying a tilt of about 37$^\\circ$ with the southwest side farther in front of the SN. An echo at the same $z$ distance would occur at $\\theta=0\\farcs89$ on 1995 January 31, and indeed there is some marginal indication of such a feature (dark in Figures 1b and 1d) between PA 190 and 260.  This is far from definite, but may indicate a shrinking of the echo cloud in PA at earlier epochs, hence the possibility that the echo cloud does not extend in front of the SN.  The inner echo lies at $\\theta\\sim1\\farcs15$ (all at the same radius, to within the errors) in 2001, over 0$<$PA$<$60.  In 1995, this same $z$ would correspond to $\\theta \\approx 0\\farcs55$, which is bordering the confusion region of the bright PSF of the SN.  The echo lies at a foreground distance $z = 81$~pc, and is perpendicular to the line of sight to within about 25$^\\circ$.  The geometry of all echoes is shown in Figure 2. Radial profiles through the echoes in F555W are plotted in Figure 3, demonstrating that these detections are above the background noise (e.g. $\\sigma_{\\rm F555W}=1.5$~DN~pix$^{-1}$).  Using the naming convention of \\citet{xu95}, we denote the outer echo as SW770 and the inner echo as NE260.  M81, with an angular momentum vector inclined 59$^\\circ$ along PA 62 \\citep{rot75} (such that the southwest side of the disk is closer to Earth), and SN~1993J, southwest of the nucleus of M81, would imply that the tilt of the NE260 dust sheet is roughly perpendicular to the disk of the galaxy. In comparison, the SW770 echo is inclined to within $\\sim30^\\circ$ of the disk plane.  The SN, with its massive progenitor contained within a dense gas cloud, is presumed to lie near the disk plane (shown in Fig.\\ 2).  We thus detect two dust structures, both extending more than one gas scale-height (c.f.\\ \\citealt{bro98}) above the plane.  SW770 sits a roughly constant 110 pc above the disk plane in M81, while NE260 appears to miss this plane by $\\ga$40~pc and extends at least 60 pc above the disk.  Without requiring the SN to lie in the disk plane, one might hypothesize that the disk plane passes near both echoing structures, implying the SN may lie $\\sim70-90$~pc behind the disk.  SW770 has measurable surface brightnesses over the PA range 160--280. Over 190$<$PA$<$250, all three bands (F450W, F555W and F814W) track each other in surface brightness consistently.  For $190<$PA$<250$, $\\langle \\mu \\rangle = 23.3\\pm0.1$, $23.3\\pm0.1$ and $24.6\\pm0.2$ in STMAG, respectively transforming to vegamag colors\\footnote{ Color transformations from STMAG to Johnson-Cousins (vegamags) were determined using {\\em synphot} and the Bruzual-Persson-Gunn-Stryker Spectrophotometry Atlas.} of roughly $B-V=0.6\\pm0.2$ and $V-I_C=0.0 \\pm 0.3$.  The spatial variation in surface brightness is similar in $B$ and $V$, both rising with inreasing PA values.  At PA$\\approx270$, $\\mu_{450}=23.0\\pm 0.1$ and $\\mu_{555}=23.1 \\pm 0.1$, while at PA$\\approx$180, $\\mu_{450}=23.6 \\pm 0.2$ and $\\mu_{555}=24.1 \\pm 0.2$.  In contrast $\\mu_{814}$ is nearly equally faint at both extremal PAs, $\\mu_{814}=25.1 \\pm 0.4$, implying colors at PA$\\sim270$ of $B-V=0.5 \\pm 0.3$ and $V-I_C=-0.9 \\pm 0.5$. In comparison, NE260 has approximately the same global colors as SW770 (over 190$<$PA$<$250): $B-V=0.5 \\pm 0.4$ and $V-I_C=0.2 \\pm 0.4$, and there is little evidence for such a color gradient.  The surface brightnesses themselves are fainter by about 0.5 mag arcsec$^{-2}$ compared to SW770.  In order to interpret these surface colors in terms of reflectivity, we must know the colors of the incident echoing flux.  Integrating over the entire SN lightcurve \\citep{ben94,ric94} from $\\sim3$ to 127 days after core collapse yields $B-V=0.73$, $V-I_C=0.69$ for the fluence of the (nearly) entire event. The color change due to dust reflectivity for NE260 and SW770 (with PA$<$250) is $\\Delta (B-V)=-0.1$ and $\\Delta (V-I_C)=-0.7$, while for SW770 with PA$>$250, $\\Delta (B-V)=-0.2$ and $\\Delta (V-I_C)=-1.6$. The largest changes in color are imparted in the Rayleighan scattering regime by very small particles. Integration of the scattering efficiency $S(\\lambda,a)$ \\citep{xu94} using the dust scattering parameters of  \\citet[and references therein]{wd01}, we find $S \\propto \\lambda^{-4.3}$ for $a<0.01 \\mu m$, yielding $\\Delta (B-V)_{\\rm max}=-0.96$ and $\\Delta (V-I_C)_{\\rm max}=-1.72$. As the observed color shifts are smaller, % the echoes should be consistent with a galactic dust distribution. Dust modeling will be examined in detail by \\citet{Sug03}.  Attributing the blue color of SW770 at PA$>$250 to a small-grain-only hypothesis is sufficiently improbable as to warrant alternative explanations.  Any extinction mechanism for F814W light should more-strongly affect the F450W and F555W bands.  One might invoke additional flux from mechanisms beyond direct reflection to increase the flux in F450W and F555W over F814W for PA$<$250.  Extended red emission \\citep{wit90a} would contribute significantly to the F814W band but not the bluer bands.  However, we caution that this very blue $V-I_C$ SW770 color is near the detection threshold, and we cannot rule out a ``hot pixel'' or variable star producing an erroneous image-subtraction residual in the F814W image.  For the sake of this discussion, we adopt a dust model with isotropically-scattering grains. We note that this assumption disagrees with some measurements of Galactic interstellar dust \\citep{wit90b,mat79,tol81}.  Using this model, we calculate the ratio of dust densities in the two echoes from the surface brightnesses and echo geometry. For an echo cloud that is thick relative to the the depth of dust echoing at a given time ($t_{s} \\bigl | {{dz} \\over {dt}} \\bigr |$), the surface brightness is predicted by \\cite{che86}: $$\\mu (\\theta) = {{n_{d} Q_{s} \\sigma_{d} } \\over {4 \\pi D^{2}}} ~ {{L_{\\nu} t_{s}} \\over {4 \\pi R^{2}}} ~ \\left| {{dz} \\over {dt}} \\right| ~ F(\\alpha),$$  \\noindent where the average apparent luminosity $L_{\\nu}/4 \\pi D^{2}$ over the SN light pulse duration $t_{s}$ is observed directly from the SN.  The geometric factors $R = \\sqrt{(z^2 + r^2)}$ (the SN-to-cloud distance), $\\alpha$ (the scattering angle), and $|dz/dt| = r^2/2ct^2 + c/2$ (describing the depth of the echoing region), are determined precisely by $\\theta$ from the light-travel-delay equation of an echo. This leaves: the grain scattering efficiency $Q_s$ and geometric cross-section $\\sigma_d$ (assumed equal in the two echoing clouds); the dust number density $n_{d}$; and the scattering phase function $F(\\alpha)= (1-g^2) / (1+g^2-2g \\cos{\\alpha})^{3/2}$ [Henyey \\& Greenstein 1941 -- which does quite well for small $\\alpha$ such as here \\citep{wit89}], where $g=\\overline {\\cos{\\alpha}}$ is the degree of forward scattering, here assumed to be zero. The ratio of the geometric factors $F(\\alpha)|\\dot{z}|/R^2$ for the two echoes is 1.52 (inner/outer), but SW770 is spread over $r_{outer}/r_{inner} = 1.71$ times the area.  Since the echoes are underresolved by WFPC2, SW770 should be about 1.13 times higher observed surface brightness than NE260, if they have the same dust properties and density.  In truth, we measure SW770 to be about 1.6 times brighter, implying either a 40\\% higher dust density in SW770, or the SW770 cloud is 40\\% thicker along the line of sight.  If, instead, the clouds are geometrically thin compared to $t_{s} \\left|\\frac{dz}{dt} \\right|$, geometric factors would predict an inner echo brightness 2.56 times that of SW770 (for equivalent dust), implying that the SW770 cloud contains 4.1 times higher dust surface density. How do these numbers change if we instead invoke a non-zero $g$? For reasonable values ($g \\approx 0.5$), the above brightness ratios change by only 1\\% or less, since the scattering angles for the two echoes are very similar ($\\alpha = 10\\fdg5$ for NE260 versus $8\\fdg5$ for SW770.)  The echo's effective width is given by $w = t_s ~ dr/dt$ where $dr/dt = c(z+ct/2)/\\sqrt{(2z+ct/2)ct}$, and $t_s$ is the total fluence in $V$ ($6.31 \\times 10^{-7}$ ergs cm$^{-2}$ \\AA$^{-1}$) divided by the maximum light flux, yielding $t_s = 3.5 \\times 10^6$s and $w = 0.23$~pc. A thin sheet of isotropic reflectors at the position of SW770 in year 2001 (with 170$<$PA$<$290) diverts no more than 0.0026\\% of the SN flux seen at Earth. We observe a flux in F555W of $\\sim4.3\\times 10^{-18}$ ergs cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$ from the echo, versus a corresponding maximum flux from the SN of $1.8\\times 10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$, or about 0.0024\\%.  SW770, over the position observed, appears to be optically thick.  This implies NE260 also has $A_v \\ga \\frac{1}{4}$.  For a dust albedo of 0.5 and grain diameter of 0.1~$\\mu$m (the largest Rayleigh-like particle), unit optical depth corresponds to $\\sim 15 ~\\mu$g~cm$^{-2}$ for grains of density 1~g~cm$^{-3}$.  If the gas-to-dust ratio is 100, this corresponds to $N_H \\approx 8\\times 10^{20}$~cm$^{-2}$ for SW770. At the position of the SN, $N_{\\mbox{\\footnotesize{H{\\sc i}}}} \\approx 10^{21}$~cm$^{-2}$, so SW770 structure is consistent with the dominant locus of gas along the Earth-SN sightline.  The velocity structure of this region of the galaxy has been studied in \\ion{H}{1} 21~cm emission and optical/UV absorption (of the SN itself).  This structure is relatively smooth and locally centered near $v_{lsr} = -135$ km s$^{-1}$ \\citep{rot75}, with some gas over the range $-155 < v_{lsr} < -115$ km s$^{-1}$ \\citep{ros75}.  In absorption against SN~1993J, the predominant M81 interstellar components are $-$119 and $-$135~km~s$^{-1}$, with possible lesser components at $-110$ and $-100$~km~s$^{-1}$ \\citep{vla94}. The former two interstellar components are each at least twice as strong in \\ion{Ca}{2} column density as the latter two (hence containing about 56\\% and 26\\% of the interstellar gas), and appear to be cold.  In IUE spectra of UV absorption lines \\citep{mar00}, a strong component at $-$130~km~s$^{-1}$ and a weaker one at $-$90~km~s$^{-1}$ is seen in low-ionization species, probably consistent with the \\ion{Ca}{2} components.  It is possible that the inner and outer echoes correspond to the two dominant absorption features ($-$119 and $-$135~km~s$^{-1}$), but this is difficult to state with certainty given the limited amount of data and the partial covering factor of the structures involved.  The structure and strength of the echoes seems to imply, however, that major portions of interstellar material in this part of the disk may be broken into fragments and perhaps even propelled a scale height or more out of the disk plane."
    },
    "0207/astro-ph0207204_arXiv.txt": {
        "abstract": "Laser guide star systems based on Rayleigh scattering require some means to deal with the flash of low altitude laser light that follows immediately after each laser pulse.  These systems also need a fast shutter to isolate the high altitude portion of the focused laser beam to make it appear star-like to the wavefront sensor.  We describe how these tasks are accomplished with UnISIS, the Rayleigh laser guided adaptive optics system at the Mt. Wilson Observatory 2.5-m telescope.  We use several methods:  a 10,000 RPM rotating disk, dichroics, a fast sweep and clear mode of the CCD readout electronics on a 10 $\\mu$s timescale, and a Pockel's cell shutter system.  The Pockel's cell shutter would be conventional in design if the laser light were naturally polarized, but the UnISIS 351 nm laser is unpolarized.  So we have designed and put into operation a dual Pockel's cell shutter in a unique bow tie arrangement. ",
        "introduction": "The UnISIS Rayleigh laser guide star system has been commissioned as described in \\citet{tho02}.  The system is built around a 30 Watt excimer laser that emits short 90 mJ pulses of 351 nm light at rates up to 333 Hz.  These pulses are projected in the \"full- aperture broadcast\" mode and are focused in the stratosphere at an altitude of $\\sim$20 km above mean sea level ( $\\sim$18 km above the telescope).  Immediately after the outgoing pulse of laser light hits the telescope primary mirror, a bright flash of low altitude Rayleigh scattered light fills the near-field of the telescope.  As described below, the science cameras can be shielded from this near-field burst of light by dichroics because the projection method -- a 10,000 RPM rotating glass disk -- blocks the adaptive optics system from seeing a large fraction of this light in the first few microseconds after the laser pulse.  It is more difficult to protect the laser guide star wavefront camera because it is highly sensitive to the laser wavelength.  We use two methods to accomplish the latter task:  a continuous high speed read and flush of the CCD wavefront camera and a Pockel's cell shutter.  We start in Sec. 2 with a general discussion of ways that other laser guide star projection systems handle the contaminating Rayleigh scattered light, and then we review the UnISIS rotating disk projection system.  In Sec. 3 we describe briefly the dichroic isolation of the UnISIS science cameras.  In Sec. 4 we describe how the UnISIS wavefront sensor is hidden from the low altitude near-field burst of Rayleigh light with our new bow tie Pockel's cell shutter system, and then in Sec. 5 we describe the second level of protection for the wavefront sensor:  a continuous high-speed read and flush mode of the UnISIS wavefront CCD.  The general characteristics of the UnISIS CCD system and its future upgrades are also described in Sec. 5. An up to date block diagram of the UnISIS optical layout will be included in the next paper of this series, but a close approximation (not showing the bowtie shutter) can be found in \\citet{tho98}. ",
        "conclusions": "Laser guided adaptive optics is still in an early phase of development, and there are many ways yet to be found to configure both the lasers and the cameras to minimize interference from the inevitable low-altitude Rayleigh scattered light, whether this comes from sodium resonance laser guide stars or from Rayleigh laser guide stars.  We have presented here several examples of how this is accomplished in UnISIS.  We find that the greatest design flexibility comes with pulsed lasers rather than CW systems, and one of the more significant characteristics of the system design is whether aircraft and satellite avoidance can be handled in a fashion that can be called \"Stealth\".  Pulsed lasers at 351 nm provide a distinct advantage in both regards.  However, new technological innovations will continue to alter the balance for some time to come.  One simple example is the potential ability of embedding the shutter function in the silicon substrate of the wavefront sensor CCD, thereby making Pockel's cell switching unnecessary.  MIT / Lincoln Laboratory has the ability to produce sensors of this type, but no UV sensitive versions are currently available for wavefront sensing \\citep{rei02}.  Other technological advances will, no doubt, come along as laser guided adaptive optics systems mature.   A number of people have contributed to the work reported here.  This includes Richard Castle, Dr. E. Harvey Richardson and Bill Knight and his machine shop crew.  At Mt. Wilson Observatory we acknowledge support provided by Mount Wilson Institute Director Dr. Robert Jastrow and technical assistance by Robert Cadman, Sean Hoss, Chris Hodge, Joe Russell, Victor Castillo and Thomas Schneider of Schneider Engineering.  Telescope operator support at the 2.5-m telescope was provided by Kirk Palmer, Michael Bradford, and Jim Strogen.  Marconi CCD39 wavefront software support was provided by Jamie Erickson and Scott Striet.  The Pockel's cells used in UnISIS were supplied by ThermoTrex Incorporated (San Diego, CA) as part of a surplus equipment transfer, and we thank Dr. David Sandler for assistance in that transfer.  This work was supported by grants from the National Science Foundation:  AST-9220504, AST-0096741 and by funds from both the University of Illinois and the New Mexico Institute of Mining and Technology.  All support is very gratefully acknowledged."
    },
    "0207/astro-ph0207518_arXiv.txt": {
        "abstract": "We review some of the characteristics of irradiated extrasolar giant planets (EGPs), in anticipation of their direct detection from the ground and from space.  Spectral measurements are the key to unlocking their structural and atmospheric characteristics and to determining the true differences between giant planets and brown dwarfs. In this spirit, the theoretical spectral and atmospheric calculations we summarize here are in support of the many searches for EGPs to be conducted in the coming decade by astronomers from around the world. ",
        "introduction": "\\label{intro}  Since the discovery of 51 Pegasi b (Mayor and Queloz 1995) and the nearly one hundred extrasolar giant planets (EGPs) that have been detected subsequently by radial velocity techniques (see this proceedings and references in Burrows et al. 2001), an increasing fraction of the world's astronomers has been engaged in determining the best means to detect such planets directly.  While the orbital elements of substellar-mass objects with M$_{\\rm p}$sin(i)s that range from $\\sim$0.2 to $\\sim$10 \\mj can constrain formation mechanisms and dynamical evolution, they are no substitute for direct spectral measurements. It is by imaging the planet and obtaining optical, near-infrared, and mid-infrared spectra that an EGP's atmospheric structure, composition, gravity, radius, and mass can be determined.  Such physical characteristics are essential data if the study of EGPs is to mature in the next decade into a major astronomical field.  They are also essential if the distinctions between brown dwarfs and giant planets of the same mass are to be determined.  It may be that, for a given primary star, different origins and histories at birth translate into different compositions and rotation rates.  Spectra will be essential in determining this.  From space, SIM will provide accurate astrometric masses (not merely M$_{\\rm p}$sin(i)s) for the known EGPs.  However, from the ground spectral deconvolution techniques, adaptive optics, interferometry (e.g., using the LBT, VLTI, KIA) and a host of promising and novel methods summarized during this conference encourage one to believe that light from an EGP will soon be detected.  From space, optical coronagraphs with ultrasmooth mirrors (e.g., Eclipse, ESPI, JPF), infrared interferometers, and precision transit missions (e.g., MOST, Eddington, Kepler, COROT) are in various stages of planning or preparation.   The space-based transit missions will be preceded by a host of ground-based transit searches (e.g., STARE, BEST, WASP, STEPSS, TeMPEST). Transit data married with precise stellar and Doppler wobble measurements can provide mass-radius relations for the close-in EGPs (``roasters\") from which one can extract structural and evolutionary information (Guillot et al. 1996; Burrows et al. 2000). The first discovered transiting extrasolar planet, HD209458b, for which a large radius of $\\sim$1.4 \\rj and a mass of 0.69 \\mj were obtained (Cody and Sasselov 2002; Brown et al. 2001; Charbonneau et al. 2000; Henry et al. 2000) has jump-started the scramble to understand EGPs under severe irradiation regimes.  The demonstration that the depth of HD209458b's transit is wavelength-dependent (Hubbard et al. 2001) and that neutral sodium resides in its atmosphere (Charbonneau et al. 2002) is an indication of the vast potential of transit studies. The discovery of a collection of transiting planets, not just one, will be a milestone in the study of EGPs.  This tempo of activity focussed on obtaining direct measurements of EGP properties demands a corresponding effort by the theoretical community to calculate the spectra of EGPs around a variety of stars, at a variety of orbital distances, and with a variety of masses and ages.  We have undertaken such a project and in this short contribution present some of our preliminary results.  A more comprehensive treatment can be found in Sudarsky, Burrows, and Hubeny (2002, in preparation). We have calculated EGP spectra for 51 Peg b, $\\tau$ Boo b, HD209458 b, $\\upsilon$ And b,c,d, GJ 876 b,c, $\\epsilon$ Eridani b, 55 Cnc b,c,d, HD114762 b, HD1237 b, and a host of other radial-velocity EGPs, as well as for theoretical objects at the full potential range of orbital distances, around a collection of stellar subtypes, and employing a variety of cloud models.  The fluxes at the Earth, as well as the planet/star contrasts as a function of wavelength from 0.4 \\mic to 30 \\mic have been determined. ",
        "conclusions": "\\label{sum}  NASA and ESA are poised to spend 100's of millions of dollars in the next decade to detect and characterize extrasolar planets.  The best data in the short term will be for EGPs, not terrestrial planets.  Hence, for those most interested in discovering Earths, the natural and inevitable path is by way of the giant planets now being discovered in profusion in the solar neighborhood. For those for whom EGPs are not mere stepping stones, the next decade promises a rich harvest of new worlds and stimulating finds.  Our calculations are designed to provide the necessary theoretical underpinnings for this quest at one of astronomy's newest frontiers."
    },
    "0207/astro-ph0207032_arXiv.txt": {
        "abstract": "I present an overview of strong and weak gravitational lensing by galaxy clusters. After briefly introducing the principles of gravitational lensing, I discuss the main lessons learned from lensing on the mass distribution in clusters and their relation to cosmology. ",
        "introduction": "Gravitational lensing phenomena due to galaxy clusters can naturally be split into two categories, strong and weak. Strong lensing was detected in 1986, when highly elongated, curved, long features of low surface brightness were found in two clusters, Abell~370 and Cl~2244 (Lynds \\& Petrosian 1986; Soucail et al.~1987a,b). Of the three possible explanations suggested, spectroscopy selected gravitational lensing when it turned out that these ``giant arcs'' had substantially higher redshifts than the clusters (Soucail et al.~1988). Weak lensing gives rise to much less spectacular distortions of background-galaxy images, termed ``arclets'' (Fort et al.~1988; Tyson et al.~1990). Since galaxies are not intrinsically symmetric, such distortions can only be quantified statistically by averaging over many images, commonly adopting the assumption that galaxy ellipticities average to zero in absence of lensing.  While arcs require compact, dense cluster cores and thus probe their central mass distribution, arclets can be found everywhere across clusters and allow their mass distribution to be mapped even at clustercentric distances comparable to the virial radius. Unlike other methods for quantifying the mass distribution in clusters, lensing has the advantage that it is sensitive only to the surface mass density projected along the line-of-sight, irrespective of its composition or physical state.  I review here the main lessons that have been learned from both weak and strong lensing by clusters. I first summarise the physical assumptions and principles underlying interpretations of lensing phenomena, keeping the formalism to the necessary minimum. I then turn to strong lensing and explain the key results and a number of open problems. After explaining the principle of weak-lensing techniques, I describe results obtained from weak lensing in clusters and conclude with an outlook and a summary. ",
        "conclusions": "The main lessons learned from lensing by clusters can be summarised as follows:  \\begin{itemize}  \\item Clusters are dominated by dark matter. They are asymmetric, substructured, and highly concentrated, and their mass-to-light ratios in the optical are of order $M/L\\sim200-400\\,h$ in solar units.  \\item Masses determined from X-ray data and gravitational lensing tend to disagree in unrelaxed clusters, for which lensing masses are biased high, and X-ray masses biased low. In relaxed clusters, the different mass estimates tend to agree well.  \\item The statistics of giant arcs constrain cosmology. Detailed numerical simulations indicate that there is a disagreement with other cosmological experiments in that arc statistics prefers an open, low-density model \\emph{without} cosmological constant. Arc statistics can also place an upper bound on interaction cross sections for dark-matter particles.  \\item Weak gravitational lensing allows the projected distribution of the total cluster matter to be mapped. Cluster density profiles are compatible with numerical simulations. There are considerable fluctuations in cluster mass-to-light ratios.  \\item Clusters, more generally dark-matter haloes, can be detected through weak-lensing techniques. Several independent observations suggest that very faint and possibly dark clusters exist. Massive and compact clusters exist at redshifts as high as $\\sim0.8$.  \\item Joint analyses of different types of cluster data (e.g.~from gravitational lensing, X-ray, and thermal Sunyaev-Zel'dovich observations) allow clusters to be deprojected along the line-of-side.  \\end{itemize}  Perhaps the most exciting issues in lensing-related cluster studies are the detection of dark-matter haloes irrespective of the radiation they emit or absorb, detailed cluster analyses jointly using different types of data, possible constraints on the nature of dark matter, and the relation between the statistics of giant arcs, cluster formation, and cosmology.  Due to lack of time, I was unable to touch on several additional exciting aspects of cluster lensing. To mention just one, clusters have been used highly successfully as gravitational telescopes for studying populations of faint sources at high redshift. For instance, the gravitational magnification by clusters was used to study high-redshift galaxies in the sub-millimetre regime (Smail et al.~1997) or spectroscopically in the optical (Pell\\'o et al.~1998). While such work does not primarily target clusters, it shows how gravitational lensing by clusters can be used as a powerful astronomical tool."
    },
    "0207/astro-ph0207562_arXiv.txt": {
        "abstract": "We report the results of a $FUSE$ study of high velocity \\ion{O}{6} absorption along complete sight lines through the Galactic halo in directions toward 100 extragalactic objects and 2 halo stars. The high velocity \\ion{O}{6} traces a variety of phenomena, including tidal interactions with the Magellanic Clouds, accretion of gas, outflowing material from the Galactic disk, warm/hot gas interactions in a highly extended Galactic corona, and intergalactic gas in the Local Group. We identify 85 high velocity \\ion{O}{6} features at $\\ge3\\sigma$ confidence at velocities of $-500 < v_{LSR} <  +500$ \\kms. There are an additional 6 confirmed or very likely ($>90$\\% confidence) features plus 2 tentative detections between $v_{LSR} = +500$ and +1200 \\kms; these very high velocity \\ion{O}{6} features trace intergalactic gas beyond the Local Group. The 85 \\ion{O}{6} features have velocity centroids ranging from $-372 \\lesssim \\bar{v}_{LSR} \\lesssim -90$ \\kms\\ to $+93 \\lesssim \\bar{v}_{LSR} \\lesssim +385$ \\kms, line widths b~ $\\sim 16-81$ \\kms\\ with an average of $\\langle$b$\\rangle$ = $40\\pm14$ \\kms, and an average \\ion{O}{6} column density $\\langle \\log N \\rangle = 13.95\\pm0.34$ with  a median value of 13.97. Values of b greater than the $17.6$ \\kms\\ thermal width expected for \\ion{O}{6} at $T \\sim 3\\times10^5$\\,K indicate that additional non-thermal broadening mechanisms are common. The \\ion{O}{6} $\\lambda1031.926$ absorption is detected at $\\ge 3\\sigma$ confidence along 59 of the 102 sight lines surveyed. The high velocity \\ion{O}{6} detections indicate that $\\sim60$\\% of the sky (and perhaps as much as $\\sim85$\\%, depending on data quality considerations)  is covered by high velocity H$^+$ associated with the \\ion{O}{6}. $N({\\rm H}^+)\\gtrsim4\\times10^{16}$ cm$^{-2}$ if the high velocity hot gas has a metallicity similar to that of the Magellanic Stream.  About 30\\% of the sky is covered by the hot, high velocity H$^+$ at a level of $N({\\rm H}^+)\\gtrsim4\\times10^{17}$ cm$^{-2}$, which is similar to the detection rate found for \\ion{H}{1} 21\\,cm emission produced by warm neutral gas at a comparable column density level. Some of the high velocity \\ion{O}{6} is associated with known \\ion{H}{1} structures (the Magellanic Stream, Complex~A, Complex~C, the Outer Spiral Arm, and several discrete \\ion{H}{1} HVCs). Some of the high velocity \\ion{O}{6} features have no counterpart in \\ion{H}{1} 21\\,cm emission, including  discrete absorption features and positive velocity absorption wings extending from $\\sim100$ to $\\sim300$ \\kms\\ that blend with lower velocity absorption produced by the Galactic thick disk/halo. The discrete features may typify clouds located in the Local Group, while the \\ion{O}{6} absorption wings may be tidal debris or material expelled from the Galactic disk. Most of the \\ion{O}{6} features have velocities incompatible with those of the Galactic halo, even if the halo has decoupled from the underlying Galactic disk. The reduction in the dispersion about the mean of the high velocity \\ion{O}{6} centroids when the velocities are converted from the LSR to the GSR and LGSR reference frames is necessary (but not conclusive) evidence that some of the clouds are located outside the Galaxy. Most of the \\ion{O}{6} cannot be produced by photoionization, even if the gas is irradiated by extragalactic ultraviolet background radiation. Several observational quantities indicate that collisions in hot gas are the primary ionization mechanism responsible for the production of the \\ion{O}{6}.  These include the ratios of \\ion{O}{6} column densities to those of other highly ionized species (\\ion{C}{4}, \\ion{N}{5}) and the strong correlation between $N$(\\ion{O}{6}) and \\ion{O}{6} line width. Consideration of the possible sources of collisional ionization favors production of some of the \\ion{O}{6} at the boundaries between cool/warm clouds of gas and a highly extended ($R \\gtrsim 70$ kpc), hot ($T > 10^6$\\,K), low-density ($n \\lesssim 10^{-4}-10^{-5}$ cm$^{-3}$) Galactic corona or Local Group medium. The existence of a hot, highly extended Galactic corona or Local Group medium and the prevalence of high velocity \\ion{O}{6} are consistent with predictions of current galaxy formation scenarios. Distinguishing between the various phenomena producing high velocity \\ion{O}{6} in and near the Galaxy will require continuing studies of the distances, kinematics,  elemental abundances,  and physical states of the different types of high velocity \\ion{O}{6} found in this study. Descriptions of galaxy evolution will need to account for the highly ionized gas, and future X-ray studies of hot gas in the Local Group  will need to consider  carefully the relationship of the X-ray absorption/emission to the complex high velocity absorption observed in \\ion{O}{6}. ",
        "introduction": "Understanding galaxy formation and evolution requires observational information about hot, highly ionized gas in and near galaxies.   Numerical simulations of cosmological structure formation in the presence of cold dark matter indicate that a significant fraction of the baryonic material at low redshift should be shock-heated to temperatures of $10^5-10^7$\\,K as the gas collapses and forms clusters and groups of galaxies (e.g., Cen \\& Ostriker 1999; Dav\\'e et al. 1999, 2001).  This reservoir of hot gas is detectable through ultraviolet absorption-line measurements (Tripp, Savage, \\& Jenkins 2000; Savage et al. 2002a), but a complete picture of how the hot gas in galaxies and the intergalactic medium (IGM) are related does not yet exist because a variety of internal and external processes affect the heating and distribution of the interstellar gas in and around galaxies.  In addition to the galaxy formation process, the accretion of satellite galaxies, tidal interactions, star formation, galactic winds, and galaxy-IGM interactions may all contribute to the production of hot gas.  To some extent, all of these activities operate in and around our own galaxy, so studying the hot gas in the immediate environment of the Milky Way is a logical step in assessing the relevance and roles of these processes locally.   One of the primary science objectives of the {\\it Far Ultraviolet Spectroscopic Explorer} ($FUSE$) mission is to determine the  properties of hot, highly ionized gas in the low-redshift universe.  A key component of this research has been the study of \\ion{O}{6} absorption along many sight lines through the Galactic halo.  Such observations were not possible until $FUSE$ was launched because previous observatories lacked either the spectral resolution or sensitivity to study the velocity structure of the \\ion{O}{6} absorption toward distant background sources. Among the most interesting results to date is the detection of \\ion{O}{6} in high velocity clouds (HVCs) (Sembach et al. 2000; Murphy et al. 2000). For decades astronomers have studied the neutral (\\ion{H}{1}) content of the clouds and have debated the origin and location of the high velocity gas.  While it is generally accepted that no single model can account for all of the observed properties of the \\ion{H}{1} HVCs (see Wakker \\& van~Woerden 1997 for a review), there is no unanimity on the fundamental properties of the different types of HVCs. With the re-introduction of the idea that some of the \\ion{H}{1} HVCs could be located outside of the Milky Way if they are embedded in halos of dark matter (Blitz et al. 1999), the possible range of locations and cloud properties has widened. Determining the ionization and hot gas content of the high velocity gas bears directly on the locations of the clouds and their interactions with other components of the gaseous interstellar and intergalactic media.  The \\ion{O}{6} $\\lambda\\lambda1031.926, 1037.617$ doublet lines in the far-ultraviolet spectral region are the best  lines to use for kinematical investigations of hot ($T \\gtrsim 10^5-10^6$\\,K) gas in the low-redshift universe.  Oxygen is the most abundant element heavier than helium, and the \\ion{O}{6} lines have large oscillator strengths ($f_{1032} = 0.133$, $f_{1038} = 0.0661$; Morton 1991).  X-ray spectroscopy of higher ionization lines (e.g., \\ion{O}{7}, \\ion{O}{8}) is possible with {\\it XMM-Newton} and the {\\it Chandra X-ray Observatory} for a small number of sources, but the spectral resolution (R~$\\equiv \\lambda/\\Delta\\lambda \\lesssim 1000$) is modest compared to that afforded by $FUSE$ (R~$\\sim 15,000$). Lower ionization lines observable at high spectral resolution at ultraviolet wavelengths are generally either much weaker than the \\ion{O}{6} lines (e.g., \\ion{N}{5}~$\\lambda\\lambda1238.821, 1242.804$) or are better tracers of collisionally ionized gas at temperatures $T \\lesssim 10^5$\\,K (e.g., \\ion{C}{4} $\\lambda\\lambda1548.195, 1550.770$, \\ion{C}{3} $\\lambda977.020$, \\ion{Si}{4} $\\lambda\\lambda1393.755, 1402.770$, \\ion{Si}{3} $\\lambda1206.500$). This latter set of ions is also considerably more susceptible to photoionization than \\ion{O}{6} since the photoionization cross sections of the ions are large and the ionization potentials are less than 54\\,eV, the energy of the \\ion{He}{2} ionization edge.   \\ion{O}{6} can be produced by photoionization under special conditions involving a hard radiation field and a very low gas density, but as we will show later, this does not appear to be a viable production mechanism for {\\it most} of the \\ion{O}{6} observed in and near the Milky Way.  This article is one in a series of three papers devoted to $FUSE$ observations of \\ion{O}{6} absorption along complete paths through the Galactic halo in the directions of quasars, active galactic nuclei (AGNs), and BL Lac objects. The other two papers include a catalog with basic measurements and illustrations of all of the \\ion{O}{6} profiles obtained in the survey (Wakker et al. 2002) and a companion study of the \\ion{O}{6} absorption associated with the thick disk of the Milky Way (Savage et al. 2002b).  Here, we concentrate on the properties of the \\ion{O}{6} absorption at high velocities with respect to those expected for gas participating in differential Galactic rotation (generally, ``high velocity'' refers to $100 < |v_{LSR}| < 400$ \\kms\\ in this paper). High velocity \\ion{O}{6} absorption associated with the Magellanic Clouds is discussed elsewhere (Friedman et al. 2000; Danforth et al. 2002; Hoopes et al. 2001, 2002; Howk et al. 2002b).  This paper is organized as follows.  In \\S2 we describe the observations and data reduction.  We present the \\ion{O}{6} HVC measurements in \\S3 and describe the general types of high velocity gas seen in \\S4.  Section 5 contains information about the column densities, sky covering fractions, velocities, and line widths of the \\ion{O}{6} absorption.  In \\S6 we discuss the high velocity \\ion{O}{6} features associated with previously identified high velocity clouds, while \\S7 specifies the few sight lines that show high velocity \\ion{H}{1} 21\\,cm emission with no corresponding \\ion{O}{6} absorption.  Section~8 highlights new high velocity gas detected only in \\ion{O}{6} absorption.  In \\S9 we consider how the high velocity gas is ionized. Section~10 describes the general kinematical behavior of the high velocity gas. Section 11 contains a discussion of the high velocity \\ion{O}{6} features and the implications of this work for understanding hot gas in the low-redshift universe.  We conclude  with a summary of results in \\S12. ",
        "conclusions": "In this paper we report the results of an extensive  $FUSE$ study of high velocity \\ion{O}{6} absorption along complete sight lines through the Galactic halo in the directions toward 100 AGNs/QSOs and two distant halo stars.  Companion studies describe the observations and data (Wakker et al. 2002) and the results for \\ion{O}{6} absorption produced by the thick disk and halo of the Milky Way (Savage et al. 2002b). In this study, the cutoff between high velocity gas and Galactic thick disk/halo absorption was generally near $|v_{LSR}| \\sim 100$ \\kms.  Our main conclusions regarding the high velocity gas are as follows:   \\noindent 1) We identify 85 individual high velocity \\ion{O}{6} features along the 102 sight lines in our sample. A critical part of this identification process involved detailed consideration of the absorption produced by \\ion{O}{6} and other species (primarily H$_2$) in the thick disk and halo of the Galaxy, as well as the absorption produced by low-redshift intergalactic absorption lines of \\ion{H}{1} and ionized metal species.  \\noindent 2)  We searched for absorption in a velocity range of $-1200 < v_{LSR} < +1200$ \\kms\\ around the \\ion{O}{6} $\\lambda1031.926$ line. With few exceptions, the high velocity \\ion{O}{6} absorption is confined to $|v_{LSR}| < 400$ \\kms, indicating that the \\ion{O}{6} features observed are either associated with the Milky Way or nearby clouds within the Local Group. The 85 high velocity \\ion{O}{6} features have velocity centroids ranging from $-372 \\lesssim \\bar{v}_{LSR} \\lesssim -90$ \\kms\\ to $+93 \\lesssim \\bar{v}_{LSR} \\lesssim +385$ \\kms. There are an additional 6 confirmed or very likely ($>90$\\% confidence) detections and 2 tentative detections of \\ion{O}{6} between $v_{LSR} = +500$ and +1200 \\kms; these very high velocity features probably trace intergalactic gas beyond the Local Group.  \\noindent 3) The 85 high velocity \\ion{O}{6} features have logarithmic column densities of 13.06 to 14.59, with an average of $\\langle \\log N \\rangle = 13.95\\pm0.34$ and a median of 13.97. The average \\ion{O}{6} column density is a factor of 2.7 times lower than the typical column density for a sight line through the \\ion{O}{6} layer in the thick disk/halo of the Galaxy.  \\noindent 4) The line widths of the 85 high velocity \\ion{O}{6} features range from $\\sim 16$ \\kms\\ to 81\\kms, with an average of $\\langle {\\rm b} \\rangle = 40\\pm14$ \\kms. The lowest values of b are close to the thermal width expected for \\ion{O}{6} at $T \\sim 3\\times10^5$\\,K, while the higher values of b indicate that additional non-thermal broadening mechanisms are common.  \\noindent 5) We detect high velocity \\ion{O}{6} $\\lambda1031.926$ absorption with integrated (total) values of $W_\\lambda \\gtrsim 30$ m\\AA\\ at $\\ge 3\\sigma$ confidence along 59 of the 102 sight lines surveyed.  For the highest quality sub-sample of the dataset, the high velocity detection frequency increases to 22 of 26 sight lines.  Forty of the 59 sight lines have high velocity \\ion{O}{6} $\\lambda1031.926$ absorption with  $W_\\lambda > 100$ m\\AA, and 27 have $W_\\lambda > 150$ m\\AA. Converting these \\ion{O}{6} detections into estimates of $N$(H$^+$) in the hot gas indicates that $\\sim60$\\% of the sky (and perhaps as much as $\\sim85$\\%) is covered by hot ionized hydrogen at a level of $N({\\rm H}^+) \\gtrsim4\\times10^{16}$ cm$^{-2}$ if the high velocity gas has a metallicity similar to that of the Magellanic Stream ($Z\\sim0.2-0.3$), and $\\sim30$\\% of the sky is covered at a level of $N({\\rm H}^+)\\gtrsim4\\times10^{17}$ cm$^{-2}$.  The covering factor of the hot, high velocity H$^+$ associated with the  \\ion{O}{6} is similar to that found for high velocity \\ion{H}{1} 21\\,cm emission of warm neutral gas at a comparable column density level.  \\noindent 6) High velocity \\ion{O}{6} absorption is observed in almost all cases where high velocity \\ion{H}{1} 21\\,cm emission is observed.  The 6 sight lines for which this is not true may contain small-scale structure within the field of view of the 21\\,cm observations that is not traced by the \\ion{O}{6} absorption measures.  In a few cases, ionization effects may also explain the absence of \\ion{O}{6}.  Three of the six cases are good candidates for follow-up STIS observations to determine if low metallicity may be the reason for the absence of \\ion{O}{6} absorption.  \\noindent 7) Some of the high velocity \\ion{O}{6} is associated with well-known high velocity structures.  These include the Magellanic Stream, Complex~A, Complex~C, the Outer Arm, and several discrete \\ion{H}{1} HVCs.  \\noindent 8) Some of the high velocity \\ion{O}{6} features have no counterpart in \\ion{H}{1} 21\\,cm emission.  These include discrete high velocity features as well as broad positive velocity \\ion{O}{6} absorption wings that blend with lower velocity absorption features produced by the Galactic thick disk/halo. The discrete features may typify clouds located in the Local Group.   \\noindent 9) The broad, high velocity \\ion{O}{6} absorption wings are concentrated mainly in the northern Galactic hemisphere (18/22 sight lines) and may trace either tidal debris or thick disk/halo gas that has been accelerated to high velocities by star-formation activity in the Galactic disk.  If the latter interpretation is correct, the gas may be related to the northern hemisphere thick disk/halo \\ion{O}{6} enhancement observed by Savage et al. (2002b).  \\noindent 10) Most of the high velocity \\ion{O}{6} features have velocities incompatible with those of Galactic rotation (by definition).  The kinematics of the high velocity \\ion{O}{6} are not described adequately by models in which the Galactic halo decouples from the underlying disk.  There is also no obvious signature of a hot Galactic wind emanating from the Galactic center, but selection effects preclude a definitive statement about its existence.   \\noindent 11) The dispersion about the mean of the high velocity \\ion{O}{6} centroids  decreases when the velocities are converted from the LSR to the GSR and LGSR reference frames.   While this reduction is expected if the \\ion{O}{6} is associated with gas in a highly extended Galactic corona or in the Local Group,  it does not stand alone as sufficient proof of an extragalactic location.   The clear separation of  the various phenomena  producing high velocity \\ion{O}{6} in and near the Galaxy will require continuing studies of the distances, kinematics,  elemental abundances,  and physical states of the different types of \\ion{O}{6} HVCs.  \\noindent 12) We find it unlikely that many of the observed \\ion{O}{6} features are produced by photoionization, even if the gas is irradiated by extragalactic ultraviolet background radiation. Rather, several observational constraints indicate that collisional ionization in hot ($T\\sim10^5-10^6$\\,K) gas is likely the dominant ionization process for most of the \\ion{O}{6}.  The constraints include the amount of \\ion{O}{6} observed, the ratios of \\ion{O}{6} column densities to those of other highly ionized species (\\ion{C}{4}, \\ion{N}{5}), and the strong correlation between $N$(\\ion{O}{6}) and \\ion{O}{6} line width.  \\noindent 13) Consideration of the possible sources of collisional ionization favors production of some of the high velocity  \\ion{O}{6} at the turbulent boundaries between cool/warm clouds of gas and a highly-extended, hot ($T > 10^6$\\,K) Galactic corona or Local Group.   The corona must have a low density ($n \\lesssim 10^{-4}-10^{-5}$ cm$^{-3}$) and be very large ($R \\gtrsim70$ kpc) to explain the \\ion{O}{6} observed in the Magellanic Stream and other putative Local Group clouds. This corona is much more extensive than the Galactic thick disk/halo region considered in previous hot gas investigations.  \\noindent 14) The existence of a hot, highly extended Galactic corona or Local Group medium and the prevalence of high velocity \\ion{O}{6} are consistent with predictions that there should be a considerable amount of hot gas  left over from the formation of the Local Group. Descriptions of galaxy evolution will need to account for highly ionized gas of the type observed in this study, and future X-ray studies of hot gas in the Local Group will need to consider  carefully the relationship of the X-ray absorption/emission to the complex high velocity absorption observed in \\ion{O}{6}."
    },
    "0207/astro-ph0207081_arXiv.txt": {
        "abstract": "{ Millimetre-VLBI is an important tool in AGN astrophysics, but it is limited by short atmospheric coherence times and poor receiver and antenna performance. We demonstrate a new kind of phase referencing for the VLBA, enabling us to increase the sensitivity in mm-VLBI by an order of magnitude. If a source is observed in short cycles between the target frequency, $\\nu_{\\rm t}$, and a reference frequency, $\\nu_{\\rm ref}$, the $\\nu_{\\rm t}$ data can be calibrated using scaled-up phase solutions from self-calibration at $\\nu_{\\rm ref}$. We have demonstrated the phase transfer on 3C~279, where we were able to make an 86~GHz image with $90~\\%$ coherence compared to self-calibration at $\\nu_{\\rm t}$. We have detected M81, our science target in this project, at 86~GHz using the same technique. We describe scheduling strategy and data reduction. The main impacts of fast frequency switching are the ability to image some of the nearest, but relatively weak AGN cores with unprecedented high angular resolution and to phase-reference the $\\nu_{\\rm t}$ data to the $\\nu_{\\rm ref}$ core position, enabling the detection of possible core shifts in jets due to optical depth effects. This ability will yield important constraints on jet properties and might be able to discriminate between the two competing emission models of Blandford-K\\\"onigl jets and spherical advection-dominated accretion flows (ADAFs) in low-luminosity AGNs.}  \\authorrunning{Middelberg et al.} ",
        "introduction": "The regions where AGN jets are launched and collimated are difficult to observe with VLBI because the jets are launched very close to the black hole and most of the bright objects are very distant.  Only in the closest AGN in Sgr A*, M87, M84, Cen A and M81 can the highest resolution observations resolve several tens of Schwarzschild radii, comparable to the scale of 10-1000~$R_{\\rm s}$ where jets are predicted to be launched and collimated (eg Koide et al.  2000, Appl \\& Camenzind 1993). Whilst the collimation regions remain unresolved in Sgr A* and M81, there are hints for resolving it in M87 (Junor et al. 1999).  The highest resolution VLBI images are currently achieved from observations at 86~GHz. However, short atmospheric coherence times and poor receiver performance and antenna efficiencies severely limit the observations to sources with $S_{\\rm 86\\,GHz}\\sim0.4~{\\rm Jy}$. We have developed a new phase-referencing strategy for the VLBA that is not limited by self-calibration at the target frequency $\\nu_{\\rm t}$, but only at a lower reference frequency $\\nu_{\\rm ref}$. A source is observed switching between these two frequencies on a timescale that does not exceed half the atmospheric coherence time. After self-calibrating the $\\nu_{\\rm ref}$ data, the phase solutions are multiplied by the frequency ratio $r=\\nu_{\\rm t}/\\nu_{\\rm ref}$ and added to an instrumental, antenna-based phase offset $\\Delta\\phi$. Using these phases, the $\\nu_{\\rm t}$ data can be imaged.  In a pilot project, we have observed M81 and 3C~279 on January 5, 2002, with various frequency pairs and cycle times to develop the calibration technique. We describe the details and results of this project and what we have learnt for future observations. ",
        "conclusions": "We have developed a new phase-referencing strategy to make mm-VLBI observations of sources too weak for self-calibration. The strategy is based on rapid changes between the target frequency and a reference frequency and uses scaled-up reference frequency phase solutions to calibrate the target frequency. We have observed 3C~279 as a strong test source and M81 as a weak science target. Using this technique, we were able to image 3C~279 at 86~GHz with a $10~\\%$ coherence loss compared to conventional self-calibration, and, for the first time, we were able to make a marginal detection of M81 at 86~GHz.  We have learnt three basic lessons from our project: 1) To prevent loss of coherence occuring from phase wraps at $\\nu_{\\rm ref}$, one should select frequencies such that their ratio is an integer. Given the VLBA frequency agility, one can almost always tune the frequencies to satisfy this condition. 2) To increase efficiency, one should use asymmetric switching times. The symmetric times that we used gave us only 8~s, or $25~\\%$ of data per cycle of 30~s duration. For future projects, we will use cycles of 60~s with 16~s at $\\nu_{\\rm ref}$, 7~s for switching, 30~s at $\\nu_{\\rm t}$ and 7~s for switching back to $\\nu_{\\rm ref}$. This is only possible in good weather conditions, when the atmospheric coherence time is 60~s or longer, but the VLBA's dynamic scheduling ensures a high probability. 3) One should use regular scans on a strong calibrator to monitor the phase offset between the frequencies of interest.  Fast frequency switching turns out to be a powerful tool for highest resolution VLBA images of weak sources. In some of the nearest, but unfortunately radio-weak AGNs, it might reveal the jet collimation region on scales of tens of Schwarzschild radii. However, it can also be used at lower frequencies, eg. to phase-reference 15~GHz observations using interleaved 5~GHz scans. Atmospheric coherence and frequency switching times should be less problematic at these frequencies.  Fast frequency switching is basically a phase referencing technique, so it will reference the $\\nu_{\\rm t}$ data to the core position at $\\nu_{\\rm ref}$. After phase transfer, the $\\nu_{\\rm t}$ phase on a given baseline will contain a time-dependent offset which comprises the sum of a constant instrumental term, $\\Delta\\Phi_{\\rm instr}$, and a term due to the source geometry at $\\nu_{\\rm t}$, $\\Delta\\Phi_{\\rm geo}$. If the core at $\\nu_{\\rm t}$ is shifted with respect to the $\\nu_{\\rm ref}$ core, $\\Delta\\Phi_{\\rm geo}$ will be a sinusoid whose period is 24~h, whose amplitude is a measure for the core shift between the two frequencies and whose zeros give the direction of the shift. $\\Delta\\Phi_{\\rm instr}$ and $\\Delta\\Phi_{\\rm geo}$ can be separated by monitoring $\\Delta\\Phi_{\\rm instr}$ on a calibrator and fitting a sinusoid to the residual phase solutions on the source.  Using this technique, it should be possible to look for core shifts due to optical depth effects in radio jets, because the bulk of emission in jets comes from the $\\tau=1$ surface, and this surface is expected to move closer to the central mass when higher frequencies are observed. In low-luminosity AGNs like Sgr~A* and M81, two emission models have been suggested, using either scaled-down jets like in quasars (Falcke et al.  1993), or spherical advection-dominated accretion flows (ADAFs; Rees 1982; Melia 1994; Narayan et al. 1995). These models can be distinguished because only the jet-like emission should have a core shift, and we have proposed to look for it in Sgr~A* and M81 using a variety of frequencies."
    },
    "0207/astro-ph0207048_arXiv.txt": {
        "abstract": "We present new spectroscopic observations of the gravitational arcs and the brightest cluster galaxy (BCG) in the cluster MS2137-23 ($z=0.313$) obtained with the Echelle Spectrograph and Imager on the Keck II telescope. We find that the tangential and radial arcs arise from sources at almost identical redshifts ($z=1.501,1.502$). We combine the measured stellar velocity dispersion profile of the BCG with a lensing analysis to constrain the distribution of dark and stellar matter in the central 100 kpc of the cluster. Our data indicate a remarkably flat inner slope for the dark matter profile, $\\rho_d\\propto r^{-\\beta}$, with $\\beta<0.9$ at 99\\% CL. Steep inner slopes obtained in cold dark matter cosmological simulations -- such as Navarro Frenk \\& White ($\\beta=1$) or Moore (1.5) universal dark matter profiles -- are ruled out at better than $99$\\%CL. As baryon collapse is likely to have steepened the dark matter profile from its original form, our data provides a powerful test of the cold dark matter paradigm at the cluster mass scale. ",
        "introduction": "A fundamental result arising from cold dark matter (CDM) numerical simulations is that the density profiles of DM halos are universal in form across a wide range of mass scales from dwarf galaxies to clusters of galaxies (Navarro, Frenk \\& White 1997, hereafter NFW). Internal to some scale radius $r_{sc}$, the dark matter profile assumes a power law form, $\\rho_{d} \\propto r^{-\\beta}$. Whilst there is some dispute amongst the simulators about the precise value of $\\beta$ with values ranging from 1.0 to 1.5, (Moore et al.\\ 1998, hereafter M98; Ghigna et al.\\ 2000, Power et al. 2002), a clear measurement of $\\beta$ in a range of objects would offer a powerful test of the CDM paradigm.  The largest observational effort in this respect to date has been via dynamical studies of low surface brightness (LSB) and dwarf galaxies, suggesting softer ($\\beta<1$) DM cores than expected on the basis of the numerical simulations (e.g. de Blok \\& Bosma 2002, Salucci \\& Burkert 2000), although the issue remains somewhat controversial (e.g. van den Bosch \\& Swaters 2001). Similar tests have recently been extended to regular spiral (Jimenez, Verde \\& Oh 2002) and elliptical galaxies (Treu \\& Koopmans 2002).  Some observational constraints are available at the scale of massive clusters, from lensing (e.g. Tyson, Kochanski \\& Dell'Antonio 1998; Williams et al.\\ 1999; Smith et al.\\ 2001), X-ray analysis (Mahdavi \\& Geller 2001) and dynamics of cD galaxies (Kelson et al. 2002).  Since massive clusters probe a totally different scale and physical conditions than galaxies, it is crucial to understand their mass distribution to test the universality of the DM profiles.  In this paper we present the first application of a new method to determine the luminous and dark mass distribution in the inner regions of massive clusters with giant arcs around a central BCG. The method combines lensing analysis with stellar kinematical measurements of the BCG. The two ingredients provide complementary information on the relevant scales ($\\sim 100$ kpc), allowing us to disentangle the luminous and dark components of the total mass distribution.  We have chosen the cluster MS2137-23 as a first application of our method since it is an approximately round system, has an isolated BCG and a very well-studied arc system. Fort et al.\\ (1992) first pointed out the potential significance of the radial and tangential gravitational arcs as a means of constraining the mass distribution on $\\simeq$100 kpc scales, and mass models have been developed subsequently (Mellier et al.\\ 1993, Hammer et al.\\ 1997; hereafter M93, H97). The redshifts of the radial and tangential arcs were predicted to lie in the range 1$<z<$2 (M93). A key issue in the earlier work is whether the radial arc is, in fact, a lensed feature. M93 and Miralde-Escud\\'e (1995; hereafter ME95) also point out the importance of determining the stellar velocity dispersion profile of the BCG to weigh the stellar contribution to the mass.  Following the earlier suggestions, we present new observations of the cluster MS2137-23 with the Keck II telescope.  We provide spectroscopic confirmation of the arcs and measure a velocity dispersion profile for the central BCG.  The spectroscopic data are used together with archival HST images to constrain the luminous and DM distribution of the cluster. In the following, $r$ is the radial coordinate in 3-D space, while $R$ is the radial coordinate in 2-D projected space. We adopt H$_0$=65~km s$^{-1}$,Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ for the cosmological parameters.    \\clearpage ",
        "conclusions": "We have secured spectroscopic redshifts for the radial and tangential arcs in the cluster MS 2137-23 and determined the velocity dispersion profile of the BCG to a physical radius of $\\simeq$8 kpc. We have combined these measurements with the lensing geometry in order to construct a self-consistent model of the mass distribution in the cluster core. Using a spherically symmetric luminous and DM mass distribution, we rule out the presence of a DM halo with an inner slope $\\beta>0.9$ at greater than 99\\% confidence, including systematics.  M98 and NFW-type halos with $\\beta\\ge1.0$ are inconsistent with the mass distribution in the core of MS2137-23.  Since the infall of baryons associated with the BCG are likely to {\\em steepen} the DM halo (Blumenthal et al 1986; Mo, Mao and White 1998), our measured profile may imply the original DM profile was even flatter.  A full modeling of this process is beyond the scope of this letter, but this strengthens our conclusion that the inner regions of the DM halo of MS2137-23 cannot be described by CDM-motivated universal halos. A potential concern is that our models have only two mass components, but X-ray emitting gas could be a non-negligible third massive component.  Using ROSAT observations of MS2137-23 (Ettori \\& Fabian 1999) we estimate that removing the X-ray component will steepen the resulting DM halo slope by less than $\\simeq0.1$ and therefore does not change dramatically the result.  Finally, individual halo shapes can depart from the ensemble average behavior. Therefore it is necessary to apply such a test to a sample of clusters. Our simple method is applicable to all approximately round clusters with a massive galaxy at their center, provided that they have at least a giant tangential arc (radial arcs further enhance the sensitivity but are not required). We are in the process of collecting data for a dozen clusters with the aim of performing such a statistical test."
    },
    "0207/astro-ph0207612_arXiv.txt": {
        "abstract": "Criticality in graviational microlensing is an everyday issue because that is what generates microlensing signals which may be of photon-challenged compact objects such as black holes or planetary systems ET calls home. The criticality of these quasi-analytic lenses is intrinsically quadratic, and the critical curve behaves as a mirror generating two mirror images along the image line ($\\parallel \\pm E_-$) at the same distances from the critical curve in the opposite sides.  At the (pre)cusps where the caustic curve ``reflects\" and develops cusps, however, ``would-be\" two pairs of quadratic images ``superpose\" to produce three mirror images because of the degenerate criticality. The critical curve behaves as a parabolic mirror, and the image inside the parabola is indeed a superposed image having the sum of the magnifications of the other two that are outside the parabola. All three images lie on a parabolic image curve shaped by a property ($\\nabla J$) of the (pre)cusp, and {\\it two distances} of the system determine the position of the image curve and positions of the three images on the image curve. The triplet images satisfy $\\sum J^{-1} =0$, and the function $\\sum J^{-1}$ is discontinuous at a caustic crossing where a pair of quadratic images disappear into a critical point. The reflection symmetry of the image curve is a manifestation of the symmetry of the cusp which is also respected by a trio of parabolic curves that are tangnet at the (pre)cusp and define the image domains. The symmetry is guaranteed when $J_{+-}$ vanishes or can be ignored, and the cusps on the lens axis of the binary lenses are strongly symmetric having $J_{+-}=0$ because of the global reflection symmetry of the binary lenses. ``$E_\\pm$-algebra\" is laid out for users' convenience. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207338_arXiv.txt": {
        "abstract": "Separation of the $E$ and $B$ components of a microwave background polarization map or a weak lensing map is an essential step in extracting science from it, but when the map covers only part of the sky and/or is pixelized, this decomposition cannot be done perfectly. We present a method for decomposing an arbitrary sky map into a sum of three orthogonal components that we term ``pure $E$'', ``pure $B$'' and ``ambiguous''. The fluctuations in the pure $E$ and $B$ maps are due only to the the $E$ and $B$ power spectra, respectively, whereas the source of those in the ambiguous map is completely indeterminate. This method is useful both for providing intuition for experimental design and for analyzing data sets in practice. We show how to find orthonormal bases for all three components in terms of bilaplacian eigenfunctions, thus providing a type of polarized signal-to-noise eigenmodes that simultaneously separate both angular scale and polarization type. The number of pure and ambiguous modes probing a characteristic angular scale $\\theta$ scales as the map area over $\\theta^2$ and as the map boundary length over $\\theta$, respectively. This implies that fairly round maps (with short perimeters for a given area) will yield the most efficient $E/B$-decomposition and also that the fraction of the information lost to ambiguous modes grows towards larger angular scales. For real-world data analysis, we present a simple matrix eigenvalue method for calculating nearly pure $E$ and $B$ modes in pixelized maps. We find that the dominant source of leakage between $E$ and $B$ is aliasing of small-scale power caused by the pixelization, essentially since derivatives are involved. This problem can be eliminated by heavily oversampling the map, but is exacerbated by the fact that the $E$ power spectrum is expected to be much larger than the $B$ power spectrum and by the extremely blue power spectrum that CMB polarization is expected to have. We found that a factor of 2 to 3 more pixels are needed in a polarization map to achieve the same level of contamination by aliased power than in a temperature map. Oversampling is therefore much more important for the polarized case than for the unpolarized case, which should be reflected in experimental design. ",
        "introduction": "Detecting polarization of the cosmic microwave background (CMB) radiation has become one of the main goals of the CMB community.  Numerous experimental groups are currently searching for CMB polarization \\cite{Keating01,staggs,hedman,peterson,angelica}. CMB polarization can potentially offer a vast amount of information about our Universe.  In general, polarization in very sensitive to the ionization history of the Universe. For example, on large scales it can provide insight into the way the Universe reionized \\cite{zalreio}. On degree scales, once the temperature anisotropies are well measured, the predicted polarization can serve as a test of how and when recombination happened and could potentially lead to an important confirmation of the Big Bang model \\cite{peebles,landau}. Moreover, because the bulk of the polarization is produced at the last-scattering surface, it should exhibit no correlation on scales larger than about one degree unless there were super-horizon perturbations at decoupling. Polarization can thus become a good test of inflation \\cite{sperzal}.  Most of the recent interest in polarization is based on its ability to provide evidence for a stochastic background of gravity waves. It has been shown that the polarization field on the sky can be decomposed into two parts, a scalar part usually called $E$ and a pseudoscalar part usually called $B$ \\cite{2.kks,3.spinlong}. The pseudoscalar part cannot be created by density perturbations to linear order in perturbation theory. A detection of the $B$ component on large scales would thus indicate the presence of a background of gravity waves, a prediction of inflationary models \\cite{spinlett,kkslett}. Such a detection would determine the energy scale of inflation and could provide a stringent test of inflationary models \\cite{kinney}. On smaller scales, the $B$ modes will most probably be dominated by secondary contributions produced after last scattering, the leading one being gravitational lensing \\cite{pollens}. A detection of these contributions could provide information about the distribution of matter all the way up to the last-scattering surface. There are many proposals for how to detect and use this effect \\cite{jacek,karim,wayne}.  In standard models, however, the $B$ component is likely to be quite difficult to detect \\cite{JKW,maxangelica,ted}.  It is clear that a separation of the observed polarization into $E$ and $B$ parts is crucial to much of the CMB polarization scientific program. It has been realized, however, that real-world complications such as the finite size of the observed patch can significantly reduce our ability to do a clean separation between the two components: when using a quadratic estimator method for measuring the $E$ and $B$ power spectra, substantial ``leakage'' between the two was found on large angular scales \\cite{maxangelica}. In \\cite{ted} it was shown that naive estimates of the sensitivity needed to detect the $B$ component that ignore the such leakage can significantly underestimate the required sensitivity for an experiment aimed at detecting the $B$ modes. In \\cite{lewis} it was shown that in a finite patch, modes that are only $E$ or only $B$ can be constructed but that there are also ambiguous modes, modes that receive contributions to their power from both $E$ and $B$. The construction of the modes was done for a round patch working in harmonic space. It was shown for each value of $m$ there are two ambiguous modes.  The issue of separating $E$ and $B$ has also generated interest in the field of weak gravitational lensing\\cite{kaiserlens,huwhitelens,Crittenden02}, where the basic cosmological signal is expected to produce only an $E$-pattern in cosmic shear maps, and the $B$-mode therefore serves as an important test for other signals due to intrinsic galaxy alignment or systematic errors. Although we do not discuss weak lensing explicitly in this paper, our results are relevant to that case as well since the lensing $E/B$ problem is mathematically analogous.  In this paper we revisit the issue of $E$ and $B$ mode separation, with two goals: to provide intuition for experimental design and for efficiently analyzing data sets in practice. We present a general derivation of the pure $E$, pure $B$ and ambiguous modes in real space, and relate them to the eigenfunctions of the bilaplacian on a finite patch. We then introduce a way to obtain modes that are very nearly ``pure'' in a pixelized map by solving a generalized eigenvalue problem and discuss how this can be used to analyze real-world data sets.  The paper is organized as follows.  Section II establishes some notation and reviews the mathematics underlying the $E/B$ decomposition of a polarization field.  In Section III, we show how to decompose the space of all polarization fields on a finite patch of sky into pure $E$ modes, pure $B$ modes, and modes that are ambiguous with respect to the $E/B$ decomposition.  Section IV presents examples of this decomposition. In Section V, we present a method for finding (nearly) pure $E$ and $B$ modes numerically for pixelized maps by solving a generalized eigenvalue problem.  Section VI presents examples.  In Section VII we show that aliasing of small-scale power is the dominant source of ``leakage'' between the $E$-modes and the $B$-modes.  We summarize our conclusions in Section VIII. ",
        "conclusions": "We have developed a formalism for measuring the $E$ and $B$ components of polarized CMB maps or weak lensing maps given the real-world complications of finite sky coverage and pixelization.  We have shown that by expanding a map in a particular basis, obtained by differentiating bilaplacian eigenfunctions, it can be decomposed as a sum of three orthogonal components that we term pure $E$, pure $B$ and ambiguous.  The pure $E$-component is orthogonal to all $B$-modes and are therefore guaranteed to be caused by an $E$ signal (on the uncut sky), and conversely for the pure $B$-component.  The ambiguous component is the derivative of a biharmonic function, and the original map contains no information about whether it is due to $E$- or $B$-signal in the uncut sky.  We also derived a discrete analogue of these results, applicable to pixelized sky maps.  Our results are useful both for providing intuition for survey design and for analyzing data sets in practice.  \\subsection{Implications for survey design}   To maximize our ability to separate $E$ and $B$, we clearly want to minimize the fraction of modes that are ambiguous. We found that the ambiguous modes are specified along the boundary of the map rather than in the two-dimensional interior. This means that the number of pure and ambiguous modes probing a characteristic angular scale $\\theta$ scales as the map area over $\\theta^2$ and as the map boundary length over $\\theta$, respectively. It is therefore best to minimize the ratio of circumference to area, i.e., to make the patch as round as possible.  Almost all pure modes  (all except the ones with $m=0$ for the spherical cap example) are a combination of both $Q$ and $U$ Stokes parameters, so to achieve unambiguous $E/B$ separation, one needs to measure both, with comparable sensitivity throughout the map.  With pixelized maps, we found that aliasing of small-scale power was a serious problem. Although it can in principle be eliminated by heavily oversampling the map, the required oversampling is greater than for the unpolarized case, both because derivatives are involved and because CMB polarization is expected to have an extremely blue power spectrum. This has important implications for, e.g., the Planck satellite, where bandwidth constraints on the telemetry have been mentioned as reasons to reduce the oversampling.  It is crucial to bear in mind that the usual Nyquist rule-of-thumb that applies to unpolarized maps may be insufficient for realizing the full scientific potential of Planck's CMB polarization measurements because one needs roughly a factor of 2 to 3 more pixels in a polarization map to achieve the same level of contamination by aliased power.   \\subsection{Implications for data analysis}  In \\cite{maxangelica}, it was shown how a quadratic estimator method could produce uncorrelated measurements of the $E$ and $B$ power spectra from real-world data sets with arbitrary sky coverage, pixelization and noise properties, and this method has been applied to both the POLAR \\cite{Keating01} and PIQUE \\cite{angelica} data.  The one annoying problem with this method was that it gave $E/B$ leakage. Our present results allow us to understand and eliminate this problem.  We now know that leakage is caused by the ambiguous modes.  The above-mentioned scaling tells us that the fraction of modes probing a given angular scale $l\\sim \\theta^{-1}$ that are ambiguous scales as $l^{-1}$, in good agreement with the asymptotic behavior empirically found in \\cite{maxangelica}.  Although \\cite{maxangelica} presented a technique for removing most of the leakage, we now know how to remove it completely: by eliminating the ambiguous modes.  In practice, the way to do this is to compute two projection matrices $\\PP_E$ and $\\PP_B$ that project onto the subspaces given by the eigenvectors $\\e$ of \\eq{eqe} with $\\lambda_E>\\lambda_*$ and $\\lambda_E<1/\\lambda_*$, respectively, for some large eigenvalue cutoff $\\lambda_*$, say $\\lambda_*=100$.  The three maps $\\PP_E\\P$, $\\PP_B\\P$ and $[\\I-\\PP_E-\\PP_B]P$ will then be approximately the pure $E$, pure $B$ and ambiguous components of the original map $\\P$, which can be directly used for visual inspection, cross-correlation with other maps and systematic error tests.  To measure the $E$ and $B$ power spectra, one compresses the original data vector $\\P$ into two shorter ones $\\P_E$ and $\\P_B$ by expanding it into the above-mentioned pure $E$ and pure $B$ eigenvectors, respectively. Since this is a mere matrix multiplication, the corresponding noise and signal covariance matrices (which the quadratic estimation method takes as input) are trivially computed as well.  These two data vectors will each have less than half the length of $\\P$. Since the time required by the quadratic estimator method scales as $n^3$, the final $E$ and $B$ power spectrum calculations are therefore about an order of magnitude faster than in the original \\cite{maxangelica} approach.  It should be noted that the ambiguous modes are not useless in all circumstances.  If it has been established that $E$ dominates over $B$ (as is expected theoretically) by observing the pure modes, then it is safe to assume that most of the power in the ambiguous modes is $E$ power as well.  In this case, the ambiguous modes can be used to reduce the errors on estimates of the $E$ power spectrum. This could be particularly useful when attempting to constrain reionization with $E$-power on the very largest angular scales attainable with a galaxy-cut all-sky map, where a substantial fraction of the modes will be ambiguous.    \\vskip 1cm {\\bf Acknowledgments:} The authors thank Ue-Li Pen for asking questions that stimulated this work. Supported was provided by NSF grants AST-0071213, AST-0134999, AST-0098048, AST-0098606 and PHY-0116590, NASA grants NAG5-9194 \\& NAG5-11099, and two Fellowships from the David and Lucile Packard Foundation. MT and EFB are Cottrell Scholars of the Research Corporation."
    },
    "0207/astro-ph0207424_arXiv.txt": {
        "abstract": "After a discussion of the properties of degenerate fermion balls, we analyze the orbit of the star S0-1, which has a projected distance of $\\sim$ 5 light-days to Sgr A$^{*}$, in the supermassive black hole as well as in the fermion ball scenarios of the Galactic center. It is shown that both scenarios are consistent with the data, as measured during the last 6 years by Genzel and coworkers and by Ghez and coworkers. The free parameters of the projected orbit of a star are the unknown components of its velocity $v_{z}$ and distance $z$ to Sgr A$^{*}$ in 1995.4, with the $z$-axis being in the line of sight. We show, in the case of S0-1, that the $z - v_{z}$ phase-space, which fits the data, is much larger for the fermion ball than for the black hole scenario. Future measurements of the positions or radial velocities of S0-1 and S0-2, which could be orbiting within such a fermion ball, may reduce this allowed phase space and eventually rule out one of the currently acceptable scenarios. This could shed some light on the nature of the supermassive compact dark object, or dark matter in general, at the center of our Galaxy. ",
        "introduction": "Self-gravitating degenerate neutrino matter has been suggested as a model for quasars, with neutrino masses in the 0.2 keV $\\simlt m \\simlt$ 0.5 MeV range \\cite{mark1} even before the black hole hypothesis of the quasars was conceived \\cite{lyn16}. More recently, supermassive compact dark objects consisting of weakly interacting degenerate fermionic matter, with fermion masses in the 10 $\\simlt m$/keV $\\simlt$ 20 range, have been proposed \\cite{viol2,bil3,bil4,tsik5,mun6} as an alternative to the supermassive black holes that are believed to reside at the centers of many galaxies.  The masses of $\\sim$ 20 supermassive compact dark objects at the centers of inactive galaxies \\cite{ho7} have been measured so far. The most massive compact dark object ever observed is located at the center of M87 in the Virgo cluster, and it has a mass of $\\sim$ 3 $\\times$ 10$^{9} M_{\\odot}$ \\cite{macc8}. NGC 3115, NGC 4594 and NGC 4374 are galaxies harbouring compact dark objects with the next smaller mass of $\\sim$ 10$^{9} M_{\\odot}$. If we identify the object of maximal mass with a degenerate fermion ball at the Oppenheimer-Volkoff (OV) limit \\cite{opp9}, i.e. $M_{\\rm OV}$ = 0.54$M_{\\rm Pl}^{3} \\; m^{-2} g^{-1/2} \\simeq$ 3 $\\times$ 10$^{9} M_{\\odot}$ \\cite{bil4}, where $M_{\\rm Pl} = \\sqrt{\\hbar c/G}$, this allows us to fix the fermion mass to $\\simeq$ 15 keV for a spin and particle-antiparticle degeneracy factor of $g$ = 2. Such a relativistic object would have a radius of $R_{\\rm OV}$ = 4.45 $R_{\\rm S} \\simeq$ 1.5 light days, where $R_{\\rm S}$ is the Schwarzschild radius of the mass $M_{\\rm OV}$. It would thus be virtually indistinguishable from a black hole of the same mass, as the closest stable orbit around a black hole has a radius of 3 $R_{\\rm S}$ anyway.  At the lower end of the observed mass range are the compact dark objects located at the center of NGC 4945, M32 and our Galaxy \\cite{eck10} with masses of about 1,3 and 2.6 million solar masses, respectively. Interpreting the Galactic object as a degenerate fermion ball consisting of $m \\simeq$ 15 keV and $g$ = 2 fermions, the radius is $R_{\\rm c} \\simeq$ 21 light-days $\\simeq$ 7 $\\times$ 10$^{4} R_{\\rm S}$ \\cite{viol2}, $R_{\\rm S}$ being the Schwarzschild radius of the mass $M_{\\rm c}$ = 2.6 $\\times$ 10$^{6} M_{\\odot}$. Such a nonrelativistic object is far from being a black hole. \\begin{figure}[h] \\begin{center} \\includegraphics[width=.8\\textwidth]{fig1a.ps} \\end{center} \\caption{ Right ascension of S0-1 versus time for various $|z|$ and $v_{x}$ = 340 km s$^{-1}$, $v_{y}$ = - 1190 km s$^{-1}$ and $v_{z}$ = 0 in 1995.4. } \\label{fig1a} \\end{figure} The observed motion of stars within a projected distance of $\\sim$ 5 to $\\sim$ 50 light-days from Sgr A$^{*}$ \\cite{eck10}, the powerful and enigmatic radio source at the Galactic center, yields, apart from the mass, an upper limit for the radius of the fermion ball $R_{\\rm c} \\simlt$ 22 light days. Matter orbiting in an optically thick and geometrically thin accretion disk in or around such a fermion ball will only emit radiation at distances larger than $\\sim$ 10 mpc from the center, as both the density and the circular frequency become nearly constant near the center of the fermion ball \\cite{bil3}. The spectrum emitted by the disk will thus have a cut-off at frequencies larger than $\\sim$ 10$^{13}$ Hz, as is actually observed. Of course, there will be a pile-up and instability of matter within $\\sim$ 10 mpc, perhaps leading to the formation of stars, as the gravitational tidal forces on nascent stars is much smaller in the fermion ball than in the black hole scenario. These stars may be eventually ejected from the central star cluster by intruder stars in close binary encounters. The formation of such a fermion ball as well as its coexistence at finite temperature with a Galactic halo composed of the same fermions has been discussed by Lindebaum \\cite{lind17} and Bili\\'{c} \\cite{bil18} respectively, at this conference. \\begin{figure}[h] \\begin{center} \\includegraphics[width=.8\\textwidth]{fig2a.ps} \\end{center} \\caption{ Declination  of S0-1 versus time for various $|z|$ and $v_{x}$ = 340 km s$^{-1}$, $v_{y}$ = - 1190 km s$^{-1}$ and $v_{z}$ = 0 in 1995.4. } \\label{fig2a} \\end{figure} The required weakly interacting fermion of $\\sim$ 15 keV mass cannot be an active neutrino, as it would overclose the Universe by orders of magnitude \\cite{kolb11}. However, the $\\sim$ 15 keV fermion could very well be a sterile neutrino, contributing $\\Omega_{\\rm d} \\simeq$ 0.3 to the dark matter fraction of the critical density today. Indeed, as has been shown for an initial lepton asymmetry of $\\sim$ 10$^{-3}$, a sterile neutrino of mass $\\sim$ 10 keV may be resonantly produced in the early Universe with near closure density, i.e. $\\Omega_{\\rm d} \\sim$ 1 \\cite{shi12}. As an alternative possibility, the required $\\sim$ 15 keV fermion could be the axino \\cite{goto13} or the gravitino \\cite{lyth14} in soft supersymmetry breaking scenarios. ",
        "conclusions": ""
    },
    "0207/astro-ph0207430_arXiv.txt": {
        "abstract": "The supernova type Ia observational data are fitted using a model with cold dark matter and the Chaplygin gas. The Chaplygin gas, which is characterized by a negative pressure varying with the inverse of density, represents in this model the dark energy responsible for the accelaration of the universe. The fitting depends essentially on four parameters: the Hubble constant, the velocity of sound of the Chaplygin gas and the fraction density of the Chaplygin gas and the cold dark matter. The best fitting model is obtained with $H_0 = 65\\,km/Mpc.s$, $c_s^2 \\sim 0.92\\,c$ and $\\Omega_{c0} = 1$, $\\Omega_{m0} = 0$, that is, a universe completely dominated by the Chaplygin gas. This reinforces the possibility that the Chaplygin gas may unify dark matter and dark energy, as it has already been claimed in the literature. \\vspace{0.7cm} \\newline PACS number(s): 98.80.Bp, 98.65.Dx, 98.80.Es ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207606_arXiv.txt": {
        "abstract": "If X-ray flashes are due to the forward jet emissions from gamma ray bursts (GRBs) observed with large viewing angles, we show that a prompt emission from a counter jet should be observed as a {\\it delayed flash} in the UV or optical band several hours to a day after the X-ray flash. Ultraviolet and Optical Telescope on {\\it Swift} can observe the delayed flashes within $\\sim13$ Mpc, so that (double-sided) jets of GRBs may be directly confirmed. Since the event rate of delayed flashes detected by {\\it Swift} may be as small as $\\sim6\\times10^{-5}$events yr$^{-1}$, we require more sensitive detectors in future experiments. ",
        "introduction": "Several observations suggest that gamma-ray bursts (GRBs) are caused by relativistic jets (e.g., Frail et al. 2001). However, in order to establish that GRBs are collimated, other observations are indispensable, such as polarization observations (Ghisellini \\& Lazzati 1999; Sari 1999) and microlensing observations (Ioka \\& Nakamura 2001b). Some theoretical models of jet emissions have been discussed (Totani \\& Panaitescu 2002; Huang, Dai \\& Lu 2002; Dado, Dar \\& De~R\\'{u}jula 2001). If GRBs are due to forward jet emissions, there should most likely be counter jet emissions, as in the AGN (Begelman, Blandford \\& Rees 1984) and the microquasar (Mirabel \\& Rodr\\'{\\i}guez 1999). Therefore the detection of counter jet emissions will give us direct evidence for the jet model of GRBs.  The confirmation of a counter jet has been by far the most important factor in the jet model of astrophysical objects. A mysterious spot was found in SN1987A using the speckle technique (Meikle et al. 1987; Nisenson et al. 1987). Many models including the jet model were proposed (Rees 1987; Piran \\& Nakamura 1987). At that time, it was difficult to distinguish each model from observations since only one spot was found. In the jet model, the counter jet should be observed although it is dim due to redshifting (Piran \\& Nakamura 1987). However, later in 1999, two spots were confirmed using new software to analyze the speckle data (Nisenson \\& Papaliolios 1999). Very recently the jet feature of the ejecta of SN1987A whose position angle is the same as the mysterious spot was confirmed by the HST (Wang et al. 2002).  As a result the jet model by Piran \\& Nakamura (1987) took the advantage. Furthermore the observation of a counter jet may enable us to estimate the Lorentz factor of the jet, as for the AGN and microquasar. Therefore it is important to argue the observational properties of the emission from the counter jet of a GRB.  Let us consider the emission from a counter jet with a Lorentz factor $\\gamma$. The observed typical frequency of the counter jet emission is about $\\gamma^2$ times smaller than that of the forward jet (i.e., the GRB). Since the typical frequency of the GRB is $\\sim100$ keV, the typical observed frequency of the counter jet emission is $\\sim10(\\gamma/100)^{-2}$ eV, which is in the UV or optical band. This transient phenomenon should be observed about several tens of hours after the forward jet emission, since it is at a radius of order $10^{14}$--$10^{15}$ cm that photons are emitted from each jet leaving, almost simultaneously, the central engine. We call this event the {\\it delayed flash} (DF).  Any attempt to detect the DF might be difficult since the afterglow of the forward jet might be brighter than the DF. The afterglow of the GRB, i.e., the afterglow of the on-axis forward jet, is much brighter than the DF. However if the forward jet is observed with a large viewing angle, there is a chance to observe the DF since the forward jet emission is also dim at the time of the DF.  Recently, we studied the emission from the off-axis jet (Yamazaki, Ioka, \\& Nakamura 2002a, b; see also Nakamura 2000; Ioka \\& Nakamura 2001a). We proposed that if we observe a GRB with a large viewing angle, it looks like an X-ray flash (XRF), a new class of X-ray transients which has been recently recognized as a phenomenon related to the GRB (e.g., Heise et al. 2001; Kippen et al. 2002; Barraud et al. 2002). The off-axis jet model can explain the typical observed frequency and other observational characteristics of the XRF, such as the peak flux ratio and the fluence ratio between the $\\gamma$-ray and the X-ray band, the X-ray photon index, the typical duration, and the event rate (Yamazaki, Ioka, \\& Nakamura 2002a, b). Although the origin of XRFs is uncertain, we assume that XRFs arise from prompt off-axis jet emissions hereafter.   In this paper, we will show that the DF can be observed after the XRF in principle. We will calculate the light curves of the XRF, DF, and the afterglow of the XRF, and discuss whether the DF can be detected by the {\\it Swift} satellite. We will find that we need more sensitive detectors to detect the DF. In \\S~\\ref{sec:model}, we describe a simple forward-/counter-jet model for the XRF and DF. In \\S~\\ref{sec:lightcurve} and \\S~\\ref{sec:afterglow}, we show the light curves of the XRF, the DF, and the XRF afterglow. \\S~\\ref{sec:dis} is devoted to a discussion. ",
        "conclusions": "\\label{sec:dis}  We have calculated the light curves of the DF, XRF, and the afterglow of the XRF. We have shown that in principle, the DF emission can be seen in the UV band about $10^4$--$10^5$ seconds after the XRF if the viewing angle is large enough (about 0.2--0.3 rad) for the afterglow of the XRF to be dimmer than the DF. Since the UV flux of the GRB afterglow is much larger than that of the DF, only the DF associated with an off-axis jet, i.e., an XRF has any chance to be observed. The preceding XRF should have a low peak energy of a few keV, a small variability, and a short duration for the DF to be detected. Due to the relativistic beaming effect, the flux of the DF is so small that only nearby events ($\\lesssim13$~Mpc for the canonical parameters) can be observed by UVOT on {\\it Swift}. Following Yamazaki, Ioka, \\& Nakamura (2002a), we can roughly estimate the event rate of the DF for the instruments on {\\it Swift} as $R_{DF}\\sim 6\\times10^{-5}\\,{\\rm events}\\ {\\rm yr}^{-1}$, where we adopt the event rate of the GRBs $r_{GRB}={5\\times 10^{-8}\\,{\\rm events}\\ {\\rm yr}^{-1}\\ {\\rm galaxy}^{-1}}$ and the number density of galaxies $n_g=10^{-2}\\,{\\rm galaxies}\\ {\\rm Mpc}^{-3}$. Therefore, we need next-generation detectors, which are more sensitive than the instruments on {\\it Swift}, to detect the DF associated with very dim XRF more frequently.  The DF may be obscured by dust extinction. In fact, about half of accurately localized GRBs do not produce a detectable optical afterglow (Fynbo et al. 2001; Lazzati, Covino, \\& Ghisellini 2002). One explanation for these ``dark GRBs'' is that most GRBs occur in giant molecular clouds (e.g., Reichart \\& Price 2002). In this picture, a GRB has a detectable optical afterglow only if the burst and the afterglow destroy the dust along the line of sight to the observer (Waxman \\& Draine 2000; Fruchter, Krolik, \\& Rhoads 2001), as suggested by the comparison between X-ray and optical extinction (Galama \\& Wijers 2001). In this case the DF is obscured since the flux of the XRF is too dim to carve out a path for the DF. However this picture may have some problems, such as no evidence of an ionized absorber (Piro et al. 2002) and variable column density in the X-ray afterglow (Djorgovski et al. 2001c). There are other explanations for dark GRBs, such as high redshift effects, dust extinction in the interstellar medium of the host galaxy (Ramirez-Ruiz, Trentham, \\& Blain 2002; Piro et al. 2002) and so on. Therefore at present we cannot conclude that the DF is obscured.   If we assume that the absolute magnitude of the host galaxy is about $\\sim -20$ mag (Djorgovski et al. 2001a, b), the apparent magnitude is about $\\sim 20 + 5 \\log D_{\\rm Gpc}$. Since a host galaxy with a size $\\sim 10$ kpc has an angular size of $\\sim 10 D_{\\rm Gpc}^{-1}$ arcsec, we can observe a point source which is dimmer than the host galaxy by $\\sim 10^{-4} D_{\\rm Gpc}^{2}$ if the angular resolution is $\\sim 0.1$ arcsec. Therefore the DF has to be brighter than $\\sim 30$ mag, and we can observe the DF if $D \\siml 13$ Mpc.  If the GRB is associated with a supernova (SN), the emission from the SN may hide the DF. The UV flux of SN1998bw was about $\\sim 17$ mag at the distance $D \\sim 40$ Mpc (Galama et al. 1998), i.e., $\\sim 6\\times 10^{-15} D_{\\rm Gpc}^{-2}$ erg s$^{-1}$ cm$^{-2}$, so that a SN like SN1998bw is brighter than the DF. However at present it is not clear whether all GRBs are associated with the SNe or not (e.g., Price et al. 2002). In any cases, deep searches following the XRF will give us valuable information.  If the DF associated with an XRF is observed, we will be able to estimate the Lorentz factor and the viewing angle of the jets. Let the typical frequency or the break energy of the DF (XRF) be $\\nu_{DF}=\\delta_{DF}\\nu'_0$ ($\\nu_{XRF}=\\delta_{XRF}\\nu'_0$), where $\\delta\\equiv 1/\\gamma(1-\\beta\\cos\\theta_v)$ is the Doppler factor. When $\\theta_v\\ll1$, $\\gamma\\gg1$ and $(\\gamma\\theta_v)^2\\gg1$, we can derive $\\delta_{DF}\\sim1/(2\\gamma)$ and $\\delta_{XRF}\\sim2\\gamma/(\\gamma\\theta_v)^2$. Since we assume that the XRF is the GRB observed from the off-axis viewing angle, we may use the observational consequence for the break energy $\\delta_{GRB}h\\nu'_0\\sim200\\,\\xi$ keV, where $\\xi\\sim$ 0.5--2 (Preece et al. 2000). In our model, $\\delta_{GRB}$ becomes $\\sim2\\gamma$. Then, we obtain $\\gamma\\sim100\\,\\xi^{1/2}(h\\nu_{DF}/5{\\rm eV})^{-1/2}$. On the other hand, we can derive $\\nu_{DF}/\\nu_{XRF}\\sim(\\theta_v/2)^2$, which implies that we can also estimate the viewing angle."
    },
    "0207/astro-ph0207661.txt": {
        "abstract": "We have reanalyzed the existing data on Zinc abundances in damped Ly$\\alpha$ (DLA) absorbers to investigate whether their mean metallicity evolves with time. Most models of cosmic chemical evolution predict that the mass-weighted mean interstellar metallicity of galaxies should rise with time from a low value $\\sim 1/30$ solar at $z \\sim 3$ to a nearly solar value at $z \\sim 0$. However, several previous analyses have suggested that there is little or no evolution in the global metallicity of DLAs. The main problem is that the effective number of systems that dominate the $N({\\rm H \\, I})$-weighted mean metallicity is very small.  We have used a variety of statistical techniques to quantify the global metallicity-redshift relation and its uncertainties, taking into account both measurement and sampling errors. Three new features of our analysis are: (a) an unbinned $N({\\rm H \\, I})$-weighted nonlinear $\\chi^{2}$ fit to an exponential relation; (b) survival analysis to treat the large number of limits in the existing data; and (c) a comparison of the data with several models of cosmic chemical evolution based on an unbinned $N({\\rm H \\, I})$-weighted $\\chi^{2}$. We find that a wider range of evolutionary rates is allowed by the present data than claimed in previous studies. The slope of the exponential fit to the $N({\\rm H \\, I})$-weighted  mean Zn metallicity vs. redshift relation is $ -0.20 \\pm 0.11$ counting limits as detections and $ -0.27 \\pm 0.12$ counting limits as zeros. Similar results are also obtained if the data are binned in redshift, and if survival analysis is used. These slopes are marginally consistent with no evolution, but are also consistent with the rates predicted by several models of cosmic chemical evolution (e.g., slopes of $-0.61$ to $-0.25$ for the models of Pei \\& Fall 1995, Malaney \\& Chaboyer 1996, and Pei et al. 1999). The $\\chi^{2}$ values obtained for most of these models are somewhat worse than that for the exponential model because the models lie above the observed data points, but still suggest that the present DLA data could indicate some evolution of the metallicity with redshift. Finally, we outline some future measurements necessary to improve the statistics of the global metallicity-redshift relation. ",
        "introduction": "The evolution of stars and gas in galaxies are topics of great interest in modern astrophysics. The average star-formation history of the universe has been estimated from emission properties of galaxies detected in deep imaging and redshift surveys (e.g. Lilly et al. 1996; Madau et al. 1996, 1998). This emission history of galaxies is connected intimately with the histories of gas consumption and metal production, because the global densities of gas, metals, and stars are coupled through conservation-type relations (e.g., Fall 2001). In particular, because the global rate of star formation is known to be high at $1 \\lesssim z \\lesssim 4$, we expect the mean interstellar metallicity of galaxies to rise rapidly in that interval. Direct observational constraints on the evolution of the mean metallicity in galaxies are therefore important for pinning down the histories of star formation and gas consumption.  Abundance measurements in gas traced by quasar absorption lines can directly probe the evolution of metals in galaxies. The damped Ly$\\alpha$ (DLA) absorbers (log $N({\\rm H \\, I}) \\gtrsim 20$) are particularly important, since they contain a large fraction of the neutral gas in galaxies, nearly enough to form all of the stars visible today (e.g., Wolfe et al. 1995). DLAs are the only class of high-redshift objects in which abundances of a large number of elements have been measured (e.g., Meyer \\& York 1992; Pettini et al. 1994, 1997, 1999, 2000; Lu et al. 1996; Kulkarni et al. 1996, 1997, 1999; Prochaska \\& Wolfe 1996, 1997, 1999; Prochaska et al. 2001b; and references therein). In previous studies, Zn II has been the most commonly used ion for estimating the total (gas + solid phase) metallicity in DLAs, for a number of reasons: (1) Zn is relatively undepleted on interstellar dust grains; (2) it tracks Fe closely in Galactic stars (for [Fe/H] $\\gtrsim -2$); (3) its absorption lines are usually unsaturated and often lie outside the Ly$\\alpha$ forest; and (4) ionization corrections are relatively small for Zn II. Abundances of depleted elements such as Cr or Fe relative to Zn are used to estimate the dust content of DLAs.  The quantity of interest here is the $N({\\rm H \\, I})$-weighted mean metallicity, which corresponds to the global $\\Omega_{\\rm{metals}}^{\\rm ISM}/\\Omega_{\\rm{gas}}^{\\rm ISM}$ ratio in galaxies. The average interstellar properties of galaxies can be determined from the statistics of quasar absorption-line systems as follows. Let $f(N_x,z)$ be the distribution in column density and redshift, for particles of any type $x$ that absorb or scatter light. These might, for example, be hydrogen atoms ($x=$~H I), metal ions ($x=m$), or dust grains ($x=d$). By definition, $H_0 (1+z)^3 |dt/dz| f(N_x,z)dN_xdz$ is the mean number of absorption-line systems with column densities of $x$ between $N_x$ and $N_x+dN_x$ and redshifts between $z$ and $z+dz$ along the lines of sight to randomly selected background quasars. These lines of sight are very narrow (the projected size of the continuum- emitting regions of quasars, less than a light-year across) and pierce the absorption-line systems at random angles and impact parameters. One can show that the mean comoving density of $x$ is given by \\begin{equation} \\Omega_x(z) = {8 \\pi G m_x \\over 3 c H_0} \\int_0^\\infty N_x f(N_x,z) dN_x, \\end{equation} where $m_x$ is the mass of a single particle (atom, ion, or grain). Equation~(1) plays a central role in this subject. It enables us to estimate the mean comoving densities of many quantities of interest without knowing anything about the structure of the absorption-line systems. In particular, we do not need to know their sizes or shapes, whether they are smooth or clumpy, and so forth. Furthermore, in the absence of selection biases, equation (1) gives exactly the correct weighting to lines of sight at different impact parameters, i.e. distances from the centers of galaxies. This point is sometimes confused in the literature and even the opposite is sometimes asserted (e.g., Edmunds \\& Phillips 1997). A corollary of equation~(1) is that the global interstellar metallicity, $Z\\equiv\\Omega_{\\rm metals}^{\\rm ISM}/\\Omega_{\\rm gas}^{\\rm ISM}$, is given simply by an average over the metallicities of individual absorption-line systems weighted by their H I column densities.  Most models of cosmic chemical evolution predict that the global interstellar metallicity rises with time, from a low value at high redshifts to a near-solar value at zero redshift (e.g., Lanzetta, Wolfe, \\& Turnshek 1995; Pei \\& Fall 1995; Malaney \\& Chaboyer 1996; Pei, Fall, \\& Hauser 1999; Tissera et al. 2001). Moreover, from emission-line observations of nearby galaxies, it appears that the present-day mass-weighted mean interstellar metallicity of galaxies, averaged over all morphological types, is near-solar. We show this explicitly in the Appendix using the mean relation between galaxy luminosity and interstellar metallicity as derived from H II regions, and integrating over the luminosity function of galaxies. However, it is not clear from the absorption-line observations whether or not the mean metallicity in DLAs actually increases with decreasing redshift as predicted by the models.  There have been several attempts to estimate the $N({\\rm H \\, I})$-weighted mean Zn metallicity of DLAs at $0.4 \\lesssim z \\lesssim 3.5$ (e.g., Pettini et al. 1997, 1999; Vladilo et al. 2000; Savaglio 2001). These studies have claimed that there is little or no evolution in the global Zn metallicity of DLAs in this redshift range. Prochaska \\& Wolfe (1999, 2000, 2002) and Prochaska, Gawiser, \\& Wolfe (2001a) have found no evolution in the mean Fe abundance for $2 \\lesssim z \\lesssim 4$. However, these previous studies have not made consistent quantitative estimates of the mean slope of the global metallicity-redshift relation and the observational uncertainties in that relation. There are, in fact, several uncertainties in the present DLA metallicity data. First of all, the metallicities for individual DLAs show a large intrinsic or cosmic scatter at any given redshift, reflecting the different rates of chemical enrichment of different galaxies. Furthermore, the current samples are relatively small. More importantly, the estimates of the $N({\\rm H \\, I})$-weighted mean metallicity for these samples are dominated by only the few DLAs with the highest $N({\\rm H \\, I})$. To discriminate between evolution vs. no evolution from these intrinsically noisy samples, it is crucial to estimate quantitatively the mean metallicity-redshift relation and the effect of the observational uncertainties on that relation. Here we present a reexamination of the Zn data with quantitative estimates of the mean metallicity-redshift relation and comparison with predictions of cosmic chemical evolution models. We include several recent low-redshift measurements, which are crucial for determining the slope of the metallicity-redshift relation. Using a number of different statistical techniques, we find that the existing data could support significant evolution of the global metallicity with redshift. ",
        "conclusions": "Most previous studies have claimed that there is no evolution in the global metallicity of DLAs. However, these studies have not been definitive for various reasons summarized in Table 1. Our analysis has demonstrated that the present DLA Zn data are consistent with some evolution of the global interstellar metallicity. The main reason for the large uncertainties is that the effective number of measurements that dominate the $N({\\rm H \\, I})$-weighted  mean metallicity is very small.  A complete lack of evolution in the mean metallicity would be quite surprising. Such a trend would be hard to reconcile with the inference that the global rate of star formation was high at $1 \\lesssim z \\lesssim 4$, based on the luminosity density of galaxies observed in deep surveys such as the Canada-France Redshift Survey and the Hubble Deep Field (e.g. Lilly et al. 1996; Madau et al. 1996, 1998). The metallicity of interstellar matter in galaxies is thus expected to rise with time. Indeed, the census of the metals in nearby galaxies shows a luminosity-weighted mean interstellar metallicity close to the solar value. (See the Appendix.) DLAs are believed to represent the interstellar matter of galaxies, and are therefore expected to show an increase in the mean metallicity with decreasing redshift. As we have shown, the DLA data are, in fact, consistent with the metallicity evolution predicted by some cosmic chemical evolution models.  \\subsection{Dust Obscuration Bias}  Dust in DLA galaxies could bias the empirical estimates of the mean metallicity $\\bar Z$ and its evolution with redshift. This is because the DLA galaxies with the highest column densities of metals are likely to be those with the highest column densities of dust, and these may obscure background quasars to such a degree that some of them are omitted from optically selected samples (Fall \\& Pei 1993; Boisse et al. 1998). Evidence for dust in DLA galaxies comes from the statistical reddening of background quasars (Fall, Pei, \\& McMahon 1989; Pei, Fall, \\& Bechtold 1991) and the depletion patterns of heavy elements, especially Zn and Cr (e.g., Pettini et al. 1994, 1997). Like the mean metallicity, the mean dust-to-gas ratio in the observed DLA galaxies is low at high redshifts, where most of the measurements have been made. The mean dust-to-metals ratio appears to be roughly independent of redshift and about equal to that in the Milky Way and the Magellanic Clouds (see Figure 1 of Pei et al. 1999). As a result of obscuration, the observed mean metallicity in the DLA galaxies, and hence the data points in Figure 1, may lie systematically below the true mean metallicity. The curves in Figure 1, however, are predictions for the true mean metallicity, without corrections for obscuration, from the models of cosmic chemical evolution.  The severity of this bias depends on several factors, including the extinction curve of the dust, the distribution of dust column densities in the DLA galaxies, the luminosity function of quasars, and the passband in which they are observed (see the Appendix of Fall \\& Pei 1993 for a detailed analysis). Together, these factors determine the fraction of the sky covered by dust as a function of the optical depth. The main difficulty in correcting for the bias stems from the unknown distribution of dust column densities in the DLA galaxies or, for a given (observed) distribution of H I column densities, the unknown distribution of dust-to-gas ratios. For an assumed shape of this distribution, one can, however, compute the expected bias as a function of the width of the distribution. Fall \\& Pei (1993) have made such calculations for a log-normal distribution of the dust-to-gas ratio $k$ in the DLA galaxies. They find that the true mean dust-to-gas ratio $\\bar k_t$ exceeds the observed mean $\\bar k_o$ by a factor that increases from $\\bar k_t/\\bar k_o = 1.2$ to 1.7 to 4 as the dispersion in the natural logarithm of the dust-to-gas ratio increases from $\\sigma(\\ln k) = 0.5$ to 1.0 to 1.5. This may provide an indication of the corresponding bias in the mean metallicity, since we expect $\\bar Z_t/\\bar Z_o \\approx \\bar k_t/\\bar k_o$. For reference, nearby normal galaxies have $\\sigma (\\ln k) \\approx 0.5$ (Pei 1992). Thus, if this dispersion also applies to the DLA galaxies, we might expect $\\bar Z_t \\approx \\bar Z_o$. If the dispersion in the dust-to-gas ratio were larger, however, $\\bar Z_t$ could differ significantly from $\\bar Z_o$.  Unfortunately, the bias caused by obscuration is difficult to quantify from theory or numerical simulation alone, because it depends on the the small-scale structure of the interstellar medium in the DLAs. Absorption-line observations sample the DLAs on scales comparable to the continuum-emitting regions of the background quasars, typically smaller, and possibly much smaller, than a light-year. Some of these lines of sight will pass through dense interstellar clouds, with high optical depths, while others, even nearby, will pass through diffuse intercloud material, with low optical depths. The obscuration of background quasars in this case may differ substantially from that in simulations with the same average interstellar density and hence optical depth but with lower spatial resolution and hence little or no structure on the relevant scales.  In principle, a comparison of metal abundances in DLAs from optical vs. radio-selected quasars can help to quantify the dust selection effect. Ellison et al. (2001) find slightly more DLAs in the foreground of radio-selected and optically faint quasars than in optically selected and optically bright quasars, in the sense that may be caused by dust obscuration. However, the sample of radio-selected quasars is still small, and neither of these differences is statistically significant. The strongest conclusion that can be drawn at present is that obscuration reduces the number density and/or $\\Omega_{\\rm H \\, I}$ of DLAs at $1.8 < z < 3.5$ in optically selected samples by a factor of two or less. However, the implications of this result for the mean metallicity have not yet been quantified. Abundance studies for a large sample of DLAs in radio-selected and optically faint quasars will help to improve the constraints on the extent of the dust obscuration bias.  \\subsection{Iron vs. Zinc}  All of our analysis has been based on Zn alone, since we believe Zn is a more reliable metallicity indicator for DLAs than Fe, which has been used in some other studies (e.g., Prochaska \\& Wolfe 1999, 2000; Savaglio 2001; Prochaska et al. 2001a). The main reasons advocated by these studies for using Fe are: (a) the ease of measuring Fe lines compared to the weaker Zn lines; (b) the ability to probe higher redshifts ($z > 3.5$) with Fe than with Zn; and (c) the better understanding of the nucleosynthetic origin of Fe than of Zn. We address these issues one by one. (a) As our ability to measure weak absorption lines is getting better in the present age of large telescopes, Zn measurements are becoming easier. Furthermore, this is not a concern for studies of the global $N({\\rm H \\, I})$-weighted metallicity as the systems that are too weak to give detections of Zn lines do not contribute much to the global metallicity anyway. As our analysis has shown, there would be almost no difference in the results even if the limits could be improved in future studies (as there is little difference whether the limits are treated as detections or zeros). (b) The redshift range $z > 3.5$ accessible with Fe is of some interest since the rate of metallicity evolution at high redshifts could in fact be different from that at low redshifts. However, the redshift range $z > 3.5$ represents only $13 \\, \\%$ of the cosmic history. Therefore, even though Fe is useful for tracing the early chemical evolution of galaxies, it does not add much to studies of the last $87 \\, \\%$ of the age of the universe. (c) While the nucleosynthetic origin of Zn is harder to understand than that of Fe, observationally Zn and Fe track each other perfectly well for most Galactic disk and halo stars with metallicities between $10^{-2}$ solar and solar, the metallicity range relevant to DLAs (e.g., Sneden, Gratton, \\& Crocker 1991). Some models of massive star explosive nucleosynthesis, such as the neutrino-driven wind models (e.g., Hoffman, Woosley, \\& Qian 1997) provide ways to understand the origin of Zn and the reason for the tight correlation between Fe and Zn in Galactic stars. Thus, there is no strong reason to suspect that the use of Zn biases studies of cosmic chemical evolution in any significant way.  On the other hand, the well-known problem with Fe is its strong depletion on dust grains. Moreover, as with many other elements, the depletion of Fe differs substantially between the cool and warm diffuse interstellar gas, but is strong in both phases (see, e.g., Savage \\& Sembach 1996). By comparison, Zn is essentially undepleted in the warm interstellar clouds and depleted $\\sim 40$ times less than Fe in the cool interstellar clouds. It is difficult to model unambiguously the dust depletion effects for Fe, because the structure and composition of the dust grains present in the DLAs is not known a priori. Furthermore, the line of sight can pass through a mixture of warm and cold gas, so that the dust depletion within a given DLA can be quite different in different parts of the absorbing gas. Even after averaging over a number of DLAs at a given redshift, it is possible that the mean correction for dust depletion may itself change as a function of redshift. Indeed, as the interstellar metallicity of the DLA galaxies rises with decreasing redshift, their dust-to-gas ratio should also rise more or less in step with the metallicity. Such a redshift-dependent dust depletion can introduce an error in estimates of the metallicity-redshift relation based on Fe data alone, or on a combination of Fe and Zn data. Although the extent of this error is hard to quantify and, in fact, may turn out to be small, it is safer to focus on an element such as Zn that does not have this problem in the first place.  \\subsection{Future Work}  Part of the reason our analysis finds a stronger evidence for evolution than previous studies such as Pettini et al. (1999) is the addition of new data at low redshifts that have recently become available. These new data have resulted in a higher mean metallicity in the lowest redshift bin than found in the earlier analyses. This highlights the need for caution in drawing conclusions from the limited DLA data sets. Indeed, it is obvious that the current samples need significant improvement in the number of measurements at $z < 2$. A large fraction of the Zn measurements so far have focussed on the redshift range $z > 2$. Possible drops in the mean Zn metallicity at $3 < z < 3.5$ and the mean Fe metallicity at  $3.5 < z < 4.5$ have been suggested (although on the basis of small samples) by Pettini et al. (1997) and Prochaska \\& Wolfe (2002), respectively. If verified with future data, a drop in the mean metallicity at high redshifts could signal the epoch of the onset of star formation in DLAs. It is, nevertheless, also essential to increase substantially the lower redshift samples, since the redshift range $z < 2$ probes the cosmic epochs when the bulk of the metals in galaxies were produced. At present, Zn measurements exist for only seven absorbers at $z < 1$ and only two absorbers at $z < 0.5$. This is especially problematic because the redshift ranges $z < 1$ and $z < 0.5$ represent $57 \\, \\%$ and $37 \\, \\%$, respectively, of the age of the universe. Even the number of measurements at $1 < z < 2$ is somewhat limited. Clearly, it is crucial to increase the number of Zn measurements at low and intermediate redshifts. This requires more measurements at $z < 0.6$ with the {\\it Hubble Space Telescope} (HST) and at $0.6 < z < 2$ with ground-based telescopes. It is particularly important to study the high-$N({\\rm H \\, I})$ systems since, as we have emphasized here, these systems dominate the global metallicity.  The large numbers of quasar spectra becoming available with surveys such as the Sloan Digital Sky Survey and the FIRST survey will increase the sample of known DLAs and 21-cm absorbers by a large factor. Obtaining element abundances of these new DLAs will be of great importance for pinning down the metallicity-redshift relation for DLAs. Zn abundances in  DLAs in front of quasars of different apparent magnitudes and colors will provide constraints on the amount of dust obscuration. Accurate determination of the metallicity-redshift relation for DLAs will, thus, be important for quantifying the selection effects in samples of quasar absorbers and for understanding the nature of DLAs, in addition to providing important constraints on the global star formation history of galaxies."
    },
    "0207/astro-ph0207095_arXiv.txt": {
        "abstract": "We investigate the predicted present-day temperature profiles of the hot, X-ray emitting gas in galaxy clusters for two cosmological models - a current best-guess  $\\Lambda$CDM model and  standard cold dark matter (SCDM). Our numerically-simulated ``catalogs\" of clusters are derived from  high-resolution (15 h$^{-1}$kpc) simulations which make use of a sophisticated, Eulerian-based, Adaptive Mesh-Refinement (AMR) code that faithfully captures the shocks which are essential for correctly modelling cluster temperatures. We show that the temperature structure on Mpc-scales is highly complex and non-isothermal. However, the temperature profiles of the simulated $\\Lambda$CDM and SCDM clusters are remarkably similar and drop-off as $T \\propto (1+r/a_x)^{-\\delta}$ where $a_x \\sim r_{vir}/1.5$ and $\\delta \\sim 1.6$. This decrease is in good agreement with the observational results of Markevitch et al.~(1998) but diverges, primarily in the innermost regions, from their fit which assumes a  polytropic equation of state. Our result is also in good agreement with a recent sample of clusters observed by {\\it Beppo}SAX though there is some indication of missing physics at small radii ($r<0.2 \\; r_{vir}$). We discuss the interpretation of our results and make predictions for new x-ray observations that will extend to larger radii than previously possible. Finally, we show that, for $r>0.2 \\; r_{vir}$, our universal temperature profile is consistent  with our most recent simulations which include both radiative cooling and supernovae feedback. ",
        "introduction": "X-ray based cluster mass determinations generally assume that morphologically-symmetric (i.e. non-merging) clusters are isothermal. However, this accepted wisdom has recently been challenged by Markevitch et al.~(1998; hereafter M98) who found evidence for decreasing temperature profiles in a sample of nearby hot clusters ($>3.5$ keV) observed with ASCA. A subsample of 17 regular/symmetric clusters displayed remarkably similar temperature profiles (when normalized and scaled by the virial radius) consistent with $T\\propto [1+(r/r_c)^2]^{-3 \\beta (\\gamma-1)/2}$ where $\\gamma=1.24^{+.20}_{-.12}$ and $\\beta=2/3$. The typical decrease is therefore a factor of $\\sim$2 in going from 1 to 6 core radii (or .09 to 0.5 virial radii). This result remains controversial as three subsequent studies of large samples of clusters concluded that the majority of cluster temperature profiles show little, or no, decrease with radius (Irwin, Bregman, \\& Evrard 1999; White 2000; Irwin \\& Bregman 2000). Most recently, De Grandi \\& Molendi (2002) have presented a composite temperature profile based on {\\it Beppo}SAX data which exhibits an isothermal core and then decreases quickly. Here we present results derived from recent high-resolution numerical simulations and show that there appears to be a universal temperature profile which declines significantly even within half a virial radius of the cluster center. We also show that these simulated profiles are consistent with the most recent cluster observations except in the very innermost regions of a cluster ($r<0.2 \\; r_{vir}$). ",
        "conclusions": "We have used our volume-limited catalogs of numerically-simulated clusters to demonstrate that both $\\Lambda$CDM and SCDM clusters follow a universal cluster temperature profile with  $T/T_o=1.3  [ 1 + 1.5 r/r_{vir} ] ^{-1.6}$. This temperature  profile agrees very well with the observationally-determined profile of M98 and the more recent {\\it Beppo}SAX data of DeGrandi \\& Molendi (2002) (Fig.~\\ref{beppo}). Our simulations also reveal a wealth of extremely complex and non-isothermal temperature structure (Fig.~\\ref{tempmaps}) which current X-ray telescopes may be able to probe. Interestingly, our preliminary analysis suggests that the large-scale cluster temperature structure (particularly for $z>0$) is a potential discriminator of cosmological model.   Our current adiabatic catalogs of clusters were intended to serve as a well-defined baseline sample for statistical and individual comparison with both observed clusters and with future simulations that include additional physical effects. Remarkably, our initial simulations of clusters with more realistic physics (radiative cooling and galaxy feedback) give rise to temperature profiles (Fig.~\\ref{cooling}) that agree well with our adiabatic profile and suggest that this truly is a universal cluster temperature profile.     \\noindent {\\bf Acknowledgements}  This work has been partially supported by NASA grant NAG5-7404 and NSF grant AST-9803137. G.L.B. is supported by NASA through Hubble Fellowship grant HF-01104.01-98A from the Space Telescope Science Institute, which is operated under NASA contract NAS6-26555. The simulations were carried out on the SGI/Cray Origin 2000 at the National Center for Supercomputing Application, University of Illinois Urbana-Champaign."
    },
    "0207/astro-ph0207026_arXiv.txt": {
        "abstract": "{We present a compilation of spectroscopic observations of the sgB[e] star \\object{CI~Cam}, the optical counterpart of \\object{XTE~J0421+560}.  This includes data from before, during, and after its 1998 outburst, with quantitative results spanning 37 years. The object shows a rich emission line spectrum originating from circumstellar material, rendering it difficult to determine the nature of either star involved or the cause of the outburst.  We collate all available pre-outburst data to determine the state of the system before this occurred and provide a baseline for comparison with outburst and post-outburst data.  During the outburst all lines become stronger, and hydrogen and helium lines become significantly broader and asymmetric.  After the outburst, spectral changes persist for at least three years, with Fe\\,\\textsc{ii} and [N\\,\\textsc{ii}] lines still a factor of $\\sim2$ above the pre-outburst level and He\\,\\textsc{i}, He\\,\\textsc{ii}, and N\\,\\textsc{ii} lines suppressed by a factor of 2--10.  We find that the spectral properties of \\object{CI~Cam} are similar to other sgB[e] stars and therefore suggest that the geometry of the circumstellar material is similar to that proposed for the other objects: a two component outflow, with a fast, hot, rarefied polar wind indistinguishable from that of a normal supergiant and a dense, cooler equatorial outflow with a much lower velocity.  Based on a comparison of the properties of \\object{CI~Cam} with the other sgB[e] stars we suggest that \\object{CI~Cam} is among the hotter members of the class and is viewed nearly pole-on.  The nature of the compact object and the mechanism for the outburst remain uncertain, although it is likely that the compact object is a black hole or neutron star, and that the outburst was precipitated by its passage through the equatorial material.  We suggest that this prompted a burst of supercritical accretion resulting in ejection of much of the material, which was later seen as an expanding radio remnant.  The enhanced outburst emission most likely originated either directly from this supercritical accretion, or from the interaction of the expanding remnant with the equatorial material, or from a combination of both mechanisms. ",
        "introduction": "\\label{Intro}  On 1998 April 2 Smith et al.\\ (\\cite{Smith:1998a}) reported an {\\it RXTE} All-Sky Monitor (ASM) detection of a bright, rapidly rising X-ray transient designated \\object{XTE~J0421+560}.  Subsequently, Marshall et al.\\ (\\cite{Marshall:1998a}) used the {\\it RXTE} Proportional Counting Array (PCA) to refine the best fit position with an error circle of 1\\,arcmin radius. The bright ($V\\sim11$) B[e] star \\object{CI~Cam}(=\\object{MWC~84}) was found to lie near to the centre of this error circle.  Spectroscopic observations by Wagner et al.\\ (\\cite{Wagner:1998a}) on 1998 April 3 revealed a rich emission line spectrum, similar to that reported by Downes (\\cite{Downes:1984a}; see Sect.~\\ref{QuiescentSpec}), but with the presence of He\\,\\textsc{ii} emission features. These features had not been reported in previous spectra, and so by analogy to other X-ray binaries Wagner et al.\\ (\\cite{Wagner:1998a}) proposed it to be the optical counterpart of \\object{XTE\\,J0421+560}. Photometric observations of the source at this time (e.g.\\ Robinson et al.\\ \\cite{Robinson:1998a}, Garcia et al.\\ \\cite{Garcia:1998a}, Hynes et al.\\ \\cite{Hynes:1998a}) showed that \\object{CI~Cam} was some 2--3\\,mag brighter than had previously been reported.  Hjellming \\& Mioduszewski (\\cite{Hjellming:1998a}) reported the detection of a transient 19\\,mJy radio source at 1.4\\,GHz, corresponding to the optical position of \\object{CI~Cam} on 1998 April 1, thus confirming the identification of \\object{CI~Cam} as the optical counterpart.  Subsequent observation of rapid radio variability established that the radio emission was of non-thermal (synchrotron) origin (Hjellming \\& Mioduszewski \\cite{Hjellming:1998b}).  Long term observations indicate that after the initial flare the radio emission underwent an unusually slow decay, with a 15\\,GHz flux of $\\sim1.5$\\,mJy about 40 months after the initial outburst (Pooley, priv.\\ comm.).  High spatial resolution maps obtained after the outburst indicated the presence of a clumpy ejection nebula (Mioduszewski et al., in preparation).  These ejecta expand at $\\sim 1.0-1.5$\\,mas\\,d$^{-1}$, corresponding to an expansion velocity $\\sim5000$\\,km\\,s$^{-1}$ for a distance of 5\\,kpc.  Given the distance estimates (and optical luminosity implied) for \\object{CI~Cam} (e.g.\\ $\\log L/L_{\\odot} \\geq4.86$; Clark et al.\\ \\cite{Clark:2000a}; Robinson, Ivans \\& Welsh \\cite{Robinson:2002a}) it is clear that if it is a binary, as is likely, it is a high mass X-ray binary (HMXB). However \\object{CI~Cam} does not sit comfortably within the traditional divisions of HMXB mass donors into classical Be stars ($\\sim70$\\,per cent) and OB supergiants ($\\sim30$\\,per cent).  While its luminosity suggests that it belongs to the later subset, such systems are typically short period binaries which accrete via Roche lobe overflow or direct wind fed accretion producing {\\em persistent} X-ray emission (typically modulated at the orbital period).  The presence of a rich emission line spectrum including forbidden lines, and a near IR excess due to hot dust (the observational criteria for the B[e] phenomenon; Allen \\& Swings \\cite{Allen:1976a}; Lamers et al.\\ \\cite{Lamers:1998a}) also mark a distinction from the other supergiant HMXB systems.  Among the stars showing the B[e] phenomenon, the high luminosity of \\object{CI~Cam} makes it a Galactic counterpart to the Magellanic Cloud supergiant B[e] stars (sgB[e] stars) and we will refer to it as such for the rest of this work.  \\object{CI~Cam} therefore appears to be the first bona fide sgB[e] star HMXB known, although direct evidence for binarity has proven elusive and a chance encounter of a compact object with the sgB[e] star although highly unlikely, cannot be ruled out.  This work presents a compilation of spectroscopy obtained before, during and after the 1998 outburst.  A companion photometric compilation has been presented by Clark et al.\\ (\\cite{Clark:2000a}). Some of the outburst data included here has previously been presented by Barsukova et al.\\ (\\cite{Barsukova:1998a}) and Barsukova et al.\\ (\\cite{Barsukova:2002a}).  In Sect.~\\ref{QuiescentSpec} we summarise the available pre-outburst data, including archival data with quantitative spectroscopy spanning $\\sim30$\\,years before the X-ray outburst and additional unpublished pre-outburst spectra.  From these we identify the typical pre-outburst strengths of spectral lines and discuss their stability.  We then in Sect.~\\ref{OutburstSpec} describe a series of new spectra running from a few days after the X-ray outburst to $\\sim3$\\,years later.  In Sect.~\\ref{ExtinctionSection} we discuss extinction and distance estimates for the system and in Sect.~\\ref{CompanionSection} we review what is known about the mass donor star.  Sect.~\\ref{FluxSection} examines the changes in the continuum flux distribution and Sect.~\\ref{SpecEvol} the spectral lines and how these evolve through the outburst.  Sect.~\\ref{Rapid} tests for the presence of shorter timescale variability.  Finally in Sect.~\\ref{Discussion} we will discuss how all of these clues can help us build a picture of the nature of the system and the outburst mechanism and in Sect.~\\ref{conc} we summarise our conclusions. ",
        "conclusions": "\\label{conc} \\object{CI~Cam} is an sgB[e] star, and many of its characteristics, particularly the emission line spectrum, are representative of this class.  Comparison with other sgB[e] stars suggests that \\object{CI~Cam} is among the hottest members of the class and viewed nearly pole-on.  It differs from the other sgB[e] stars in interacting with a compact star which introduces an additional variable element. It seems most likely that the compact object is physically associated with \\object{CI~Cam}, i.e.\\ that it is an HMXB.  In this case, the compact object is probably in a long period, eccentric orbit.  However we cannot rule out the possibility that this was a chance encounter and that the two stars are not physically associated.  Resolution of this issue will likely involve waiting for a true periodicity, repeated over several cycles, or another outburst.  In considering the outburst mechanism, there is actually little difference between a long period eccentric orbit and a chance encounter, so our discussion of the outburst applies to both cases.  We suggest that the majority of the optical emission lines originate from an equatorially concentrated outflow or circumstellar disc. During outburst, hydrogen and helium lines appear to have two components, a narrow rest component and a moderately blueshifted broad component.  Metallic lines are mainly dominated by the narrow component at all times, although some asymmetry is seen in outburst suggesting that a broad component is present.  The square profiles of the Fe\\,\\textsc{ii} lines can be explained by the equatorial outflow model, if viewed pole-on, so the narrow component is likely associated with this.  The broad component becomes weaker and narrower on the decline, and almost disappears in quiescence.  This may come from material ejected in the outburst rather than from the equatorial outflow.  Forbidden lines fall into two categories; [O\\,\\textsc{i}] and [Fe\\,\\textsc{ii}] show similar line profiles to Fe\\,\\textsc{ii}, but with a lack of low velocity material.  These profiles could originate from the low density, upper layers of the equatorial outflow.  Other forbidden lines, [N\\,\\textsc{ii}] and [O\\,\\textsc{iii}] are narrower and likely come from a much more extended region.  The outburst mechanism remains undetermined, although the outburst was probably precipitated by the passage of the compact object through the equatorial material.  It is unlikely that X-ray heating of any component is responsible for the optical outburst.  Instead the optical outburst is likely associated with the expanding remnant produced by the X-ray outburst, either through direct emission from the remnant or as a result of its interaction with the circumstellar material.  The spectral shape of the outburst optical continuum, and the presence of broad, blue-shifted emission components, are both consistent with predictions for supercritical accretion resulting in ejection of much of the material (Shakura \\& Sunyaev \\cite{Shakura:1973a}), and the peak mass transfer rate for an equatorial passage of the compact object is indeed predicted to be well above the Eddington limit.  After the outburst changes in the emission lines persist for at least three years, with Fe\\,\\textsc{ii} lines stronger than before and He\\,\\textsc{i}, He\\,\\textsc{ii}, and N\\,\\textsc{ii} lines weaker.  The timescale for the extended Fe\\,\\textsc{ii} decay, at least, is similar to the expected viscous timescale of the disc of hundreds to thousands of days, so this may indicate the gradual recovery of the disc to its equilibrium state. As the system does not yet appear to have stabilised continued monitoring is important to determine if the system eventually recovers to the pre-outburst state or if it settles to a different level."
    },
    "0207/astro-ph0207483_arXiv.txt": {
        "abstract": "Number counts of galaxies are re-analyzed using a semi-analytic model (SAM) of galaxy formation based on the hierarchical clustering scenario. Faint galaxies in the Subaru Deep Field (SDF, near-infrared $J$ and $K'$) and the Hubble Deep Field (HDF, ultraviolet/optical $U$, $B$, $V$, and $I$) are compared with our model galaxies.  We have determined the astrophysical parameters in the SAM that reproduce observations of nearby galaxies, and used them to predict the number counts and redshifts of faint galaxies for three cosmological models, the standard cold dark matter (CDM) universe, a low-density flat universe with nonzero cosmological constant, and a low-density open universe with zero cosmological constant.  The novelty of our SAM analysis is the inclusion of selection effects arising from the cosmological dimming of surface brightness of high-redshift galaxies, and from the absorption of visible light by internal dust and intergalactic \\ion{H}{1} clouds.  As was found in our previous work, in which the ultraviolet/optical HDF galaxies were compared with our model galaxies, we find that our SAM reproduces counts of near-infrared SDF galaxies in a low-density universe either with or without a cosmological constant, and that the standard CDM universe is {\\it not} preferred, as suggested by other recent studies.  Moreover, we find that simple prescriptions for (1) the timescale of star formation being proportional to the dynamical time scale of the formation of galactic disks, (2) the size of galactic disks being rotationally supported with the same specific angular momentum as that of surrounding dark halo, and (3) the dust optical depth being proportional to the metallicity of cold gas, cannot completely explain all of observed data.  Improved prescriptions incorporating mild redshift-dependence for those are suggested from our SAM analysis. ",
        "introduction": "It is well known that the number of faint galaxies in a given area of sky can constrain cosmological parameters, because it depends on the geometry of the Universe (e.g., Peebles 1993).  Many efforts have been devoted to this subject using traditional galaxy evolution models assuming monolithic collapse, such as the wind model for elliptical galaxies and the infall model for spiral galaxies (e.g., Yoshii \\& Takahara 1988).  These models are in fact able to reproduce many of the observed properties of nearby galaxies, and provide a useful theoretical tool for understanding their evolution \\citep{ay86, ay87, ayt91}.  In the analyses of galaxy counts from traditional approaches, it has been found that the Einstein-de Sitter (EdS) universe, a representation of the standard cold dark matter (CDM) universe, is not reconcilable with the observed high counts of faint galaxies, and that a low-density universe is preferred \\citep{yt88, yp91, yp95, y93}.  Recently, \\citet{ty00} and \\citet{t01} compared their predictions against the observed number counts to the faint limits in the Hubble Deep Field (HDF; Williams et al. 1996) and in the Subaru Deep Field (SDF; Maihara et al. 2001), taking into account various selection effects and allowing for the possibility of number evolution of galaxies in a phenomenological way.  Note that the SDF counts are now the deepest near-infrared ones with the 5$\\sigma$ limiting magnitude of $K=23.5$ in total magnitude.  They confirmed that the EdS universe cannot reproduce the observed high counts.  However, in their best-fit models, the merger rates of HDF and SDF galaxies are a little different.  A mild merger rate is needed to reproduce the counts in the HDF, while a negligible rate was necessary for the SDF.  It should be noted that the photometric passbands for the two applications are different; ultraviolet/optical for the HDF and near-infrared for the SDF.  They suggested that the difference of the merger rate might be originated by morphology-dependent number evolution because late-type galaxies are mainly seen in shorter wavelength such as $B$-band and early-type galaxies are seen in longer wavelength such as $K$-band.  In any case, it should be explained why the merger rate depends on the observed wavelength in order to obtain a better understanding of the galaxy formation process.  On the other hand, in the study of formation of large-scale structure in the universe, both theory and observation suggest that gravitationally bound objects, such as galaxy clusters, are formed through continuous mergers of dark halos with an initial density fluctuation spectrum predicted by the CDM models.  Based on this scenario of {\\it hierarchical clustering}, the so-called semi-analytic models (SAMs) of galaxy formation have been developed by a number of authors (Kauffmann, White \\& Guiderdoni 1993; Cole et al. 1994, 2000; Somerville \\& Primack 1999; Nagashima et al. 2001, hereafter NTGY).  SAMs successfully reproduced a variety of observed features of local galaxies, such as their luminosity function, color distribution, and so on.  Faint galaxy number counts have also been analyzed using SAMs \\citep{c94, kgw94, h95, bcf96}.  These studies showed that predicted number counts in the EdS universe agree with the observed counts.  Their results, however, contradict analyses carried out with traditional galaxy evolution models.  Recently, this contradiction was resolved by NTGY, in which their SAM is compared with the galaxy counts in the HDF. They found, by matching properties of model galaxies with observation especially in local luminosity functions and cold gas mass fraction, that normalization of model parameters, related to combinations of physical processes such as star formation (SF) and supernova (SN) feedback, is very important, and that accounting for selection effects caused by cosmological dimming of surface brightness and absorption of emitted light by internal dust and intergalactic \\ion{H}{1} clouds is crucial in the analysis of galaxy counts, as shown by \\citet{ty00}.  It should be noted that recent analysis by \\citet{lypcf} also clarify the importance of the selection effects caused by the cosmological dimming of surface brightness in the observational point of view.  They introduced the star formation rate intensity distribution function, which was derived from the ultraviolet luminosity density for the HDF galaxies, at several redshifts and found that at high redshift significant fraction of ultraviolet luminosity density must be missed due to the cosmological dimming of surface brightness.  The purpose of this paper is to examine whether our SAM can {\\it simultaneously} reproduce both the ultraviolet/optical and near-infrared galaxy counts in the HDF and in the SDF.  Because the luminosity of galaxies in different passbands reflects the influence of different stellar populations, the subject of multi-band number counts provides a strong constraint on galaxy formation.  In this paper, using selection criteria for SDF galaxies based on \\citet{t01}, we compare our SAM prediction of galaxy counts with the observed counts in the SDF.  This paper is outlined as follows. In \\S2 we briefly describe our SAM, which is almost the same as our previous models (NTGY). In \\S3 we constrain the astrophysical parameters in our SAM analysis using local observations. In \\S4 we compare theoretical number counts of faint galaxies with the HDF and SDF data, and in \\S5 we discuss the range of uncertainties in our calculations of galaxy number counts.  In \\S6 we provide a summary and discussion. ",
        "conclusions": "\\label{sec:summary}  We have calculated the number counts of faint galaxies in the framework of SAM for three cosmological models, the standard CDM (EdS) universe, a low-density open universe, and a low-density flat universe with nonzero $\\Lambda$.  Our SAM includes the selection effects from the cosmological dimming of surface brightness of galaxies with criteria appropriate for the SDF and HDF, and also includes some modifications to our previous analysis (NTGY), such as the optical depth estimation of dust within a galaxy. In this paper we have shown that our SAM is fully consistent with that of the previous version, and can explain the observed multi-band galaxy counts from the UV to the near-infrared.  Comparison of theoretical predictions with the observed number counts of SDF and HDF galaxies, as well as with other ground-based observations, indicates that the standard CDM is ruled out, and a $\\Lambda$-dominated flat universe and a low-density open universe are favored. This result is consistent with that from HDF galaxies (NTGY), but is in sharp contrast with previous SAM analyses by other authors, where many of the conceivable selection effects in the faint observations have been ignored.  Some uncertainties in our SAM have been discussed.  These arise from a lack of knowledge on the galaxy formation process, and also from an insufficient survey of the physical properties of high-redshift galaxies.  We especially focused on the uncertainties in redshift dependence of SF timescale, galaxy size, and dust extinction.  We found that dust extinction hardly affects galaxy counts in the Subaru $K'$ band, but does significantly affect those in the {\\it HST} $B_{450}$ band.  Thus, the $K'$-band galaxy counts are robust against the uncertainty of dust extinction.  Two other factors affect galaxy counts even in the $K'$-band.  If the SF timescale at high redshift is shorter than one which is simply proportional to the dynamical timescale in the disk, too many galaxies are formed, and the number of galaxies at faint-end is greatly overpredicted.  We found that the SF timescale should be nearly constant against redshift, as suggested by our previous analysis (NTGY) and by other recent SAM analyses\\citep{kh00, spf01}.  The uncertainty in estimation of galaxy size results in an uncertainty in estimation of the surface brightnesses of galaxies, which is directly related to the selection effects mentioned above.  In our SAM, like usual SAMs, the size of the disk is determined under the assumption of specific angular momentum conservation of cooling gas, so that the disk size is proportional to the cooling radius.  In order to see how theoretical predictions are changed by changing the galaxy size, we introduced a free parameter $\\rho$ allowing for an additional redshift dependence in size estimation.  We found that the value of $\\rho\\ga 0.5$ is favored in order to reproduce the observed counts, which indicates that the disk radius should be extended by a factor of $(1+z)$ over the cooling radius.  We also found that this manipulation cannot be discriminated observationally because we cannot know how many undetected low surface brightness galaxies below the selection criteria are there at high redshift, which are presumably systems of large size.  Through this work, we have shown that our SAM can explain a variety of observed properties of nearby and high-redshift galaxies, and place some constraints on star formation, size evolution, and dust extinction. More stringent constraints will certainly be obtained by a greater knowledge of dynamical and kinematical properties of galaxies in near future."
    },
    "0207/astro-ph0207160_arXiv.txt": {
        "abstract": "We consider radiation of relativistic electrons accelerated within the jet extended boundary layer. Due to velocity shear across the boundary the observed jet/counterjet brightness ratio is diminished as compared to the one derived for the jet spine. Thus the jet Lorentz factor evaluated from the observed jet asymmetry can be an underestimated value influenced by observation of the slower boundary layer. We briefly discuss several consequences of the radiating boundary layer model in the context of recent Chandra and XMM observations of the large scale jets. \\textbf{} ",
        "introduction": "Jet stratification was proposed in order to interpret radio and optical observations of large scale jets in radio galaxies (e.g., \\cite{lai96}). The inferred jet morphology consists of a fast central {\\it spine} surrounded by the {\\it shear layer} extending into the broad jet {\\it cocoon}. However, the physical properties of the boundary region are not exactly known. Polarimetry of the large scale jet boundaries usually show magnetic field being parallel to the jet axis, suggesting strong shearing effects at the jet edges. 3D numerical simulations reveal turbulent character of the boundary layer and its high specific internal energy (\\cite{alo99}). Such regions are therefore promising places for particle acceleration, including both stochastic scattering in a turbulent medium and cosmic ray viscosity (cf. review by \\cite{ber90}). With the parameters characteristic for the large scale jets, the former mechanism creates at the jet boundary a characteristic two-component electron spectrum: a power-law ended by a pile-up bump (\\cite{ost00}, see also \\cite{sta01}). A role of such electron distribution for the multiwavelength large scale jet emission was studied by \\cite{sta02}.  Recently, Chandra and XMM observatories detected significant X-ray nonthermal emission connected with the large scale jets in a number of radio loud AGNs. One of the most spectacular X-ray jets is the one observed in quasar PKS 0637-752 (\\cite{cha00}). The most likely explanation of its X-ray emission is the inverse-Compton scattering of cosmic microwave background photons by the low energy nonthermal electrons (\\cite{tav00}). This mechanism, however, requires highly relativistic flow velocities at the distance of order of 0.1 - 1 Mpc from the galactic nucleus. It is inconsistent with middly relativistic velocities inferred from the jet-counterjet radio brightness asymmetries of the large scale jets in radio galaxies. One may note, that the highly relativistic jets provide the most efficient way of energy transport to the terminal hot-spots in FR II sources (\\cite{ghi01}). Therefore, there is striking disagreement between theoretical modeling and radio measurements for such sources.  Here we discuss the possibility, that the emission of electrons accelerated within the turbulent shear layer affects the jet-counterjet brightness asymmetry observations, allowing the jet spine to remain highly relativistic even far away from the active nucleus. We point out several consequences of the presented model for radiative output of large scale jets, including their X-ray emission. ",
        "conclusions": "Relativistic velocities of large scale jets (corresponding to the bulk Lorentz factor $\\Gamma_j \\sim 10$) were postulated in order to explain X-ray emission detected by Chandra and XMM from the jets with small inclination angles. Much smaller velocities of the jets in radio galaxies inferred from the jet-counterjet radio brightness asymmetry can be explained as a result of the radiating boundary layer with the velocity shear. The jet spines flowing with large bulk Lorentz factors can carry high kinetic energy, $L_{tot} \\sim 10^{47} {\\rm erg/s}$, assuming that they are composed from `normal' electron-proton plasma (cf. \\cite{ghi01}), and efficiently produce X-rays due to Compton scattering of CMB photons. The radiating boundary layer can dominate the X-ray emission of the jets observed at large viewing angles, as discussed by \\cite{cel01} and \\cite{sta02}. In the later model, such X-ray emission is produced by synchrotron radiation of the high energy electron pile-up bump. Therefore, the jet-counterjet X-ray brightness asymmetry differ from the radio one due to different spectral characters of the power-law and the spectral bump emission. One should note, that recent HST observations of the jet in 3C 273 (\\cite{jes02}) show spectral flattening starting at UV frequencies and suggest that our interpretation of the large scale jets' X-ray emission can be correct at least for some sources."
    },
    "0207/astro-ph0207356_arXiv.txt": {
        "abstract": "We apply a new algorithm, called the Unbiased Minimal Variance (hereafter UMV) estimator, to reconstruct the cosmic density and peculiar velocity fields in our local universe from the SEcat catalog of peculiar velocities comprising both early (ENEAR) and late type (SFI) galaxies. The reconstructed fields are compared with those predicted from the IRAS PSC$z$ galaxy redshift survey to constrain the value of $\\beta=\\Omega_m^{0.6}/b$, where $\\Omega_m$ and $b$ are the mass density and the bias parameters. The comparison of the density and velocity fields is carried out within the same methodological framework, and leads, for the first time, to consistent values of $\\beta$, yielding $\\beta=0.57_{-0.13}^{+0.11}$ and $\\beta=0.51 \\pm {0.06}$, respectively.  We find that the distribution of the density and velocity residuals, relative to their respective errors, is consistent with a Gaussian distribution with $\\sigma\\approx 1$, indicating that the density field predicted from the PSC$z$ is an acceptable fit to that deduced from the peculiar velocities of the SEcat galaxies. ",
        "introduction": "\\label{sec:intro}  In the gravitational instability scenario (\\eg, Peebles 1980), mass density fluctuations and peculiar velocities evolve in an expanding universe under the effect of gravity. If density fluctuations are small, linear theory is valid and a simple relation exists between peculiar velocities, $\\bv$, and mass density contrast, $\\delta_m$: \\begin{equation} \\nabla \\cdot \\bv = -\\Omega_m^{0.6} \\delta_m, \\label{eq:divv} \\end{equation} where $\\Omega_m$ is the mass density parameter. Equation~(\\ref{eq:divv}) shows why peculiar motions are so important in cosmology: they provide a direct probe of the mass density distribution in the universe. The mass density fluctuation field, $\\delta_m$, can be deduced from the galaxy observed density contrast, $\\delta_g$, assuming a relation (bias) between the distribution of galaxies and that of the underlying density. The simplest relation suggested in the literature is that of linear bias, namely $\\delta_g = b \\delta_m$, where $b$ is the linear bias parameter for a given population of mass tracers. This assumption seems to hold on very large (linear) scales and it is supported by both observational evidence (e.g. Baker \\etal, 1998 and Seaborne \\etal, 1999) and numerical studies (\\eg, Blanton \\etal, 1999).  Many authors have used galaxies' peculiar velocities and their redshift space positions to estimate the value of $\\beta=\\Omega_m^{0.6}/b$, under the hypotheses of linear theory and linear biasing. These analyses have been typically carried out using two alternative strategies. In the so-called density-density comparisons a 3-D velocity field and a self-consistent mass density field are derived from observed radial velocities and compared to the galaxy density field measured from large redshift surveys.  The typical example is the comparison of the mass density field reconstructed with the POTENT method (Bertschinger \\& Dekel 1989, Dekel \\etal, 1990) from the MARK III catalog of galaxy peculiar velocities (Willick \\etal, 1997a) with the galaxy density field obtained from the IRAS 1.2 Jy redshift catalog (Sigad \\etal, 1998). The various applications of density-density comparisons to a number of datasets have persistently led to large estimates of $\\beta$, consistent with unity (see Sigad \\etal, 1998 and references therein). The alternative approach is constituted by the velocity-velocity analyses. In this second case the observed galaxy distribution is used to infer a mass density field from which peculiar velocities are obtained and compared to the observed ones. The velocity-velocity methods have been applied to most of the velocity catalogs presently available yielding systematically lower values of $\\beta$, in the range $0.4 - 0.6$ (see Zaroubi 2002a, for a summary of the most recent results).  Both density-density and velocity-velocity methods have been carefully tested using mock catalogs extracted from N-body simulations. They were shown to provide an unbiased estimate of the $\\beta$ parameter. Yet, when applied to the same datasets, the discrepancy in the $\\beta$ estimates turned out to be significantly larger than the expected errors. Accounting for mildly nonlinear motions (e.g. Sigad \\etal, 1998 and Willick \\etal, 1996) or allowing for possible deviations from a pure linear biasing relation consistent with the observational constraints (see discussion in Somerville \\etal, 2001, Branchini \\etal 2001) does not explain this discrepancy (Berlind, Narayanan and Weinberg 2001). Velocity-velocity comparisons are generally regarded as more reliable as they require manipulation of the denser and more homogeneous, redshift catalog data. Whereas, the density-density comparisons involve manipulation of the noisier and sparser velocity data.  In any case both classes of methods are quite complicated and it is hard to understand how systematic errors can arise and propagate through the analysis. Therefore, it is likely that these systematics affect the $\\beta$ parameter estimation.  The purpose of this work is to address, and possibly solve, the density-density {\\it vs.} velocity-velocity dichotomy. We achieve this goal by using the novel Unbiased Minimal Variance estimator, recently proposed by Zaroubi (2002b). The UMV estimator allows one to reconstruct an unbiased cosmological field at any point in space from sparse, noisy and incomplete data and to map it into a dynamically-related cosmic field ($e.g.$ to go from peculiar velocities to overdensities). The UMV is applied here to the SEcat catalog of peculiar velocities (Zaroubi 2002b) to reconstruct both the mass density and peculiar velocity fields. These fields are then compared with the analogous quantities predicted from the distribution of IRAS PSC$z$ galaxies (Saunders \\etal, 2000) of density-density and a velocity-velocity analyses. In Section~\\ref{sec:method} we briefly review the basics of the UMV estimator. The SEcat and PSC$z$ catalogs are presented in Section~\\ref{sec:data}. Error estimation from mock catalogs is described in \\S~\\ref{sec:mock}. The density and velocity fields obtained by applying UMV to SEcat are compared in Section~\\ref{sec:results} with the analogous quantities deduced from the PSC$z$ dataset. Finally, in Sections~\\ref{sec:discussion} we discuss the results and present our conclusions. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have applied the new UMV estimator to recover the density and velocity fields in the local universe from the SEcat catalog of galaxy peculiar velocities. In order to obtain the so called $\\beta$ parameter, these fields have been compared with those modeled from the spatial distribution of IRAS PSC$z$ galaxies assuming linear theory and biasing. Previous estimates of $\\beta$ from density-density comparisons, mainly based on the POTENT algorithm (Bertschinger \\& Dekel 1989, Dekel \\etal, 1990), have yielded a large value ($\\beta\\approx 1$ \\cf, Sigad \\etal, 1998), inconsistent with the smaller values ($\\approx 0.5$) independently obtained from all recent velocity-velocity VELMOD (Willick \\etal, 1996, Willick \\& Strauss 1998, and Branchini \\etal, 2001) and ITF (da Costa \\etal, 1998 and Nusser \\etal, 2000) comparisons.  For the first time the UMV method provides a common methodological framework in which to perform velocity-velocity and density-density comparisons.  The velocity-velocity comparison yields a value of $\\beta$ consistent with that measured in the VELMOD and ITF analyses. However, the value of the same parameter obtained from our density-density comparison is significantly smaller than those obtained from the POTENT analyses (\\cf, Sigad \\etal, 1998).  The $\\beta$ parameters from both $v-v$ and $\\delta-\\delta$ comparisons presented here are in agreement, yielding a $\\beta \\approx 0.55$ with an estimated error of the order of $0.1$.  In contrast with the POTENT algorithm, the new UMV method reconstructs the density field from peculiar velocities while taking into account their underlying correlation properties. The regularization aspect of the UMV estimator significantly improves the stability of the inversion, which is especially important given the low signal-to-noise ratio of peculiar velocity data.  The regularization obtained by this method is very similar to the one provided by the Wiener filter method (Zaroubi \\etal, 1999). However, the UMV has been designed to provide an unbiased estimator of the underlying signal, a property that the Wiener filtering method lacks. These two aspects make the UMV estimator a very useful tool for reconstruction from peculiar velocity data.  In our error analysis we have shown that for the best fit value of $\\beta$ the density and velocity residuals are normally distributed. This indicates that the PSC$z$ density and velocity fields constitute an adequate model for those reconstructed with the UMV estimator. The fields only differ by a monopole term, corresponding to a mismatch in the the mean density within $60 \\hmpc$ which is caused by the known incompleteness of the PSC$z$ catalog at faint fluxes (Teodoro \\etal, 2000).  This also implies that the effect of the nonlinear dynamics and amount of nonlinear and stochastic biasing on the scales involved in our analysis is negligible relative to the measured peculiar velocity errors.  The results presented in this paper are quite encouraging since for the first time the two ways of estimating the value of the $\\beta$ parameter give a consistent result. Our results also suggest that the UMV estimator is a promising tool for the problem of reconstructing the dynamical fields from the observed radial peculiar velocities and therefore could be applied to other datasets. In particular, to reconstruct the large scale structure from the incoming large and uniform surveys that will provide both the spatial distribution and peculiar velocities of extra-galactic objects \\eg, SDSS and large cluster surveys with kinematic Sunyaev-Zel'dovich measurements.  Our present density-density comparison results are in marked contrast to those obtained by earlier work, including those from the recent analysis of the Mark~III catalog using the POTENT method (\\eg, Sigad \\etal, 1998), and raises the question on the origin of this discrepancy. The Mark III catalog, as shown for example by Davis \\etal\\ (1996) and more recently by Courteau \\etal\\ (2000), suffers from systematic calibration errors that would cause a systematic error in the estimation of $\\beta$.  However, these errors are not expected to overestimate the value of $\\beta$ by more than a factor of two.  An application of the UMV method to the Mark III catalog shows that the obtained values of $\\beta$ are somewhat higher than those obtained from the SEcat catalogs by $0.1-0.2$. Moreover, the $v-v$-like VELMOD analysis yield consistent values of $\\beta$ when applied to Mark III and SFI datasets (Willick \\etal, 1997b, Willick and Strauss 1998, Branchini \\etal, 2001).  Based on these arguments, one could speculate that the most likely explanation to the inconsistent results is a conspiracy of both the systematics errors in Mark III and some noise-driven inversion instability in the POTENT reconstructions."
    },
    "0207/astro-ph0207411_arXiv.txt": {
        "abstract": "The origin of the diffuse extragalactic, high-energy gamma-ray background (EGRB) filling the Universe remains unknown. The spectrum of this extragalactic radiation, as measured by the {\\sl Energetic Gamma Ray Experiment Telescope} (EGRET) on-board the Compton Gamma-Ray Observatory (CGRO), is well-fit by a power law across nearly four decades in energy, from 30 MeV to 100 GeV. It has been estimated that not more than a quarter of the diffuse gamma-ray background could be due to unresolved point sources. Recent studies have suggested that much of the diffuse background could originate from the up-scatter of cosmic microwave background (CMB) photons by relativistic electrons produced by shock waves in the intergalactic medium (IGM) during large-scale structure formation.  In this work we search for evidence of gamma-ray emission associated with galaxy clusters by cross-correlating high Galactic latitude EGRET data with Abell clusters. Our results indicate a possible association of emission with clusters at a $\\geq 3\\sigma$ level. For a subset of the 447 richest ($R\\geq 2$) clusters the mean surface brightness excess is $1.2\\times 10^{-6}$ ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ ($>100$MeV), corresponding to a typical non-thermal bolometric luminosity of $L_{\\gamma}\\sim 1\\times 10^{44}$ erg s$^{-1}$. Extrapolating this measurement and assuming no evolution we conservatively estimate that $\\sim 1-10$\\% of the EGRB could originate from clusters with $z<1$. For this cluster population the predicted non-thermal luminosity is in excellent agreement with our measurement, suggesting that the clusters have experienced mass accretion within the last $10^9$ yrs. If correct, then future gamma-ray missions, such as the Gamma-ray Large Area Space Telescope (GLAST) should be able to directly detect nearby galaxy clusters. ",
        "introduction": "During the nine years (1991 - 2000) EGRET was operational on-board the Compton Gamma-Ray Observatory, it detected a diffuse gamma-ray emission above 30 MeV filling the Universe. This gamma-ray background is composed of an intense Galactic component, due to cosmic ray interactions with local interstellar gas and radiation \\citep{hun97}, as well as a diffuse and isotropic extragalactic component. The extragalactic radiation is well described by a simple power law with index $-2.1\\pm 0.3$ from 30 MeV to 100 GeV \\citep{sre98}.  The nature of this extragalactic gamma-ray background (EGRB) has long been debated.  It is not clear if the EGRB has a purely diffuse origin or is made up of the superposition of discrete, unresolved, point sources. The majority of the identified EGRET sources are active galaxies in the blazar class, and this has led to the suggestion that the EGRB is produced by an unresolved population of blazars \\citep{pad93}. However, estimates of contributions from unresolved gamma-ray blazars range from nearly all ~\\citep{ste96,ste01} to less than about a quarter \\citep{chi98,muc00}. It seems likely that other sources of the diffuse extragalactic radiation must exist.  Recent work by ~\\citet{loe00} has suggested that some, if not most of the diffuse gamma-ray background could be a result of large-scale structure formation in the Universe. Gravitationally-induced shock waves produced during cluster mergers and large-scale structure formation give rise to highly relativistic electrons which are responsible for Compton up-scattering the cosmic microwave background photons to high energy gamma-rays. The fraction of the shocks' thermal energy transferred to relativistic electrons could range from 1-10\\% ~\\citep{wax00}. The gamma-ray background is then expected to be produced in filaments, sheets, and extended ($> 1^\\circ$) gamma-ray halos associated with newly formed massive clusters ~\\citep{wax00}. In some scenarios, the shock radii for clusters could be large, with gamma-rays detectable in the form of 5-10 Mpc-diameter ring-like emission tracing the cluster virialization shock \\citep{kes02}. Other work has suggested that some $0.5-2$\\% of the EGRB could arise solely from the population of clusters ~\\citep{col98}. Recent simulation studies of a $\\Lambda$CDM universe suggest that future high energy gamma-ray experiments such as GLAST or the next generation of atmospheric Cherenkov telescopes should be able to resolve gamma-rays from individual clusters in the local Universe ($z\\simeq 0.025$, \\citet{kes02}). The association of the EGRB with the large scale structure of the Universe can then be tested by cross-correlating gamma-ray maps with known galaxy clusters.  So far, individual galaxy clusters have not been found to be correlated with discrete gamma-ray sources detected by EGRET, although it has been suggested that they may be candidates for unidentified EGRET sources \\citep{tot00}. In most cases the predicted gamma-ray fluxes from nearby individual clusters \\citep{dar95,col98} are below the sensitivity limit of EGRET, and indeed, no excess gamma-ray emission is seen from any individual cluster. For example, a $2\\sigma$ upper limit of $ 4\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ at energies greater than 100 MeV can be derived for Coma, the closest rich cluster \\citep{sre96}. However, in some cases the predicted emissivity from individual clusters is in disagreement with the $2\\sigma$ EGRET upper limits \\citep{ens97}.  Previous efforts to constrain the gamma-ray emission of clusters used a small sample (58) of X-ray luminous clusters \\citep{rei99}, and found no significant signal when the EGRET data centered on the clusters were co-added. Corrections for Galactic diffuse emission were not made, and a likelihood analysis of the data found no statistically significant emission associated with the clusters. \\citet{col02} has recently reported preliminary evidence of the association of unidentified EGRET sources with galaxy clusters at high latitudes, claiming that these sources also show strong radio emission. Relativistic particles responsible for the radio emission are probably also the source of inverse-Compton gamma-rays. However, the number of claimed associations is small, and no account was taken of the known deviation from Poissonian statistics due to cluster-cluster correlations. As we discuss below, even at high Galactic latitudes - and even with the removal of the best model of Galactic emission - residual Galactic signatures remain, which is a known feature of these models \\citep{hun97}. Furthermore, by utilizing the much larger (albeit more inhomogeneous) optical Abell catalog of clusters and evaluating emission in radial bins, we increase our search sensitivity by a factor of 100-1000.   In the present analysis we search for angular correlations of the EGRET extragalactic diffuse emission with cluster position, using the complete Abell catalog of clusters (\\S2), and radially binning the gamma-ray emission around the Abell clusters (\\S3,4). We then estimate the average gamma-ray luminosity per cluster and address the question of whether this is consistent with recent predictions of the gamma-ray energy flux arising from intergalactic shocks, based on semi-analytic predictions and hydrodynamical cosmological simulations (\\S5,6). ",
        "conclusions": "We have detected evidence of a  positive correlation of unresolved gamma-ray emission with clusters of galaxies in the local Universe. The correlation signal is broadly consistent with emission localized within $\\sim 1^{\\circ}$ of clusters, and specifically is more strongly correlated with optically rich clusters.  There is some evidence that the gamma-ray emission may be more spatially extended up to $\\sim 3^{\\circ}$, or up to $\\sim 15$Mpc at the mean redshift of rich clusters in the Abell catalog. We suggest that the mean non-thermal luminosity associated with rich clusters is $\\bar{L_{\\gamma c}}\\sim 1\\times 10^{44}$ erg s $^{-1}$ (30eV-1TeV) in the local Universe.  Recent theoretical predictions and modeling have suggested that 10-80\\% of the EGRB arises from upscattering of CMB photons by relativistic electrons in the large-scale gravitational shocks of the IGM, including clusters. If such models are valid our results suggest that $\\sim 1-10$\\% of this component of the EGRB originates from structures associated with rich clusters with $0<z<1$. Specific predictions for the population of clusters alone have estimated a contribution to the EGRB of $0.5-2$\\%, dominated by systems with $z\\leq0.2$ \\citep{col98}.  Several caveats are worth noting. First, diffuse Galactic gamma-ray emission is very hard to remove, even with the current best model for its distribution. We believe we have accounted for the residual effects of this emission by assessing significance via our Monte Carlo simulations, but it may still introduce systematic errors to our results. Second, we cannot, with these data, prove that the excess emission associated with clusters is not due to some population of discrete sources, themselves positively correlated with clusters, although there is no a priori reason to believe this to be the case. Third, estimating the cluster luminosity by measuring the intensity in the innermost 1$^{\\circ}$ bin is likely to be an underestimate, given both the width of the EGRET PSF and the evidence for more extended, correlated, emission. Finally, several potential sources of random error exist serving to increase the uncertainty in quoted numbers - which, therefore, should be considered as broad estimates.  It should also be noted that another possible mechanism for producing non-thermal emission in clusters is synchrotron emission in radio galaxies contained in the clusters, that produce relativistic electrons which then up-scatter to produce the high energy gamma-rays (e.g. \\citet{gio93,pet01}). In this case the particles are accelerated in the nucleii of radio galaxies, and this could account for the observed gamma-ray luminosities. In this case, one would expect to see a stronger correlation signal with rich clusters containing radio galaxies. \\citet{led95} find that some 79\\% of $R=2$ Abell clusters have evidence of radio emitting member galaxies (in a statistical subsample of 393 clusters of all richnesses).  However, more sophisticated analysis is required owing to the small number of radio surveyed clusters and we will investigate this in an extension to the present work.   As discussed by \\citet{col98}, \\citet{loe00}, \\citet{wax00}, and \\citet{kes02}, the appearance of non-thermal emission in IGM/cluster structures is transient, with a timescale of $\\sim 10^{9}$yrs. Therefore the emission disappears almost entirely as soon as there is no strong shocking of gas. If we consider our sample of rich Abell clusters then we know from other wavelengths that the majority of these low-$z$ systems are probably in a state closely approximating hydrostatic equilibrium (e.g., the presence of X-ray cooling flows, the nearly isothermal gas, etc.). We can consider two extreme states: either only some, non-equilibrium clusters are gamma-ray bright, or all clusters are close to equilibrium and we are simply seeing the fading emission of an earlier epoch of accretion activity. In the first case we can exploit the fact that a system with an EGRET flux of at least $\\sim 5 \\times 10^{-8}$ ph s$^{-1}$ cm$^{-2}$ ($>100$MeV) would be directly detectable in the EGRET map. Given our mean flux of $\\sim 1.1\\times 10^{-9}$ (\\S5) and a total of 447 rich clusters, we would obtain an equivalent detection if only $\\sim 10$ rich clusters had a flux of $5 \\times 10^{-8}$ and the rest were gamma-ray dark. This would then provide a lower limit to the number of actively accreting clusters in the local Universe, a 2\\% fraction by number.  Alternatively, if all clusters are assumed to be close to equilibrium with a mean temperature of $kT\\simeq 3$keV, a mean gas mass $\\simeq 10^{13}$M$_{\\odot}$, and, $\\xi_e=0.05$ then they are almost exactly as gamma-ray luminous as predicted by Equation 2. This implies that in fact all clusters should have been actively accreting recently, certainly within $z<0.3-0.4$, and possibly at the present time ($z=0$).  Even in the low-density simulation of \\citet{kes02} it appears to be the case that $z=0$ clusters can have active gamma-ray emission. Semi-analytic predictions for low density cosmologies ($\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$) also suggest that even at $z\\sim 0$, for massive clusters ($10^{15}$ M$_{\\odot}$) several $10^{13}$M$_{\\odot}$ in baryons should accrete per $10^{9}$yr \\citep{lac93}. The shock regions may form some 5-10Mpc from the cluster core, creating a gamma-ray `ring' of emission \\citep{kes02}. The suggestion of a rather more extended emission pattern from our cross-correlation (Figure 9) supports this scenario. In this case our estimated $\\bar{L}_{\\gamma c}$ can be considered a good measure of the typical active non-thermal emission for rich clusters.  This would then imply that the efficiency of transfer of energy from the shocks to relativistic electrons is similar to the value of $5$\\% inferred from non-relativistic shocks in the ISM.  The simulations of \\citet{kes02} predict that for $\\xi_e\\geq 0.03$, future high resolution gamma-ray telescopes with threshold sensitivities $>10^{-10}$ for energies $> 10$ GeV should be able to resolve some gamma-ray halos associated with large scale structures. The prospect of direct detection of gamma-ray sources with emission attributed to intergalactic shocks with GLAST, VERITAS, HESS, MAGIC, or other atmospheric Cherenkov telescopes is an exciting one.  Such gamma-ray halos could be a new source class of high energy sources waiting to be discovered. Their existence would allow an entirely new, and direct, probe of structure formation processes, leading to an improved understanding of inter- and intra-cluster gas dynamics, magnetic fields, and energy partitioning. While extrapolations from EGRET data and simulations for future instruments predict at least a dozen or more detectable sources \\citep{kes02}, direct determination of gamma-ray sources due to shocks can be used as an independent calibration for $\\xi_e$. In fact, if the value of $\\xi_e$ were lower than the inferred 5\\% (see discussion above), this would have an impact on the number of sources resolvable by future telescopes. \\citet{kes02} predict that the number drops to $\\sim 1$ for $\\xi_e\\sim 0.2$. GLAST will have the spectral and spatial resolution to confirm whether galaxy clusters can be directly detected in gamma-rays."
    },
    "0207/astro-ph0207282_arXiv.txt": {
        "abstract": "{ The extended photometry available for XZ Dra, a Blazhko type RR Lyrae star, makes it possible to study its long-term behavior. It is shown that its pulsation period  exhibit cyclic, but not strictly regular variations with $\\approx7200$ d period. The Blazhko period ($\\approx76$ d) seems to follow the observed period changes of the fundamental mode pulsation with ${\\rm d}P_{\\rm B}/{\\rm d}P_{0}=7.7\\times 10^{4}$ gradient. Binary model cannot explain this order of period change of the Blazhko modulation, nevertheless it can be brought into agreement with the $O-C$ data of the pulsation. The possibility of occurrence of magnetic cycle is raised. ",
        "introduction": "One of the unsolved problems, which is perhaps the most intriguing one in RR Lyrae star research, is the Blazhko effect, the amplitude and/or phase modulation of light curves with periods of $10 - 500$ days. Although theoretical interpretations of the phenomenon have been suggested \\citep{shiba,vanh}, a clear explanation of it is still lacking.  \\citet{blazhko} was the very first who recognized that the light maximum of an RR Lyrae star (namely RW Dra) showed phase modulation on a long time-scale (around 40 days). Subsequent investigation has revealed that the phase modulation is accompanied by modulation of the light curve and the amplitude of light variation. About a third of the fundamental mode RR Lyrae stars show the effect \\citep{szeidl}, however, only few have been investigated in detail yet. There are less than 10 galactic field stars which have sufficient observations to permit deeper insight into their Blazhko properties, but these studies have not been enough to expose the physical background of the observed modulation. The number of Blazhko stars which have been observed for a long enough time to detect any changes in their modulation properties is even fewer. Most of the stars listed in the summary review paper of Blazhko variables \\citep[completed in \\citealt{smith}]{szeidl} still have not been studied in detail.  XZ Dra (BD$+64\\degr 1332$, HIP\\,94134, $\\alpha_{2000}=19^{\\rm h} 09^{\\rm m} 42\\fs6$, $\\delta_{2000}=+64\\degr 51\\arcmin 32\\arcsec$) is one of the best observed RRab stars. \\citet{schn} discovered the star's variability on Babelsberg plates. Soon after the announcement of the discovery, \\citet{beyer} observed the star visually and determined the correct value of the fundamental period. He also commented on the strong oscillation in brightness of the individual light maxima. \\citet{bd}, based on the rough estimates of their photographic observations showed that these oscillations had a period of 76 days and the star's behaviour resembled that of AR Her.  During the last century, continuous effort has been made at the Konkoly Observatory to regularly observe RR~Lyrae stars with Blazhko effect. Collection of photometric and some radial velocity observations of XZ Dra has been recently published in \\citet{mitteil}. These data, together with all the published measurements of XZ~Dra, made it possible to follow its photometric behaviour during a remarkably long (70-year) period. Due to the extended data now available a detailed analysis of the properties of its pulsation and Blazhko behaviour has become feasible. ",
        "conclusions": "The 70 years photometric observations of XZ Dra have revealed that long period (7200 d) cyclic changes in the pulsation period have been occurring. The Blazhko period seems to follow this period change, exhibiting $3-5$ days full range of period change. This means that, besides RR Lyrae, XZ Dra is the second Blazhko variable which clearly shows indication of long-term cyclic behaviour. The manifestation of these long-term changes is, however, completely different for the two stars.  The long-term cyclic changes favour the magnetic rotator-pulsator model of the Blazhko modulation, by explaining the observed phenomena with changes in the global magnetic field structure and/or strength. However, to check the reality of this explanation, detailed theoretical work is needed.  Binary interpretation of the observations, although giving acceptable good fit to the data, fails to explain the detected range of the Blazhko period variation.  Rapid $O-C$ and radial velocity changes of XZ Dra are predicted to occur next time in the years $2011-2018$, when coordinated photometric and spectroscopic observations would greatly help to give correct answers to the presently unsolved questions."
    },
    "0207/astro-ph0207557_arXiv.txt": {
        "abstract": "We investigate, if the cosmic ray positron fraction, as reported by the HEAT and AMS collaborations, is compatible with the annihilation of neutralinos in the supergravity inspired Constrained Minimal Supersymmetric Model (CMSSM), thus complementing previous investigations, which did not consider  constraints from unification, electroweak symmetry breaking and the present Higgs limits at LEP. We scan over the complete SUSY parameter space and find that in the acceptable regions the neutralino annihilation into $b\\overline{b}$ quark pairs is the dominant channel and improves the fit to the experimental positron fraction data considerably compared to a fit with background only. These fits are  comparable to the fit for regions of parameter space, where the annihilation into $W^+W^-$ pairs dominates. However, these latter regions are ruled out by the present Higgs limit of 114 GeV from LEP.   \\vspace{1pc} ",
        "introduction": "The cosmic ray positron fraction at momenta above 7 GeV, as  reported by the HEAT collaboration, is difficult to describe by the background only hypothesis\\cite{HEAT}. The data at lower momenta agree with previous data from the AMS experiment\\cite{ams}. A contribution from the annihilation of neutralinos, which are the leading candidates to explain the cold dark matter in the universe, can improve the fits considerably\\cite{edsjo,kane}. The neutralinos are the Lightest Supersymmetric Particles (LSP) in  supersymmetric extensions of the Standard Model, which are stable, if R-parity is conserved.  This new multiplicative quantum number for the supersymmetric partners of the Standard Model (SM) particles is needed to prevent proton decay and simultaneously prevents the LSP a) to decay into the lighter SM particles and b) can only interact with normal matter by producing additional supersymmetric particles. The cross sections for the latter are typically of the order of the weak cross sections, so the LSP is ``neutrinolike'', i.e. it would form halos around the galaxies and consequently, it is an excellent candidate for dark matter.  In addition to being of interest for cosmo\\-logy, supersymmetry solves also many outstanding problems in particle physics, among them\\cite{rev}: 1) it provides a  unification of the strong and electroweak forces, thus being a prototype theory for a Grand Unified Theory (GUT) 2) it predicts spontaneous electroweak  symmetry breaking (EWSB) by radiative corrections through the heavy top quark, thus providing a relation between the GUT scale, the electroweak scale and the top mass, which is perfectly fullfilled 3) it includes gravity 4) it cancels the quadratic divergencies in the Higgs mass in the SM 5) the lightest Higgs mass can be calculated to be below 125 GeV in perfect agreement with electroweak precision data, which prefer indeed a light Higgs mass. Therefore it is interesting to study the positron fraction from neutralino annihilation in the reduced region of SUSY parameter space, where these constraints are satisfied and compare it with the AMS and HEAT data, as  will be done in the next section.     \\begin{figure} \\begin{center} \\includegraphics [width=0.49\\textwidth,clip]{mass-evol.ltb.eps} \\caption[]{\\label{mass} \\it The running of the masses between the GUT scale and the electroweak scale. } \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0207/hep-ph0207254_arXiv.txt": {
        "abstract": "The observed interaction energy of cosmic rays with atmospheric nuclei reaches up to a PeV in the center of mass. We compute nucleon-nucleon and nucleon-neutrino cross sections for various generic parton cross sections appearing in string and brane world scenarios for gravity and compare them with cosmic ray data. Scenarios with effective energy scales in the TeV range and parton cross sections with linear or stronger growth with the center of mass energy appear strongly constrained or ruled out. String-inspired scenarios with infinite-volume extra dimensions may require a fundamental scale above $\\simeq100\\,$eV for which they are probably in conflict with gravity on parsec scales. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/nlin0207021_arXiv.txt": {
        "abstract": "We present a dimension analysis of a set of solar type~I storms and type~IV events with different kind of fine structures, recorded at the Trieste Astronomical Observatory. The signature of such types of solar radio events is highly structured in time. However, periodicities are rather seldom, and linear mode theory can provide only limited interpretation of the data. Therefore, we performed an analysis based on methods of the nonlinear dynamics theory.  Additionally to the commonly used correlation dimension, we also calculated local pointwise dimensions. This alternative approach is motivated by the fact that astrophysical time series represent real-world systems, which cannot be kept in a controlled state and which are highly interconnected with their surroundings. In such systems pure determinism is rather unlikely to be realized, and therefore a characterization by invariants of the dynamics might probably be inadequate.  In fact, the outcome of the dimension analysis does not give hints for low-dimensional determinism in the data, but we show that, contrary to the correlation dimension method, local dimension estimations can give physical insight into the events even in cases in which pure determinism cannot be established. In particular, in most of the analyzed radio events nonlinearity in the data is detected, and the local dimension analysis provides a basis for a quantitative description of the time series, which can be used to characterize the complexity of the related physical system in a comparative and non-invariant manner.  In this frame, the degree of complexity we inferred for type~I storms is on the average lower than that relevant to type~IV events. For the type~IV events significant differences occur with regard to the various subtypes, whereas pulsations and sudden reductions can be described by distinctly lower values than spikes and fast pulsations. ",
        "introduction": "Nonlinear time series analysis based on the theory of deterministic chaos has turned out to be a powerful tool in understanding complex dynamics from measurements and observational time series. In particular it can provide descriptions and interpretations for irregular times series, which nevertheless might not be governed by a stochastic physical process and which are only poorly understood by linear methods. A number of recent reviews and conference proceedings shows the great interest in the field of nonlinear time series analysis (see, for instance, Grassberger et al. \\cite{GrassbergerEtal91}; Casdagli \\& Eubank \\cite{CasdagliEubank92}; Weigend \\& Gershenfeld \\cite{WeigendGershenfeld93}; Kugiumtzis \\cite{KugiumtzisEtal94a}, \\cite{KugiumtzisEtal94b}; Abarbanel \\cite{Abarbanel96}; Kantz \\& Schreiber \\cite{KantzSchreiber97}; Schreiber \\cite{Schreiber99}).  Since the development of chaos theory, it is well known that even simple dynamical systems, described by few nonlinear differential equations, can reveal a complex and quasi-irregular behavior. A central concept to characterize such systems is the so-called {\\em attractor}. Under the dynamics of a deterministic system the trajectories do not cover the whole phase space, but, after all transient phenomena have faded out, converge to a subset of the phase space, the attractor. The attractor itself is invariant to the dynamical evolution. Simple examples of attractors are fixed points and limit cycles. However, when the related dynamical system is {\\em chaotic}, the attractor can have a complex geometry with a {\\em fractal}, i.e., non-integer, dimension. Different invariant parameters exist to describe the geometry and the dynamics of an attractor, such as dimensions, Lyapunov exponents and entropies. Besides the fractal geometry, chaotic systems have the striking property that initially neighboring trajectories diverge exponentially under the dynamics, and the growth rate is given by the {\\em Lyapunov exponent}. This phenomenon results from the folding and stretching of the trajectories under the dynamics, the folding leading to the convergence of the trajectories to the attractor and the stretching to the divergence in certain directions. While the average stretching rate is given by the Lyapunov exponent, the loss of information due to the folding is quantified by the {\\em entropy}. Lyapunov exponents and entropies characterize the dynamics on the attractor, and the {\\em dimension} characterizes its geometry. The physical meaning of the dimension of an attractor is that it corresponds to the degree of freedom of the related dynamical system, i.e., in a deterministic case, the minimum number of ordinary differential equations needed to fully describe the system. Deterministic systems are characterized by a finite dimension. Deterministic chaotic systems have the additional characteristic that their dimension is fractal. Contrary to that, a stochastic system is characterized by an infinite dimension, indicating its infinite degree of freedom. Therefore, the determination of the dimension of an attractor enables to discriminate whether a dynamical system is deterministic or stochastic.  Previous papers exist concerning the investigation of fractal dimensions of solar radio bursts. It has to be noted, that the time series used by the different authors are not directly comparable as they represent different types of radio events. Kurths \\& Herzel (\\cite{KurthsHerzel86}, \\cite{KurthsHerzel87}),  Kurths \\& Karlick$\\acute{\\rm y}$ (\\cite{KurthsKarlicky89}), and Kurths et al. (\\cite{KurthsEtal91}) analyzed decimetric pulsations and ascertained finite dimension values. Contrary to that, Isliker (\\cite{Isliker92b}) and Isliker \\& Benz (\\cite{IslikerBenz94a}, \\cite{IslikerBenz94b}) investigated different types of solar radio bursts in the metric~(m) and decimetric~(dm) wavelength range (type~I storms, type~II bursts, type~III bursts, type~IV events, and narrowband spikes), which did not reveal any hints for low-dimensional determinism.\\footnote{The reported finite dimension for one of the analyzed narrowband spike events in Isliker (\\cite{Isliker92b}) was revised in a later paper (Isliker \\& Benz \\cite{IslikerBenz94a}).}  However, these investigations rely all on the correlation dimension method. The present paper additionally introduces a complementary dimension analysis, motivated by the fact that solar radio bursts represent real-world systems, which implies some major restrictions. First, the time series cannot be expected to be stationary, and second, pure determinism is rather unlikely to be realized. Therefore we do not only concentrate on the usual way of looking at the problem: ``Does the analyzed time series represent a deterministic or a stochastic system?\" but in particular focus the question: ``What statistical description can be extracted from a dimension analysis of the time series?\" With such refined formulation of the problem, we try to make use of the concepts and tools of nonlinear time series analysis even in cases in which the determination of invariants of the dynamics, as, e.g., attractor dimensions, possibly fails.  The paper is structured as follows. Sect.~\\ref{Methods} explains the used methods and discusses critical points in the determination of fractal dimensions from time series. In Sect.~\\ref{DataSets} the investigated data sets are characterized and the analysis procedure is described. Sect.~\\ref{Results} presents the results of the dimension analysis, which are discussed in Sect.~\\ref{Discussion}. Finally, the conclusions are drawn in Sect.~\\ref{Conclusion}. ",
        "conclusions": "}  In the following items we give a summary of the main results obtained by the presented dimension analysis of several types of solar radio events, based on the correlation dimension and the local pointwise dimension method. The results are relevant concerning the physics of the analyzed events as well as the different methods applied. \\begin{enumerate} \\item The analysis does not enable to claim low-dimensional determinism in the time series. This outcome is in agreement with the results obtained by Isliker (\\cite{Isliker92b}) and Isliker \\& Benz (\\cite{IslikerBenz94a}, \\cite{IslikerBenz94b}), who also, among others, investigated type~I storms, type~IV events, and spikes. We cannot confirm the results of Kurths \\& Herzel (\\cite{KurthsHerzel86}, \\cite{KurthsHerzel87}), Kurths \\& Karlick$\\acute{\\rm y}$ (\\cite{KurthsKarlicky89}), and Kurths et al. (\\cite{KurthsEtal91}), who obtained finite dimension values for decimetric pulsations. However, the outcome of the present paper does not exclude deterministic chaos in the analyzed time series but makes pure low-dimensional determinism, characterized by few free parameters, rather improbable. \\item The analyzed time series are not fully stochastic, i.e. white noise. This fact we infer from the distinctly slower increase of the dimension values with increasing embedding dimension than expected for fully stochastic processes, which always fill the whole phase space, i.e. $D(m) \\approx m$. Moreover, the surrogate data analysis suggests that the time series do not represent linear stochastic processes. \\item For most of the analyzed data sets we have evidence that nonlinearity in the time series is present (given on a $3\\sigma$~level by means of a surrogate data test). \\item A comparison of the two different methods used for the determination of fractal dimensions reveals that the local dimension method is more stable and enables more physical insight than the classical correlation dimension method. The local dimension analysis can provide a statistically significant quantity for systems, which cannot be characterized by invariants of the dynamics, probably since they are in fact not purely deterministic. Such quantities can be of special interest for comparative studies, investigating interrelations between different time series (which, e.g., can be useful for classificational purposes) or investigating intrarelations in between one time series (in order to detect dynamical changes). \\item The retrieved pointwise dimension values can be interpreted in terms of complexity of the underlying physical system. In this frame our analysis indicates that spikes and fast pulsations are the signature of systems of higher complexity than pulsations, sudden reductions and type~I storms. \\end{enumerate}  In relation with other kind of analysis of solar radio bursts the presented results might give further ideas on the physics of the events. In the following we present a short discussion in this respect, applied to pulsation and spike events, which are quite striking features associated with solar flares.  Spikes have been intensively studied during recent times. Their short duration and small bandwidth gives rise to the evidence that they are associated with the energy fragmentation process in solar flares (Benz \\cite{Benz85}, \\cite{Benz86}). Based on this connection, Schwarz et al. (\\cite{SchwarzEtAl93}) performed a nonlinear analysis by means of symbolic dynamics methods, interpreting the spikes appearance in the frequency-time domain as spatio-temporal patterns. This analysis gives indications that the simultaneous appearance of spikes at different frequencies is not a purely stochastic phenomenon but may be caused by a non\\-linear deterministic (not necessarily low-dimensional) system or by a Markov process, compatible with a scenario in which spikes at nearby locations are simultaneously triggered by a common exciter, i.e. the localized sources are causally connected. In the present paper we find evidence for the spike events analyzed, that they do not represent a purely stochastic phenomenon in their temporal order either, even if the degree of freedom of the related physical system is expected to be quite high. Interpreting this result in the frame of the scenario suggested by Schwarz et al. (1993), it might give indications that the triggering of successive spikes by a localized source is not caused by a fully stochastic process, but reveals some (possibly weak) kind of nonlinear causal connection. However, this inference is restricted to the assumption that the spikes time series rather reflect the physical conditions of the triggering mechanism than those of the emission.  Pulsations, although a rather marginal phenomenon in the course of solar flares, have reached a wealth of attention, especially due to the very regular features they sometimes reveal (for a review see Aschwanden \\cite{Aschwanden87}). In previous investigations of the dimensionality of solar pulsations (Kurths \\& Herzel \\cite{KurthsHerzel86}, \\cite{KurthsHerzel87}; Kurths \\& Karlick$\\acute{\\rm y}$ \\cite{KurthsKarlicky89}; Kurths et al. \\cite{KurthsEtal91}) the presence of low-dimensional determinism is reported, with dimensions \\mbox{$2.5 \\lesssim D \\lesssim 3.5$}. Moreover, for one single event a dynamical evolution from a limit cycle to a low-dimensional chaotic behavior was found (Kurths \\& Karlick$\\acute{\\rm y}$ \\cite{KurthsKarlicky89}). Although we cannot confirm these results, we want to stress that our analysis suggests that pulsation events, especially quasi-periodic pulsations, represent the least complex phenomena among the analyzed types of radio events, i.e. their degree of freedom is expected to be lower than that of other burst types, even if not low-dimensional. The inferred high-dimensionality and the nonlinear structures detected do not match with linear MHD oscillation models for pulsations (e.g., Rosenberg \\cite{Rosenberg70}; Roberts et al. \\cite{RobertsEtAl84}), in which only a few eigenmodes are excited, but rather favor models of self-organizing systems of plasma instabilities, which comprise periodic as well as low- and high-dimensional chaotic behavior. Such a self-organizing model for the electron-cyclotron maser instability, based on a Lotka-Volterra type equation system, is discussed in Aschwanden \\& Benz (\\cite{AschwandenBenz88}), however restricted to limit cycle solutions."
    },
    "0207/astro-ph0207377_arXiv.txt": {
        "abstract": "We present the first analysis of three-dimensional genus statistics for the SDSS EDR galaxy sample. Due to the complicated survey volume and the selection function, analytic predictions of the genus statistics for this sample are not feasible, therefore we construct extensive mock catalogs from N-body simulations in order to compare the observed data with model predictions.  This comparison allows us to evaluate the effects of a variety of observational systematics on the estimated genus for the SDSS sample, including the shape of the survey volume, the redshift distortion effect, and the radial selection function due to the magnitude limit.  The observed genus for the SDSS EDR galaxy sample is consistent with that predicted by simulations of a $\\Lambda$-dominated spatially-flat cold dark matter model. Standard ($\\Omega_0=1$) cold dark matter model predictions do not match the observations. We discuss how future SDSS galaxy samples will yield improved estimates of the genus. ",
        "introduction": "Characterizing the large-scale distribution of galaxies and clusters is critical for understanding the formation and evolution of this structure as well as for probing the initial conditions of the universe itself. The most widely used statistical measure for this purpose is the two-point correlation function (2PCF).  This statistic is easily estimated, is simply parameterized, and is well-studied.  The observed correlation function is approximately a power law over a finite range of scales (Totsuji \\& Kihara 1969), and thus can be expressed by two numbers, the power-law slope and the correlation length. This correlation function statistic has been successfully applied in the cosmological context for more than 30 years, thus its behavior is well understood theoretically. Given a set of cosmological parameters, one can predict the corresponding mass 2PCF $\\xi(r,z)$ on a scale $r$ at a redshift $z$ using accurate fitting formulae (e.g., Hamilton et al. 1991; Peacock \\& Dodds 1996).  In principle, the spatial biasing of luminous objects relative to the underlying dark matter may invalidate the straightforward comparison between observations and theoretical predictions.  However, the SDSS galaxy sample which we analyze is consistent with a (practically) scale-independent linear bias on scales of $(0.2 - 4)h^{-1}$Mpc if one adopts the currently popular $\\Lambda$-dominated spatially-flat cold dark mater (LCDM) model (Kayo et al. 2002).  While the 2PCF and its Fourier transform the power spectrum are the most convenient and useful cosmological statistics, they completely ignore information about the correlations of the phases of the density fluctuations in $k$-space.  In contrast, the topology of large-scale structure, as measured by the genus statistic (Gott, Melott, \\& Dickinson 1986) is strongly sensitive to these phase correlations.  One of the Minkowski functionals (Mecke, Buchert \\& Wagner 1994; Schmalzing \\& Buchert 1997) is closely related to the genus.  Other statistics that quantify phase correlations in the data include higher-order n-point correlation functions (e.g., Peebles 1980 and references therein), percolation analysis (Shandarin 1983), minimal spanning trees (Barrow, Bhavsar \\& Sonoda 1985), and void statistics (White 1979).  We focus on the genus mainly because its behavior is well-understood theoretically; if the primordial fluctuations are Gaussian, then the genus in the linear regime has an exact analytic expression (see eq. [\\ref{eq:genus_rd}] below). The theoretical prediction for a Gaussian random field makes it clear that only the amplitude of the genus depends on the 2PCF, while higher order correlations determine the shape as well as affecting the amplitude.  Thus, the shape of the genus statistic plays a complementary role to the 2PCF. In addition to the linear theory predictions, a perturbative expression for the genus in the weakly nonlinear regime has been obtained by Matsubara (1994) and the log-normal model is shown to be a good empirical approximation in the strongly nonlinear regime (e.g., Coles \\& Jones 1991; Hikage, Taruya \\& Suto 2002).  Previous to the SDSS, investigations of the 3D genus statistic for galaxy redshift surveys include Gott et al. (1989), Park, Gott \\& da Costa (1992), Moore et al. (1992) Rhoads et al. (1994), Vogeley et al. (1994), and Canavezes et al. (1998). As suggested by Melott (1987), the 2D variant of this statistic has been estimated for a variety of data sets by Coles \\& Plionis (1991), Plionis, Valdarnini \\& Coles (1992), Park et al. (1992), Colley (2000), Park, Gott, \\& Choi (2001), and Hoyle, Vogeley \\& Gott (2002a). These investigations generally indicate consistency with the hypothesis that the initial perturbations were Gaussian in nature. Some departures from Gaussianity have been suggested, but the statistical significance of the results was low due to the small size of the available data sets.  Hoyle et al. (2002b) find weak evidence for variation in the genus with galaxy type in the SDSS using the 2D genus statistic.  In the present paper, we describe the first analysis of the 3D genus of the SDSS early data release galaxy sample.  We evaluate a variety of observational effects using mock catalogs from N-body simulations such as the shape of the survey volume, the redshift distortion effect, and the radial selection function due to the magnitude limit. To within the uncertainties for this preliminary sample, we find that the LCDM model reasonably reproduces the observed shape and amplitude of the genus of SDSS galaxies. A complementary analysis of the 2D genus for the SDSS galaxies is presented separately in Hoyle et al. (2002b).  Other analyses of early SDSS data include measurement of the power spectrum (Szalay et al. 2002; Tegmark et al. 2002; Dodelson et al. 2002), correlation function (Zehavi et al. 2002; Connolly et al. 2002; Infante et al. 2002), and higher-order moments (Szapudi et al. 2002).  The outline of the paper is as follows. In section 2 we describe the SDSS Early Data Release sample that we analyze. To analyze systematic effects we use mock samples drawn from N-body simulations, which are described in section 3. Section 4 describes the genus statistic and the method of estimation. Observational systematics are analyzed in section 5. Genus results for the SDSS EDR sample are presented in section 6. Section 7 presents our conclusions. ",
        "conclusions": "We present the first analysis of the three-dimensional genus statistics for SDSS EDR galaxy data.  Due to the complicated survey volume and the selection function of the current sample, analytic predictions of the corresponding genus statistics are not feasible and we construct extensive mock catalogs from N-body simulations in order to explore the cosmological implications. We also use these mock catalogs to examine the effects of several possible observational systematics on the genus statistic.  We find that redshift-space distortions cause minimal changes in the genus curve.  However, the slice-like geometry of the SDSS EDR galaxy samples causes large distortions of the genus curve on smoothing scales $R_{\\rm G}> 7 h^{-1}$Mpc. We demonstrate that weighting by the inverse of the selection function with distance in apparent-magnitude limited samples allows recovery of the genus curve. Because the mock catalogs include all of these effects, comparison with theoretical models is possible.  We conclude that the observed shape and amplitude of the genus for the SDSS EDR galaxy sample are consistent with the distribution of dark matter particles in simulations of the LCDM model. In contrast, the SCDM model predictions do not match the amplitude of the observed genus.  Comparison of the observed genus curves with the theoretical prediction for a Gaussian random field shows that, within the uncertainties as estimated from the mock samples, the data are consistent with Gaussianity of the primordial density field.  A question for the future is the degree of improvement that one can expect for the larger samples of the SDSS galaxy redshift survey. To predict this, we generate larger mock catalogs by increasing the sky coverage by three times ($\\sim 7\\times 10^4$ galaxies) and thirty times ($\\sim 7\\times 10^5$ galaxies) relative to the Northern stripe in EDR. Figure \\ref{fig:future} shows the predicted genus in LCDM and SCDM models (see also Colley et al.  2000 for predictions of genus results for the SDSS). The positions of the symbols in those panels correspond to the data of one specific mock sample, but with error-bars estimated from the 15 independent mock samples. For reference, the lines in upper panels show the mean genus curves averaged over the 15 independent mock samples, while the lines in lower panels indicate the fit of the random-Gaussian prediction to the plotted sample data.  Of interest is that the smaller uncertainties in this future sample will put much stronger constraints on cosmological models, the nonlinear nature of the galaxy biasing, and primordial Gaussianity than those presented here. It may also be worth while to apply theoretical estimate of genus statistics based on Lagrangian perturbation theory after appropriate smoothing (Seto et al. 1997).  \\begin{figure}[htp] \\begin{center} \\FigureFile(80mm,80mm){figure22.eps} \\FigureFile(80mm,80mm){figure23.eps} \\FigureFile(80mm,80mm){figure24.eps} \\FigureFile(80mm,80mm){figure25.eps} \\end{center} \\caption{Prediction of the genus for future samples of SDSS galaxies, using mock catalogs in LCDM ({\\it Open circles} and {\\it Solid lines}) and SCDM ({\\it Open triangles} and {\\it Dotted lines}) From left to right, $R_{\\rm G}=5h^{-1}$Mpc with $7 \\times 10^4$ galaxies, $R_{\\rm G}=10h^{-1}$Mpc with $7 \\times 10^4$ galaxies, $R_{\\rm G}=5h^{-1}$Mpc with $7 \\times 10^5$ galaxies, and $R_{\\rm G}=10h^{-1}$Mpc with $7 \\times 10^5$ galaxies.  Upper and lower panels plot the total number of genus in terms of $\\nus$ and $\\nuf$, respectively.} \\label{fig:future} \\end{figure}  \\bigskip  We thank Y. P. Jing for kindly providing us his N-body simulation data which were used in generating mock samples.  We also thank T. Buchert for a careful reading of this manuscript. I. K. gratefully acknowledges support from the Takenaka-Ikueikai fellowship.  Numerical computations were carried out at ADAC (the Astronomical Data Analysis Center) of the National Astronomical Observatory, Japan.  This research was supported in part by the Grant-in-Aid from Monbu-Kagakusho and Japan Society of Promotion of Science (12640231, 13740150, 14102004, and 1470157). MSV and FH acknowledge support from NSF grant AST-0071201 and a grant from the John Templeton Foundation. JRG acknowledges support from NSF grant AST-9900772.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbu-Kagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions, which are the University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, the Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Princeton University, the United States Naval Observatory, and the University of Washington.  \\clearpage"
    },
    "0207/astro-ph0207141_arXiv.txt": {
        "abstract": "We report the results obtained from ISAAC/VLT near-IR spectroscopy of two low-luminosity $z \\sim 1.9$ galaxies located in the core of the lensing cluster AC114. The amplification factor allowed to obtain, for the first time, physical properties (SFR, abundance ratios, mass, etc) of star-forming galaxies, 1 to 2 magnitudes fainter than in previous studies of LBGs at $z \\sim 3$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207066.txt": {
        "abstract": "{It is generally accepted that relativistic jet outflows power the nonthermal emission from active galactic nuclei (AGN). The composition of these jets -- leptonic versus hadronic -- is still under debate. We investigate the microphysical details of the conversion process of the kinetic energy in collimated relativistic pair outflows into radiation through interactions with the ambient interstellar medium. Viewed from the coordinate system comoving with the pair outflow, the interstellar protons and electrons represent a proton-electron beam propagating with relativistic speed in the pair plasma. We demonstrate that the beam excites both electrostatic and low-frequency magnetohydrodynamic Alfven-type waves via a two-stream instability in the pair background plasma, and we calculate the time evolution of the distribution functions of the beam particles and the generated plasma wave turbulence power spectra. For standard AGN jet outflow and environment parameters we show that the initial beam distributions of interstellar protons and electrons quickly relax to  plateau-distributions in parallel momentum, transferring thereby one-half of the initial energy density of the beam particles to electric field fluctuations of the generated electrostatic turbulence. On considerably longer time scales, the plateaued  interstellar electrons and protons will isotropise by their self-generated transverse turbulence and thus be picked-up in the outflow pair plasma. These longer time scales are also characteristic for the development of transverse hydromagnetic turbulence from the plateaued electrons and protons. This hydromagnetic turbulence upstream and downstream is crucial for diffusive shock acceleration to operate at external or internal shocks associated with pair outflows. ",
        "introduction": "The detection of intense medium-energy gamma radiation from over 60 blazar active gactic nuclei (hereafter abbreviated as AGNs) with the EGRET instrument on the Compton observatory (Hartman et al. \\cite{h99}) and TeV gamma radiation from several BL-Lac AGNs (Pohl \\cite{p01}) shows that nonthermal gamma-ray production is a significant dissipation mechanism of jet energy generated by black-hole accretion. Besides the modelling of the broadband nonthermal radiation, the composition of the jet plasma -- i.e. electron-positron pair jets (leptonic jets) versus electron-proton jets (hadronic jets) -- and the acceleration of these jet particles to relativistic energies are main subjects of the current theoretical research.  Gamma radiation in leptonic models of broadband blazar emission is attributed to synchrotron self-Compton (Maraschi et al. \\cite{mgc92}; Bloom \\& Marscher \\cite{bm96}; Tavecchio et al. \\cite{tmg98}) or external Compton (Dermer \\& Schlickeiser \\cite{ds93}; Sikora et al. \\cite{sbr94}; B\\\"ottcher et al. \\cite{bms97}; Dermer et al. \\cite{dss97}; Arbeiter et al. \\cite{aps02}) processes (see B\\\"ottcher \\cite{b01} and Sikora \\& Madejski \\cite{sm01} for recent reviews). In hadronic models, secondary photopairs and photomesons and secondary mesons are produced when energetic protons and ions interact either with ambient synchrotron photons (Mannheim \\& Biermann \\cite{mb92}; Mannheim \\cite{m93}), photons of the external field (Bednarek \\& Protheroe \\cite{bp99}; Atoyan \\& Dermer \\cite{ad01}) and/or ambient matter fields (Pohl \\& Schlickeiser \\cite{ps00}). Observationally, future detection of high-energy neutrino emission correlated with high-energy photon emission (Schuster et al. \\cite{sps02}) will provide the ultimate test between the leptonic and hadronic jet models.  Most existing radiation models of AGN jets are very unspecific on the microphysical details of the conversion of the kinetic jet energy into energetic charged particles and subsequently into radiation. Without detailed discussion these models often assume that a significant fraction of the accreted kinetic energy is injected into nonthermal pairs and/or hadrons with power-law distribution functions. This efficient conversion is attributed to the scenario that the outflowing relativistic jet plasma has produced a relativistic shock with fully developed hydromagnetic turbulence in order to allow for efficient non-thermal diffusive particle acceleration at the collisionless shock fronts:  \\noindent (1) This neglects the fact that it takes a finite time to build up the necessary turbulence in the two-stream multi-fluid system consisting of the relativistically moving jet plasma and the traversed interstellar or intergalactic hydrogen plasma.  \\noindent (2) It is not clear from the beginning that the outflowing relativistic jet plasma will generate primarily (and enough) transverse magnetohydrodynamic turbulence, which is crucial for efficient particle deflections in the up- and downstream region of the shock waves. It is well conceivable that most of the kinetic blast wave energy is transferred to electrostatic plasma turbulence and not to transverse hydromagnetic turbulence.  \\noindent (3) The properties of the generated plasma turbulence are decisive both for the formation and the nature of the developing collisionless shock waves. It is known from non-relativistic shock theory that the inclusion of the finite pressure and energy flux of the generated plasma turbulence in the Rankine-Hugoniot shock relations strongly modifies the standard fluid shock properties (Vainio \\& Schlickeiser \\cite{vs99, vs01}; Lerche et al. \\cite{lps00}) and subsequently the energy spectrum of the accelerated nonthermal particles.  It is the purpose of this work to consider more thoroughly some of the microphysical details of the energy conversion in relativistic jet outflows. We consider the energisation of relativistic particles in the jet by interactions with the surrounding medium following the earlier work of three of us (Pohl \\& Schlickeiser \\cite{ps00}; Pohl et al. \\cite{pls02}). There the AGN jet has been assumed to be a cloud of dense {\\it electron-proton} plasma which moves relativistically through the electron-proton interstellar medium of the AGN host galaxy. The plasma cloud is assumed to have a cylindrical shape with thickness $d^*=10^{13}d^*_{13}$ cm, which is small compared to the radius $r^*=10^{14}r^*_{14}$ cm of the cylinder (see Fig. 1 of Pohl \\& Schlickeiser (\\cite{ps00}) for a sketch of the assumed cloud geometry). It was shown that in such {\\it hadronic jets} swept-up ambient matter is quickly isotropised in the jet cloud frame by a relativistic two-stream instability, which provides relativistic particles in the jet cloud without invoking any acceleration process. Here and in the following the index $*$ indicates the values of physical quantities in the laboratory (AGN) frame; quantities not indexed are in the jet frame.  Here we study the {\\it leptonic variant} of this jet outflow model, i.e. we model the outflowing jet cloud as a one-dimensional channeled outflow of thickness $d^*$ with relativistic bulk velocity $\\vec{V}$, consisting of {\\it pairs of electrons and positrons} of density $n_b^*$ instead of {\\it electrons and protons}. To avoid dramatic pair annihilation, the density of pairs must be limited, and we allow for a thermal pair distribution with non-relativistic temperature $\\Theta =k_BT_{\\rm pair}/(m_ec^2)<<1$ in the jet rest frame. This beam of pairs propagates into the surrounding interstellar medium that consists of cold protons and electrons at rest of density $n_i^*$.  For mathematical convenience we assume that the outflow is directed parallel to a uniform background magnetic field. The assumption of the magnetic field directed along the ejecta velocity enormously facilitates the analytical treatment of the two-stream instability, but crucially depends on the location of the particle energisation with respect to the large-scale magnetic field and the confinement of the jet ejecta. In magnetohydrodynamic models of jets in accreting systems the poloidal magnetic field component is more strongly ($\\propto R^{-2}$) suppressed by the side expansion of the ejecta than the transverse magnetic field component ($\\propto R^{-1}$), so that our assumption will hold at close distances to the central object. Moreover, in the presence of a dominating transverse magnetic field it can be argued that the ambient medium  will penetrate the ejecta by one Larmor radius only and the momentum exchange should be faster and more efficient than in the case of a dominating poloidal magnetic field considered here. {\\footnote {We are grateful to the referee for noting this difference.}}  Viewed from the coordinate system comoving with the pair outflow, the interstellar protons and electrons represent a proton-electron beam propagating with relativistic speed $-\\vec{V}$ antiparallel to the uniform magnetic field direction. Modifying the analysis of Pohl \\& Schlickeiser (\\cite{ps00}) and Pohl et al. (\\cite{pls02}) for cold electron-proton outflows to finite temperature pair outflows, we demonstrate that the beam excites both electrostatic and low-frequency magnetohydrodynamic Alfven-type waves via a two-stream instability in the pair background plasma. We study the time evolution of the beam particles, the generated plasma wave turbulence power spectra and discuss the radiation signatures of such systems. % ",
        "conclusions": "% We have investigated the microphysical details of the energy conversion in relativistic pair outflows interacting with the surrounding interstellar medium consisting of cold protons and electrons. We have represented the relativistic pair blast wave as a one-dimensional channeled outflow directed parallel to a uniform background magnetic field. Viewed from the coordinate system comoving with the pair outflow, the interstellar protons and electrons represent a proton-electron beam propagating with relativistic speed antiparallel to the uniform magnetic field direction. We demonstrate that the beam excites both electrostatic and low-frequency magnetohydrodynamic Alfven-type waves via a two-stream instability in the pair background plasma, and we calculate the time evolution of the distribution functions of the beam particles and the generated plasma wave turbulence power spectra.  For standard AGN jet outflow and environment parameters we show that the initial beam distributions of interstellar protons and electrons quickly relax to  plateau-distributions in parallel momentum, transferring thereby one-half of the initial energy density of the beam particles to electric field fluctuations of the generated electrostatic turbulence.  On considerably longer time scales, the plateaued  interstellar electrons and protons will isotropise by their self-generated transverse turbulence and thus be picked-up in the outflow pair plasma. These longer time scale are also characteristic for establishing fully developed power spectra of transverse hydromagnetic turbulence from the plateaued electrons and protons. This hydromagnetic turbulence upstream and downstream is crucial for diffusive shock acceleration to operate at external or internal shocks associated with pair outflows.  It takes a finite time period of the early AGN outflow phase to build up the magnetohydrodynamic turbulence for nonthermal particle acceleration at either internal or external shocks associated with pair jet outflows. During this initial phase -- prior to the generation of hydromagnetic turbulence -- the radiative and energy-exchange processes in the jet outflow are not dominated by acceleration processes of nonthermal charged particles at collisionless shock waves."
    },
    "0207/astro-ph0207155_arXiv.txt": {
        "abstract": "{We describe a scheme for the formation of globular cluster systems in early-type galaxies using a semi-analytic model of galaxy formation. Operating within a $\\Lambda$CDM cosmology, we assume that metal-poor globular clusters are formed at high-redshift in pre-galactic fragments, and that the subsequent gas-rich merging of these fragments leads to the formation of metal-rich clusters. We compare our results with contemporary data, and look at the particular case of the globular cluster and stellar metallicity distribution function of the nearby elliptical Centaurus A.}  \\resumen{Describimos el escenario para la formacion de sistemas de cumulos globulares en galaxias elipticas utilizando un modelo semi-analitico de formacion de galaxias. Trabajando en una cosmologia de tipo $\\Lambda$CDM, asumimos que los cumulos globulares de baja metalicidad se forman a altos redshifts en fragmentos pre-galacticos, y que la consiguiente fusion rica en gas de estos fragmentos conduce a la formacion de los cumulos globulares de alta metalicidad. Comparamos nuestros resultados con datos contemporaneos y estudiamos el caso particular de la funciones de distribucion de metalicidad estelar y de los cumulos globulares de la galaxia eliptica cercana Centaurus A.}  \\addkeyword{Stars: Star Clusters} \\addkeyword{Galaxies: Elliptical}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207363_arXiv.txt": {
        "abstract": "Massive binary star systems are not uncommon, and neither the supersonic collision of their winds. In the present work we study these shocks and the further consequences on wind structure. The post-shock gas is a warm and high-density environment, which allows dust to form and grow. We show that this growth is fast, of just a few hours. An application for $\\eta$Car shows that, probably, the decline of X-rays fluxes observed in its light curve is the consequence of its high absorption in periodic dust formation events, on the periastron passage. ",
        "introduction": "Many massive binary systems are known, and almost all present high-energy fluxes, not originated in the stars, but around them, indicating the presence of colliding winds. The source of observed X-rays high fluxes should be a hot and dense gas, generated by shocks around these objects (Zhekov \\& Skinner 2000). At high temperature and density, free-free emission in X-rays becomes highly important (Usov 1992; Ishibashi et al. 1999). X-ray emissions associated with interacting winds were observed on many objects (Corcoran et al. 2001; Thaller et al. 2001). Grains also may be formed and grow on post-shocked gases (Monnier, Tuthill, \\& Danchi 2001) when the heated gas loses high amount of energy, mainly by {\\it bremsstrahlung}, cools and becomes denser. Recent observations in IR supply this idea, showing episodic dust formation events in these regions (Marchenko, Moffat, \\& Grosdidier 1999; Harries, Babler, \\& Fox 2000). Some binary systems present, periodically, a sudden decrease of X-ray emission (Ishibashi et al. 1999). A fast dust formation and growth event in these regions due to the shocked winds might provoke a considerable increase in extinction explaining these features of the light curves. This could cause the temporary reduction on X-rays and UV fluxes, which may increase {\\it a posteriori} with the evaporation of the particles, or even by the expansion of the dust shell. ",
        "conclusions": "In this work we present a model for colliding winds where we determined the changes in physical parameters, like density, temperature and pressure in the post-shock. This gas, cooled and denser after shock, can be the site for grains formation. The model was applied to $\\eta Car$ which could explain: i-) the rapid decline on fluxes observed periodically; and ii-) the slightly decrease in opacity considering the expansion of the dust with the wind. This model agrees better with observations for $e \\sim 0.8$, and a mass loss rate of the primary star of $\\sim (3 - 5)\\times10^{-4} \\; M_\\odot \\; yr^{-1}$."
    },
    "0207/astro-ph0207649_arXiv.txt": {
        "abstract": "{\\small We follow the evolution of a low frequency QPO during the 2000 outburst of the microquasar XTE J1550--564, which was found to be present in the PCA energy range (2--65 keV) in 19 of 43 observations. The frequency of the QPO varies from 0.1 Hz to 6 Hz, and appears to follow the evolution of the soft X-ray flux. If we assume the soft X-rays represent the behavior of an accretion disk, the relation indicates that this low frequency QPO is linked to the accretion disk. We show that the non-trivial relation between the QPO frequency and the soft flux may be as expected from the Accretion Ejection Instability (AEI), when the disk approaches its last stable orbit. Furthermore, the energy dependence of the QPO may indicate the presence of a hot spot rotating in the disk as predicted by the AEI.} ",
        "introduction": "Low frequency Quasi Periodic Oscillations (LFQPO) are commonly observed in microquasars when a strong power law component is detected in their X-ray energy spectra, with a typical frequency of 0.1--10 Hz. During the Low Hard State (LHS) they have high amplitude ($\\sim 15\\%$ rms), and during the Intermediate/Very High State (IS/VHS) they have a moderate amplitude ($\\sim 5\\%$ rms). Their presence can be explained in the context of the Accretion Ejection Instability (AEI, \\cite{Tagger99}), which occurs in the innermost region of an accretion disk threaded by vertical magnetic field lines. The instability manifests as a spiral density wave (a hot point) rotating at  $10-30 \\%$ of the Keplerian frequency at the inner edge of the disk, producing the modulation detected as a LFQPO. One would then expect the frequency $\\nu$ of the QPO to vary as $\\nu \\propto r^{-3/2}$, r being the inner radius of the disk. We have shown \\cite{Varniere02} that due to general relativistic effects this relation is modified whenever the disk is close to its last stable orbit (LSO). This could explain the behavior we observed in the case of \\J1655, where the relation between the QPO frequency and the disk inner radius was inverted compared to \\G1915 \\cite{Rodriguez02a}, whereas the energy dependence of the QPO amplitude in \\G1915 may suggest the presence of a hot point in the disk \\cite{Rodriguez02b}. ",
        "conclusions": "We propose that theoretical predictions of the AEI model match many observational constraints, starting with the non trivial evolution of the QPO frequency vs. the (soft) flux, and by extension the disk radius. We show that a hot point rotating in the disk, which is an expected signature of the AEI, is compatible both with the frequency evolution of the flux modulation and the energy spectrum of the QPO."
    },
    "0207/astro-ph0207013_arXiv.txt": {
        "abstract": "The mechanism of production of a large number of universes is considered. It is shown that universes with parameters suitable for creation of life are necessarily produced as a result of quantum fluctuations. Fractal structures are formed provided fluctuations take place near a maximum of the potential. Several ways of formation of similar fractal structures within our universe are discussed. Theoretical predictions are compared with observational data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207069_arXiv.txt": {
        "abstract": "The nearby galaxy M33 hosts the most luminous steady X-ray source in the Local Group. The high spatial resolution of {\\it Chandra} allows us to confirm that this ultra-luminous X-ray source is within the nucleus and rule out at the 4.6$\\sigma$ level a previously proposed possible counterpart located 1\\arcsec\\ away. The X-ray spectrum is well fitted by a disc blackbody with $kT_{\\rm in}$=1.18$\\pm$0.02~keV and $R_{\\rm in}(\\cos\\theta)^{0.5}$=57$\\pm$3~km, consistent with earlier results from {\\it ASCA} and {\\it BeppoSAX}. The source flux is steady between 1 and $10^4$~s ($<$3\\% rms variability, 0.5-10 keV) with a 0.5-10~keV luminosity of $1.5\\times10^{39}$~erg~s$^{-1}$.  The X-ray properties and an association with a radio source are reminiscent of the galactic microquasar GRS~1915+105. ",
        "introduction": "The nearby galaxy M33 ($d\\approx 795$~kpc; \\citealt{vdb}) hosts the most luminous steady X-ray source in the Local Group. M33 X-8 dominates the X-ray emission from the galaxy with a 1-10~keV luminosity of about 1.2$\\cdot 10^{39}$~erg~s$^{-1}$, ten times more than that of the other X-ray sources in the field. The ultra-luminous X-ray source (ULX) has been detected at this level since its discovery by {\\it Einstein}. \\citep{long,markert,gottwald,peres,takano,lc96,parmar}.  Observations obtained with the {\\it ROSAT} High Resolution Imager (HRI) showed that X-8 is point-like and coincident with the optical nucleus of M33 at the HRI 5\\arcsec{} resolution \\citep{schulman}. These also revealed a 106 day $\\sim$20\\% modulation of the X-ray flux from M33 X-8 \\citep{106}.  The nature of the X-ray emitting object is rather puzzling. The measured velocity dispersion of the nucleus places an upper limit of 1500~M$_\\odot$ on the mass of a putative massive central black hole \\citep{gebhardt}.  A low luminosity active galactic nucleus (AGN) is therefore unlikely. Indeed, the UV/optical spectrum of the nucleus shows no signs of AGN-type activity and is best modelled by a combination of two starbursts at 40 Myrs and 1 Gyrs (\\citealt{lcd} and references therein).  The 106 day modulation suggests a single object is responsible for most of the X-ray emission, perhaps a persistent X-ray binary with a 10 M$_\\odot$ black hole radiating at the Eddington luminosity. {\\em Hubble Space Telescope} observations place a limit of 18\\% on the contribution of a point source to the total far-UV flux from the nucleus \\citep{dubus}.  The implied ratio $L_\\rmn{opt}$/$L_\\rmn{X}\\sim 0.05$ would be in line with expectations from a low-mass X-ray binary. The counterpart would then be mostly lost in the glare of the nucleus. However, there remains a distinct possibility that M33 X-8 is associated with a UV bright star located 1\\arcsec{} to the NNW \\citep{dubus}. The star has a O9III spectrum with no outstanding feature linking it to the powerful X-ray source but cannot be ruled out at the {\\it ROSAT} HRI resolution \\citep{lcd}.  Most studies have assumed M33 X-8 is associated with the nucleus and used this to register the X-ray images. {\\it Chandra} offers the opportunity to locate M33 X-8 to unprecedented accuracy and settle whether this source is within the nucleus or not. Three observations of M33 are available in the {\\it Chandra} archive (Table~1). We used the first two to accurately position M33 X-8 with respect to {\\it HST} Planetary Camera (PC) images of the nucleus (\\S2). The third, non piled-up observation is used to study the X-ray spectrum and timing behaviour of M33 X-8 (\\S3 and \\S4). Our results are discussed in the last section.  \\begin{table} \\caption{{\\em Chandra} Observations of M33} \\begin{tabular}{@{}lllllll@{}} Det. & ObsId & Seq.\\# & Date & Exp. & offset\\\\ \\\\ ACIS-I & 1730 & 600145 & 7/12/2000 & 49.4~ks &  0\\arcmin   \\\\ ACIS-S & 786  & 600089 & 8/30/2000 & 46.2~ks &  0\\arcmin \\\\ ACIS-S & 787  & 600090 & 1/11/2000 & 9.2~ks  &  7.7\\arcmin\\\\ \\end{tabular} \\medskip  Det. is the detector used; exp. is the total integration time; offset is the distance of M33 X-8 from the telescope axis. The CCD frame time was 3.2~s for the first two observations and 1.0~s for the last, offset observation. \\end{table} ",
        "conclusions": "The high luminosity of M33 X-8 is not sufficient to rule out a neutron star X-ray binary since magnetospheric accretion can result in super-Eddington luminosities. This was observed e.g. in the high mass X-ray binaries (HMXBs) SMC X-1 or A0535-66. A HMXB would have been an attractive possibility should M33 X-8 have been associated with the nearby 09III star. The {\\it Chandra} localisation of X-8 rules this out (\\S2). X-8 could still be a HMXB hiding in the nucleus but this would be surprising since the compactness and spectra of the nucleus make it unlikely that it hosts $>10 M_\\odot$ stars \\citep{lcd}; a HMXB evolves on the uncomfortably short thermal timescale of the donor star ($<10^5$~yrs for $M_2 >10 M_\\odot$); the X-ray spectrum of X-8 is unlike that of typical HMXBs which have flat power laws in the 0.5-10~keV range; and the power spectrum shows no evidence for pulsations.  On the other hand, the X-ray spectrum and absence of short-term X-ray variability below 1~Hz are highly reminiscent of the high or very high states of galactic black hole candidates \\citep{tl}. In both states the spectrum is dominated by a disc blackbody with a weak contribution from a steep power law. The very high state shows 3-10~Hz QPOs with a few \\% rms superposed on 1-10\\% rms flat top noise in the 1-10~Hz range. There is little variability in the high state. GRS~1915+105, a microquasar in our Galaxy with a $\\sim 14$~M$_\\odot$ black hole accreting from a $\\sim 1.2$~M$_\\odot$ K giant star \\citep{greiner}, also has such soft, low variability states \\citep{belloni}. The optical emission from such a system would be lost in the nucleus of M33.  Since its discovery in 1992, GRS~1915+105 has been steadily detected with X-ray luminosities around or in excess of $10^{39}$~erg~s$^{-1}$.  The analogy with GRS~1915+105 is even more attractive considering the 0.6~mJy (respectively 0.2~mJy) radio source detected by the VLA at 20~cm \\citep[6~cm;][]{gordon99} which we show to be consistent with the ultra-luminous X-ray source (\\S2.4). GRS~1915+105 has been known to show steady radio emission around 30-50~mJy at 2~cm, framed by powerful ejection episodes where the emission from the relativistic jets becomes $\\ga$ 500~mJy \\citep{fender}. The radio source in M33 is more powerful than that; but a small change in the beaming angle of the jet ($\\approx 70^\\circ$ from the line-of-sight for GRS~1915+105) could easily increase the observed radio flux. Follow-up observations of the radio source are needed to test this hypothesis."
    },
    "0207/astro-ph0207543_arXiv.txt": {
        "abstract": "% The flux of cosmic ray hadrons at the atmospheric depth of 820 g/cm$^2$ has been measured by means of the EAS-TOP hadron calorimeter (Campo Imperatore, National Gran Sasso Laboratories, 2005 m a.s.l.).\\\\ The hadron spectrum is well described by a single power law : \\\\ $S_{h}(E_{h}) = (2.25 \\pm 0.21 \\pm 0.34^{sys}) \\times 10^{-7} (\\frac{E_{h}}{1000})^{(-2.79 \\pm 0.05)}$ m$^{-2}$s$^{-1}$sr$^{-1}$GeV$^{-1}$\\\\ over the energy range 30 GeV $\\div$ 30 TeV. The procedure and the accuracy of the measurement are discussed. \\\\ \\vskip .8 cm \\small (*) Corresponding author. \\normalsize  The primary proton spectrum is derived from the data by using the CORSIKA/QGSJET code to compute the local hadron flux as a function of the primary proton spectrum and to calculate and subtract the heavy nuclei contribution (basing on direct measurements). Over a wide energy range $E_{0}= 0.5 \\div 50$ TeV its best fit is given by a single power law :\\\\ $S(E_{0}) = (9.8 \\pm 1.1 \\pm 1.6^{sys}) \\times 10^{-5} (\\frac{E_{0}}{1000})^{(-2.80 \\pm 0.06)}$ m$^{-2}$ s$^{-1}$ sr$^{-1}$ GeV$^{-1}$.\\\\ The validity of the CORSIKA/QGSJET code for such application has been checked using the EAS-TOP and KASCADE experimental data by reproducing the ratio of the measured hadron fluxes at the two experimental depths (820 and 1030 g cm$^{-2}$ respectively) at better than $10 \\%$ in the considered energy range. ",
        "introduction": "The spectrum of hadrons detected at different atmospheric altitudes retains significant information about the energy/nucleon spectrum of primary cosmic rays, which is dominated by the lightest component, i.e. the proton one. Its measurement has been carried on in the past, both at sea level \\cite{dig,bro,bar,fic,ash,cow} and at mountain altitude \\cite{ame,can,ren,sio,ada}, using different experimental techniques, like emulsion chambers, magnetic spectrometers and calorimeters.\\\\ The knowledge of the primary proton spectrum is of main relevance for the understanding of the cosmic rays acceleration mechanisms and of the propagation processes in the Galaxy. Moreover, the proton component is mainly responsible for the uncorrelated particle production in the atmosphere: any uncertainties on the proton spectrum reflect in an uncertainty in the calculation of the secondary particle fluxes ($\\pi$ and $K$) and thus, for example, on the knowledge of the atmospheric muon and neutrino fluxes. A precise knowledge of such spectra is of particular importance to interpret the observational data from muon and neutrino detectors deep underground \\cite{cas}.\\\\ The measurement of the primary proton spectrum has been performed by means of satellite and balloon borne experiments \\cite{boe,men,bel,san,web,alc,swo,buc,seo,rya,zat,iva,asa,apa} and indirectly derived by using ground based detectors \\cite{ino,ame1,sak,aha}. In the region of tens of TeV, however, direct measurements lack statistics and moreover their energy determinations are not calorimetric and depend on the interaction parameters and their fluctuations. The data inferred from hadron measurements at ground level can therefore provide significant new information.\\\\ On the other side, the derivation of the information on the primary cosmic ray spectrum from hadron measurements, as well as the comparison of the results from different experiments, relies on the use of simulation tools describing the interaction and propagation of primary cosmic rays in the atmosphere. The response of such hadron interaction models has therefore to be verified, especially considering that the recorded hadrons are the results of large fluctuations with respect to the average behavior. \\\\ The EAS-TOP Extensive Air Shower array was located at Campo Imperatore, 2005 m a.s.l., above the underground Gran Sasso Laboratories, with the aim of studying the cosmic ray spectrum in the energy range $10^{13} \\div 10^{16}$ eV through the detection of the different air shower components.\\\\ In this paper, we present and discuss the results obtained in the study of the uncorrelated hadrons by means of the calorimeter of EAS-TOP, namely: \\\\ a) the measurement of the hadron flux in the energy range 30 GeV $\\div$ 30 TeV;\\\\ b) the derivation of the primary proton energy spectrum in the range 0.5 $\\div$ 50 TeV; \\\\ c) the check of the propagation and interaction code (CORSIKA/QGSJET) used for the interpretation of the data. ",
        "conclusions": "The hadron flux has been measured over a wide energy range (30 GeV$\\div$30 TeV) by means of the EAS-TOP hadron calorimeter at the atmospheric depth of 820 g/cm$^{2}$. The spectrum is well described by a single power law in the whole range : \\\\ $S(E_{h}) = (2.25 \\pm 0.21 \\pm 0.34^{sys}) \\times 10^{-7} (\\frac{E _{h}}{1000})^{(-2.79 \\pm 0.05)}$ m$^{-2}$ s$^{-1}$ sr$^{-1}$ GeV$^{-1}$.\\\\ Taking into account the contamination from heavier nuclei, on the basis of direct measurements, the primary proton spectrum is obtained between 0.5 and 50 TeV and is found to be compatible with a single slope power law: \\\\ $S(E_{0}) = (9.8 \\pm 1.1 \\pm 1.6^{sys}) \\times 10^{-5} (\\frac{E_{0}}{1000})^{(-2.80 \\pm 0.06)}$ m$^{-2}$ s$^{-1}$ sr$^{-1}$ GeV$^{-1}$.\\\\ A systematic uncertainty of about 7 $\\%$ due to the uncertainty in the helium flux is included. The data match very well with the direct measurements over a wide energy range, usually not available to a single experiment, where direct measurements become statistically poor.\\\\ The reliability of the CORSIKA/QGSJET interaction and propagation code, which is used to propagate the hadrons in the atmosphere and to compute the heavy nuclei contribution, is directly checked in this energy range by comparison with accelerator data and, concerning the direct application to the present measurement, through its capability to reproduce the ratio of hadron fluxes as measured at two different atmospheric depths by EAS-TOP and KASCADE, at 820 and 1030 g/cm$^{2}$ respectively."
    },
    "0207/astro-ph0207296_arXiv.txt": {
        "abstract": "We present the    results of an \\xmm  observation of the radio-quiet X--ray pulsar \\ee\\ located at the center of the shell-like supernova remnant  \\gg\\ . The X--ray spectrum is characterized by the presence of two phase-dependent absorption lines at energies $\\sim$0.7 keV and $\\sim$1.4 keV. Moreover, these broad spectral features have significant substructure, suggesting that they are due to the blending of several  narrower lines. We interpret such features as evidence for an atmosphere containing metals and a magnetic field value of a few 10$^{12}$ G, consistent with the observed spin-down rate $\\pdot$=(1.98$\\pm$0.83)$\\times$10$^{-14}$ s s$^{-1}$. Since \\ee\\ is the only X--ray emitting pulsar showing evidence of such features, we tentatively link them to the unique combination of  age and energetics that characterize this object. We suggest that a young age and a low level of magnetospheric activity are  favorable conditions for the detection of atomic spectral features from $Z>1$ elements in neutron star atmospheres, which would be either blanketed by a thin layer of accreted hydrogen in older objects or masked by non-thermal processes in young energetic pulsars. ",
        "introduction": "The thermal emission from the surface of a neutron star traces  the star's cooling history. Its study can thus provide invaluable information on the poorly known equation of state of matter at super-nuclear densities and on physical processes in strong magnetic fields. Satellite observations carried out in the last decade have clearly shown the thermal origin of the soft X--ray emission ($\\sim$0.1-3 keV) from a handful of middle-aged radio pulsars (Becker \\& Tr\\\"{u}mper 1997) as well as from several radio-quiet neutron stars (Caraveo, Bignami \\& Tr\\\"{u}mper 1996; Treves et al. 2000).  It is expected   that the presence of an atmosphere on the   neutron star surface distort the emerging radiation  by altering the blackbody energy distribution and  introducing absorption  features (see, e.g., Zavlin \\&  Pavlov 2002 for a recent review). Fits with atmospheric models generally yield emitting regions compatible with standard neutron star dimensions, thus providing indirect evidence for the presence of an atmosphere. However, until recently no convincing evidence for the absorption features predicted by these models was found, thus leaving substantial uncertainty on the atmospheric composition and magnetic field.   Here we report on an \\xmm observation of the radio-quiet neutron star  \\ee . Previous X--ray, optical and radio observations (Bignami, Caraveo \\& Mereghetti 1992; Mereghetti, Bignami \\& Caraveo 1996; Vasisht et al. 1997) strongly suggested a  neutron star nature for  this source, located close to the geometrical center of the shell-like supernova remnant  \\gg  (Roger et al. 1988). This was confirmed by the discovery  of fast  X--ray  pulsations with period $P$=0.424 s (Zavlin et al. 2000). A  subsequent    measurement (at the $\\lsim2\\sigma$ level) of a positive  period derivative $\\pdot$ = (2.0$^{+1.1}_{-1.3}$) $\\times$ 10$^{-14}$ s s$^{-1}$  (Pavlov et al. 2002a) results in a large discrepancy between the pulsar's characteristic age $\\tau_{c}\\equiv\\frac{P}{2\\pdot}=200-900$ kyrs and  the age  of  7 kyrs (with a factor 3 uncertainty) estimated for the associated SNR (Roger et al. 1988).  Our  \\xmm data show the presence  of broad absorption features at $\\sim$0.7 and $\\sim$1.4 keV in the spectrum of  \\ee\\ . Such features have been independently discovered in data from the \\cha satellite (Sanwal et al. 2002). The high throughput of the \\xmm telescope, coupled to the good spectral and timing resolution of the European Photon Imaging Camera (EPIC) instrument, allow us to show  that the lines have a significant substructure, which varies with the phase of the pulsar. ",
        "conclusions": "The \\xmm observations reported here show unambiguously that the absorption  features independently discovered with \\cha (Sanwal et al. 2002) in the X--ray spectrum of \\ee are phase-dependent and have significant sub-structure. This supports an interpretation in terms of atomic transitions in regions with different temperature and magnetic field on the neutron star surface. A detailed analysis of phase-resolved spectra with higher statistical quality will undoubtedly provide important information on the surface composition and other  neutron star parameters.  The  period value reported here is  identical, within the uncertainties, to that obtained two weeks later (Pavlov et al. 2002a), but our smaller error reduces the uncertainty on the pulsar spin-down rate $\\pdot$=(1.98$\\pm$0.83)$\\times$10$^{-14}$ s s$^{-1}$. Since the association with the relatively young SNR remnant is extremely likely, we consider the high characteristic age of \\ee\\ as evidence that this pulsar was born with a long period.  We suggest that the combination of a young age and a low level of magnetospheric activity are  favorable conditions for the detection of atomic spectral features from Z$>$1 elements, which would be either blanketed by a thin layer of accreted hydrogen in older objects or masked by non-thermal processes in young energetic pulsars. If this scenario is correct,  the most  promising candidates for line detection are  the   central X--ray sources in young SNRs which do not show synchrotron nebulae, like Cas A (Mereghetti, Tiengo \\& Israel 2002), Puppis A (Zavlin, Tr\\\"{u}mper \\& Pavlov 1999),  and  G 266.2--1.2 (Pavlov et al. 2001b)."
    },
    "0207/astro-ph0207319_arXiv.txt": {
        "abstract": "Many models of Gamma Ray Bursts (GRBs) invoke a central engine consisting of a black hole of a few solar masses accreting matter from a disk at a rate of a fraction to a few solar masses per second. Popham et al. and Narayan et al. have shown that, for $\\Mdot \\approxgt 0.1 \\Msun$s$^{-1}$, accretion proceeds via neutrino cooling and neutrinos can carry away a significant amount of energy from the inner regions of the disks. We improve on these calculations by including a simple prescription for neutrino transfer and neutrino opacities in such regions. We find that the flows become optically thick to neutrinos inside a radius $R \\sim 6-40 R_{\\rm s}$ for $\\Mdot$ in the range of $0.1 -10 \\Msun$s$^{-1}$, where $R_{\\rm s}$ is the black hole Schwarzchild radius. Most of the neutrino emission comes from outside this region and, the neutrino luminosity stays roughly constant at a value $L_{\\nu} \\sim 10^{53} \\ergps$. We show that, for $\\Mdot \\approxgt 1\\Msun$s$^{-1}$, neutrinos are sufficiently trapped that energy advection becomes the dominant cooling mechanism in the flow. These results imply that $\\nu\\bar{\\nu}$ annihilation in hyperaccreting black holes is an inefficient mechanism for liberating large amounts of energy. Extraction of rotational energy by magnetic processes remains the most viable mechanism. ",
        "introduction": "Most popular models of Gamma Ray Bursts (GRBs) invoke a binary merger or a collapse involving compact objects. In particular, these include mergers of double neutron star binaries (Eichler et al. 1989; Narayan, Paczynski \\& Piran 1992; Ruffert \\& Janka 1999), mergers of a neutron star with a black hole (Paczynski 1991; Narayan et al. 1992; Ruffert \\& Janka 2001 and references therein), of helium star with a black hole (Fryer \\& Woosley 1998), ``collapsars'' or ``failed supernovae'' (Woosley 1993; Paczynski, 1998; MacFadyen \\& Woosley 1999; MacFadyen, Woosley \\& Heger 2001) and ``supranovae'' (Vietri \\& Stella 1998). All of the above scenarios lead to the formation of a black hole with a debris torus or disk around it. (The only exceptions are models in which the GRB energy is provided by the magnetic and rotational energy of the newly formed neutron star; e.g.; Usov 1992). In order to understand how the extraordinary amount of energy characteristic of GRBs can be extracted, we are motivated to further examine the properties of such compact and massive disks around black holes.  Accretion models in the context of GRBs have been recently discussed by Popham, Woosley \\& Fryer (1999; hereafter PWF), Narayan, Piran \\& Kumar (2001; hereafter NPK) and Kohri \\& Mineshige (2002). The typical mass accretion rates in GRB models are extremely high, of the order of a fraction of solar mass up to a few solar masses per second. Under such conditions, the gas photon opacities are also very high and radiation becomes trapped (see e.g.; Katz 1977; Begelman 1978; Abramowicz et al. 1988). However, at sufficiently high mass accretion rates, although energy advection remains important in the outer parts, the disk becomes dense and hot enough in the inner regions to cool via neutrino emission. For this reason PWF named these disks neutrino-dominated accretion flows (NDAFs). This regime is of particular interest, because neutrinos can, in principle (see e.g. NPK; Ruffert \\& Janka 1999), tap the thermal energy of the disk produced by viscous dissipation and liberate large amounts of its binding energy (via the $\\nu\\bar\\nu \\rightarrow e^+e^-$ process in regions of low baryon density).  For this mechanism to be efficient, though, the neutrinos must escape before being advected into the black hole.  In this paper we investigate the effects of neutrino transport within the context of NDAFs. By using a simple prescription to account for neutrino scattering and absorptive opacities, we find that, for accretion rates $\\Mdot \\approxgt 1 \\Msun \\ps$, the gas becomes increasingly more opaque and neutrinos become trapped. As a result, energy advection becomes the dominant cooling mechanism in the inner regions of the flows. We show that, as $\\Mdot$ increases, the disk emitting surface moves further out in radius and the neutrino luminosity of the flow remains nearly constant. We show that the accretion flow luminosity plateaus as it approaches the neutrino Eddington limit of the inner disk and as the neutrino cooling efficiency decreases.  NPK and PWF also noted the importance of neutrino opacity at $\\Mdot \\approxgt 1 \\Msun$s$^{-1}$  but did not allow for neutrino transport effects in their models. Our work therefore complements the earlier studies of NDAFs. A detailed treatment of neutrino transfer was included in the numerical simulations carried out by Ruffert \\& Janka (1999), where accretion tori are formed as a result of neutron star merging. Consistent with our findings, these authors also show that opacities can be high in such tori. However, their results are specific to the parameter space covered by the merger model.  In \\S 2 we identify the dominant neutrino opacity sources and outline the basic equations we use and the approximations we make to the neutrino transfer problem. In \\S 3 we present our numerical results. We delineate the regions of parameter space in accretion rate and radius where the flow becomes optically thick to neutrino emission. In \\S 4 we discuss the stability of the accretion flow.  Finally, in \\S 5, we compute the flow luminosity of our model, compare it with the derived neutrino Eddington limit and discuss various implications of our results for GRB energetics. ",
        "conclusions": "The relevance of accretion processes in the central engine of GRBs was highlighted by Narayan et al. (1992). PWF carried out the first detailed study of accretion around black holes with ultra-high $\\dot{m}$ (of order a solar mass per second), and identified a new mode of accretion which they named neutrino-dominated accretion flow (NDAF). In a subsequent study, NPK noted that accretion flows with such high $\\dot{m}$ are highly advection dominated at large radii ($r\\ga 100$). They suggested that the flow at these outer radii may take the form of a convection-dominated accretion flow (CDAF; Narayan, Igumenshchev \\& Abramowicz 2000; Quataert \\& Gruzinov 2000) or a related kind of flow (e.g., Blandford \\& Begelman 1999). As a result, only a small fraction of the available mass accretes on the black hole, the bulk of the gas being ejected from the system. Close to the black hole ($r\\la 100$), however, NPK confirmed PWF's result, that a cooling-dominated NDAF should be present.  PWF and NPK made the simple assumption that the accreting gas is optically thin to its own neutrino emission. We have improved on this by including the effects of neutrino transfer (via a simple prescription) and the neutrino opacities self-consistently in the model. With these improvements, we find that the central regions of NDAFs are typically opaque, so that the neutrinos emitted by the accreting gas are largely trapped. The optically-thick region extends out to $r \\sim 4-5$ at $\\mdot \\sim 0.1$ to $r\\sim 30-40$ at $\\mdot \\sim 10$ (Figure 3). We have shown that, above $\\mdot \\approxgt 1$, the majority of the energy liberated by viscous dissipation is advected with the flow instead of being emitted in the form of neutrinos, and energy advection is much more important in the inner regions of the accretion flows than previously realized (Figure~3).  In order to further explore the implications of these results we now calculate the neutrino luminosity from our model and the fraction of the binding energy carried out to infinity by neutrinos. We also derive the neutrino Eddington luminosity for such flows.  Figure 6 shows the neutrino luminosity from the accretion flow, $L_{\\nu} = \\int_{r_{\\rm min}}^{r_{\\rm max}} 2 \\pi q_{\\nu}^{-}rdr$, (in units of $10^{51} \\ergps$; solid line; upper panel). Also, using the opacities discussed in \\S 2.2 we derive the neutrino Eddington luminosity of the flow: \\beq L_{\\rm Edd,\\nu} = \\frac{4\\pi GM c}{\\kappa_{\\nu}} \\sim 9 \\times 10^{53} m_3 T_{11}^{-2} \\ergps , \\eeq where, for illustration the numerical value given on the right uses only the two dominant components of the opacity: neutron and proton absorption and scattering (for electron neutrinos this gives $\\kappa_{\\nu_e} = \\kappa_{\\nu_e, a} + \\kappa_{\\nu_e, s} = 5.5 \\times 10^{-17} T_{11}^2 \\cmsq$g$^{-1}$; Equations 8 and 11; note that for the calculation shown in Figure 6 all opacity terms are used). Because the neutrino cross sections are a function of neutrino energies, the neutrino Eddington luminosity is a function of both black hole mass and gas temperature. In Figure 6, we show the neutrino Eddington luminosity calculated for $T_{11} = T_{11} (r=3)$, the temperature at the inner radius of the flow; at larger radii the Eddington luminosity increases.  In the middle panel of Figure 6 we also show the neutrino radiative efficiency, defined, as in standard accretion theory, by $\\eta_{\\nu} = L_{\\nu}/\\Mdot c^2$.  The neutrino luminosity from the flow is $L_{\\nu} \\sim 10^{51} \\ergps$ at $\\mdot \\sim 0.01$  and increases almost linearly with $\\mdot$ up to $\\mdot \\sim 0.1$ (Figure 6). Between $0.1 \\approxlt \\mdot \\approxlt 1$, the luminosity flattens off significantly, ranging within $6$ -- $8 \\times 10^{52} \\ergps$. Above $\\mdot \\sim 1$ the neutrino luminosity stays virtually constant at $L_{\\nu} \\sim 8 \\times 10^{52} \\ergps$. In contrast, in the PWF solution, which assumes optically thin neutrino emission, the neutrino luminosity ranges from $L_{\\nu} \\sim 10^{51-52} \\ergps$ for $\\mdot = 0.1$ to up to $L_{\\nu} \\sim 10^{54} \\ergps$ for $\\mdot=10$ (see their Table 3). Thus, one important consequence of the addition of neutrino opacities is that the flow luminosity varies only by a factor of a few over the range of accretion rates appropriate for popular models of GRB progenitors.  The luminosity of the flow is almost constant because the efficiency with which energy is transported out of the flow by neutrinos, $\\eta_{\\nu}$ (bottom panel, Figure 6) decreases with increasing $\\mdot$ (or equivalently, as shown in Figure 5, decreases as neutrinos become trapped and energy advection becomes more dominant at larger $\\mdot$). We show that $\\eta \\sim 0.1$ (consistent with the efficiency expected from a geometrically thin, cooling dominated Newtonian accretion disk) up to $\\mdot \\sim 0.1$ and decreases (almost linearly) with increasing $\\dot{m}$ (reaching $\\eta \\sim $a few $ 10^{-3}$ at $\\mdot =10$). The fraction of energy transported away by neutrinos decreases as the inner regions of the flow become neutrino opaque (see Figure 3).  Figure 6 also shows the neutrino Eddington luminosity evaluated at the inner edge of the disk (dashed line). In the inner regions of the flow, the temperature is the highest, and consequently the Eddington luminosity the lowest (Eq.~20), so that $L_{\\nu} \\sim L_{{\\rm Edd},\\nu}$ for $\\mdot \\approxgt 1$. Momentum deposition in the inner regions of the flow by neutrinos may therefore be quite effective as a mechanism responsible for ejection.  We now examine the implications of our results for GRBs. A number of authors, including Eichler et al. (1989; see also Narayan et al. 1992) suggested that $\\nu\\bar\\nu$ annihilation around merging compact binaries might produce a relativistic $e^{+}e^{-}$ fireball with sufficient energy to power a GRB. Detailed numerical simulations by Janka et al. (1999) show that this process might conceivably power short GRBs (those with durations under a couple of seconds or so), provided that there is a modest level of beaming. However, we have shown that neutrino advection is important in most of the parameter space of these models, which may have a significant impact on the efficiency of this process. We estimate the luminosity due to $\\nu\\bar\\nu$ annihilation along $z$-axis above the disk (using an improved version of Eq.~10 in Ruffert et al. (1997) provided by Thomas H. Janka, private communication) to be: \\beq \\L_{\\nu \\bar{\\nu}} \\sim 6 \\times 10^{-35}  \\frac{2\\;\\langle E_{\\nu} \\rangle\\;L_{\\nu}^2}{\\pi c^2\\left[1-({R_{\\rm min}}/{R_{\\rm surf})^2}\\right]^2} \\;\\;\\frac{1}{R_{\\rm surf}} \\int_{0}^{\\infty} d\\epsilon\\;\\Phi\\; (\\cos \\theta_{\\rm min} - \\cos\\theta_{\\rm surf})^2\\;\\;\\ergps, \\eeq where the constant in front takes into account the neutrino cross sections. $L_{\\nu}$ is given above and shown in Figure 6 and $R_{\\rm surf}$ and $\\langle E_{\\nu} \\rangle $ are evaluated at the neutrinosphere, defined as the surface at which the emergent neutrinos originate. Hence, $\\langle E_{\\nu} \\rangle$ is the average energy of the escaping neutrinos given by $\\langle E_{\\nu} \\rangle = 3.7\\, k T_{11, {\\rm surf}}$ (Eq.~10) with $T_{11, {\\rm surf}} = T_{11}(r_{\\rm surf}=r(\\tau_{\\nu}=2/3))$ and $\\tau_{\\nu} = \\tau_{a,\\nu} + \\tau_{s,\\nu}$.  If the spectrum is a blackbody, the radius of this surface is located near and above the layer of optical depth $= 2/3$ (this is because the opacity is mostly dominated by absorption processes). The contour of $\\tau_{\\nu} =2/3$ and the radius of the neutrinosphere as a function of $\\mdot$ is shown in Figure 3. We find that the radius of the neutrinosphere \\footnote{To derive the analytical scalings we solve Eqs.~(13) and (14) assuming $P=P_{\\rm gas}$ and $q^{+} = q_{adv}^-$. We have shown that is a good approximation within the optically thick region.} to be given by, $r_{\\rm surf} = r_{\\tau=2/3} \\sim 17 \\mdot^{2/5}$ in the range $0.1 \\approxlt \\mdot \\approxlt 10$. At this radius the temperature of the neutrinosphere is roughly $T_{11,{\\rm surf}} \\sim 0.2 \\mdot^{-3/20}$ which is fairly insensitive to $\\mdot$ in the range from $0.1 \\approxlt \\mdot \\approxlt 10$). The integral part in equation (21) is a geometrical factor, where $\\Phi = 3/4[1-2\\langle \\mu \\rangle^2 +\\langle \\mu^2 \\rangle^2 +1/2 (1-\\langle \\mu^2 \\rangle)^2]$, where $\\langle \\mu^n\\rangle = \\int_{\\mu_{\\rm surf}}^{\\mu_{\\rm min}} d\\mu \\mu^n / \\int_{\\mu_{\\rm surf}}^{\\mu_{\\rm min}} d\\mu$ with $\\mu_{\\rm surf} = \\cos\\theta_{\\rm surf} = 1/ \\sqrt{1 + (r_{\\rm surf}^2/z^2)}$, $\\mu_{\\rm min} = \\cos\\theta_{\\rm min} = 1/ \\sqrt{1 + (r_{\\rm min}^2/z^2)}$ and $\\epsilon = z/r_{\\rm surf}$.  The bottom panel of Figure 6 shows the neutrino annihilation luminosity, $L_{\\nu \\bar{\\nu}}$ in units of $10^{51} \\ergps$.  $L_{\\nu \\bar{\\nu}}$ increases up to its maximum value of $\\sim 10^{50} \\ergps$ at $\\mdot \\approxgt 1$ and slightly decreases for larger $\\mdot$. Our estimate of $L_{\\nu,\\bar{\\nu}}$ at $\\mdot =1$ agrees well with values estimated in the simulations of Ruffert \\& Janka (1999) and PWF (although our results are less accurate than those of PWF for $\\mdot \\approxlt 0.1$). In accordance with PWF we find that energetic events can only be achieved for $\\mdot > 0.1$ (below which the neutrino annihilation efficiency decreases very sharply; see also their Table 3); but in contrast with their results we do not find that increasingly more energetic events can be achieved for larger accretion rates. Our calculations imply that the efficiency of $\\nu\\bar\\nu$ annihilation remains constant (or slightly decreases) for $\\mdot \\approxgt 1$. This is expected, as we have found that neutrinos are increasingly more trapped in the disk. (Our estimate may be uncertain by a factor $\\sim 2$; this is mostly due to the fact that our solution gives only height-averaged quantities whereas the vertical disk stratification may be important in this calculation). Note also that recent results from hydrodynamical calculations carried out by Lee \\& Ramirez-Ruiz (2002) show that $\\nu\\bar\\nu$ annihilation can only produce bursts from impulsive energy inputs, as the annihilation luminosity scales as $t^{-5/2}$. This further restricts the importance of this process to a small fraction of bursts.  Energy extraction from the disk is also possible by MHD processes. Such processes are broadly based on the expectation that the differential rotation of the disk will amplify pre--existing magnetic fields, until they approach equipartition with the gas kinetic energy. Proposed mechanisms include Parker instabilities in the disk leading to reconnection, relativistic flares and winds (Narayan et al. 1992; Meszaros \\& Rees 1997) or the Blandford-Znajek mechanism (Blandford \\& Znajek 1977; hereafter BZ). Of these, perhaps, the Blandford-Znajek efficiency is the easiest to estimate (at least roughly). We follow the common assumption that the magnetic field in the disk will rise to some fraction of its equipartition value $B^2/8 \\pi \\sim \\rho c^2_s$. Typical values of $\\rho c^2_s$ are $10^{30-32} \\ergpcmq$ for $0.1 < \\mdot < 10$ (see Figure 2), implying a field strength of $\\sim 10^{15 -16}$ G if we make the conservative assumption that $B$ is only $10 \\%$ of its equipartition value. Consistent with earlier work (PWF; Ruffert \\& Janka 1999), the BZ jet luminosity at $\\mdot=10$ is then \\beq L_{BZ} = \\left(\\frac{B^2}{4\\pi}\\right) \\pi R_{h}^2 a^2 c \\approxgt 10^{52} a^2 \\left(\\frac{B}{10^{16} \\G} \\right)^2 \\left(\\frac{M}{3\\;\\Msun}\\right)^2 \\;\\;\\; \\ergps\\;, \\eeq where $R_{h}= 2 G M/c^2$ is the black hole radius and $a = R_{h} \\Omega_{h}/c$ is the dimensionless black hole spin parameter ($0 <a <1$). $L_{BZ}$ is about 2 orders of magnitude larger than the neutrino annihilation rate, $L_{\\nu \\bar\\nu}$, at $\\mdot \\sim 10$. For larger equipartition fractions, the ratio $L_{BZ}/L_{\\nu \\bar\\nu}$ becomes large even at $\\mdot \\sim 0.1$ or $1$. It is obvious that the energy liberated by this mechanism is simply proportional to the internal energy density of the disk and hence is not affected by advection. Indeed, Livio, Ogilvie \\& Pringle (1999) have argued that the BZ mechanism might be most relevant for advection dominated flows (see also Meier 2001). Numerical simulations such as those by MacFadyen \\& Woosley (1999) and MacFadyen, Woosley \\& Heger (2001) have shown that the BZ mechanism would be an efficient mechanism capable of liberating a large fraction of the black hole spin energy.  By similar electromagnetic considerations as those used to derive the BZ luminosity, Livio et al.(1999) estimate that the electromagnetic/wind power output from a disk is given by: \\beq L_{d} = \\left(\\frac{B_{pd}^2}{4\\pi}\\right) \\pi R_{d}^2 \\left(\\frac{R_d \\Omega_d}{c}\\right) c \\sim \\left(\\frac{B_{pd}}{B}\\right)^2 \\left(\\frac{R_{d}} {R_{h}}\\right)^{3/2} a^{-2} L_{BZ} \\sim 5 \\; a^{-2} L_{BZ} \\eeq where, following Livio et al., we take ${R_d \\Omega_d}/{c} \\sim (R_h/R_d)^{1/2}$. $B_{pd}$ is the poloidal field in the disk given, approximately, by $B_{pd} \\sim (H/R) B$ (see also Merloni \\& Fabian 2002). Because $H/R \\sim 1$ in the inner region of the disk we have$ B_{pd} \\sim B$. $R_{d}$ is a factor of a few times $R_{h}$ (above we take $R_{d}$ as the radius of the innermost stable orbit, at which the disk energy density is the highest). Since $a^2 <1$ and, $R_d$ is at least a few times $R_{h}$, it is evident that in a geometrically thick disk, the electromagnetic/wind luminosity from the disk can easily be larger than the BZ luminosity, hence $L_{d} >> L_{\\nu \\bar{\\nu}}$. Our calculations therefore indicate that, with increasing accretion rates, MHD processes become significantly more efficient at releasing energy than neutrino annihilation processes, hence they are probably the most viable mechanisms for energy extraction in these systems.  Although the stability analysis presented in \\S 4 shows that the accretion flows we have studied are intrinsically stable, we note that the viscous time scale (Fig. 3) can be as short as a small fraction of a second. Therefore, if the accretion disk is fed in a variable manner, e.g., via fallback material after a supernova explosion as in the collapsar model, we may expect variations in the accretion rate. This might explain the complex light curves of GRB in at least some of the scenarios discussed above (e.g., MacFadyen \\& Woosley, 1999)."
    },
    "0207/astro-ph0207633_arXiv.txt": {
        "abstract": "We predict the level of small-scale anisotropy in the cosmic microwave background (CMB) due to the Sunyaev--Zel'dovich (SZ) effect for the ensemble of cosmological models that are consistent with current measurements of large-scale CMB anisotropy. We argue that the recently reported detections of the small-scale (arcminutes) CMB anisotropy are only marginally consistent with being the SZ effect when cosmological models are calibrated to the existing primary CMB data on large scales. The discrepancy is at more than $2-2.5\\sigma$, and is mainly due to a lower $\\sigma_8\\la 0.8$ favored by the primary CMB and a higher $\\sigma_8\\ga 1$ favored by the SZ effect. A degeneracy between the optical depth to Thomson scattering and the CMB-derived value of $\\sigma_8$ suggests that the discrepancy is reduced if the universe was reionized very early, at redshift of $\\sim 25$. ",
        "introduction": "\\label{sec:intro}  The current generation of cosmic microwave background (CMB) experiments are producing a wealth of information on a wide range of angular scales. With the recent measurements by BIMA \\citep{dawson02} and CBI \\citep{mason02} complementing earlier measurements \\citep[see][for a recent compilation]{hu02}, CMB anisotropy has been measured over a range of multipoles of $l=2-6000$ (or angular scales of $2'-90^\\circ$).   At low multipoles ($l \\la 2000$) anisotropy is primarily generated at $z\\ga 1000$ except at very low multipoles ($l \\la 10$) where late-time decay of gravitational potential contributes significantly. At higher multipoles (smaller angular scales) low-redshift sources generate a significant amount of fluctuation power. At the observing frequencies of CBI and BIMA ($\\sim$ 30 GHz), the largest sources of low-redshift anisotropy are radio point sources and the thermal Sunyaev--Zel'dovich (SZ) effect. The latter is of cosmological interest and may be large enough to be detected, depending on cosmology \\citep[e.g.,][]{cole88}.   The reported detections of power at small angular scales ($l=2000-6000$) have argued that point-source contamination is not a problem, suggesting that the detected power could be due to the SZ effect \\citep{bond02,dawson02,komatsu02}. Since the number density and brightness of the sources contributing to the SZ fluctuations (i.e., hot gas in galaxy clusters at $z \\la 1$) depend on the background cosmology, the level of the SZ fluctuations depends on cosmological parameters. The SZ angular power spectrum is sensitive to the matter-fluctuation amplitude and the baryon density of the universe but relatively insensitive to the matter density of the universe \\citep{komatsu99} or other cosmological parameters \\citep{komatsu02}. By fitting the CBI and BIMA data to theoretical predictions, \\citet{komatsu02} have found a constraint on linear r.m.s. mass fluctuations within an $8~h^{-1}~{\\rm Mpc}$ sphere, $\\sigma_8$, as $\\sigma_8(\\Omega_{\\rm b}h/0.035)^{0.29}=1.04\\pm 0.12$ at the 95\\% confidence level.   On the other hand, observations of the primary CMB anisotropy and the large-scale structure (LSS) of the universe have already provided tight constraints on cosmological parameters \\citep[e.g.,][]{wang02}. By using these constraints we can predict how much SZ power ought to be seen at the CBI and BIMA multipole bands. By doing so we can see if an SZ interpretation of the small-scale fluctuations is consistent with cosmological models favored by CMB or LSS.  In this {\\em Letter} we estimate the level of the SZ angular power spectrum expected from cosmological models consistent with CMB data at $l < 2000$, and compare it with the CBI and BIMA data at $2000<l<10000$. We use only CMB data and a prior on the Hubble constant and do not include any constraints from LSS. For the primary CMB data, we use the Markov Chain Monte Carlo (MCMC) results of \\citet{lewis02} as an estimate of the appropriate weighting for each cosmological model. For the SZ effect, we use a model for the SZ power spectrum as a function of cosmology from \\citet{komatsu02}. In \\S\\ref{sec:cmb} we outline salient features of the MCMC results. In \\S\\ref{sec:sz} we sketch out our calculation of the SZ power spectrum, including the effects of the non-Gaussian nature of the SZ fluctuations. In \\S\\ref{sec:results} we compare MCMC realizations of the SZ power spectrum weighted by the primary CMB data with the CBI and BIMA data, and check for consistency between them. In \\S\\ref{sec:disc} we discuss implications for the reported detections of power at small angular scales. ",
        "conclusions": "\\label{sec:disc}  We have presented estimates of the SZ angular power spectrum at $l \\ga 2000$ for cosmological models that are consistent with the current CMB anisotropy data at $l \\la 2000$. The mode of distribution of the predicted SZ power at $l \\sim 3000$ is 50~$\\mu{\\rm K}^2$, while we find that the weighted mean is 97~$\\mu{\\rm K}^2$ (the median is 78.5~$\\mu{\\rm K}^2$), still significantly lower than the reported CBI detection of anisotropy on these angular scales, 508~$\\mu{\\rm K}^2$. Moving out from the mode at levels of equal likelihood, the ranges of power enclosed by 68\\% and 95\\% of the distribution are $20-200~\\mu{\\rm K}^2$ and $5-450~\\mu{\\rm K}^2$, respectively. This is effectively a calculation of the primary CMB prior for the SZ power spectrum at the CBI band. When compared to the quoted 68\\% ($359-624~\\mu{\\rm K}^2$) and 95\\% ($199-946~\\mu{\\rm K}^2$) regions of the CBI detection at face value, we find that the CBI detection is only marginally consistent with the SZ effect. The BIMA detection is fully consistent with the SZ effect.  By doing a careful statistical analysis, taking into account the non-Gaussianity of SZ fluctuations and sampling variance, we find that an SZ interpretation of the CBI data is inconsistent at more than $2-2.5\\sigma$ for those cosmological models which are most consistent with the primary CMB data. The models with the highest SZ power share several characteristics: all models with more than 400~$\\mu{\\rm K}^2$ of power but a few exceptions have $0.95<\\sigma_8$ and $0.05<\\Omega_{\\rm b}$. Furthermore, such models prefer a low Hubble constant ($h<0.7$) and a high matter density ($0.3<\\Omega_{\\rm CDM}+\\Omega_{\\rm b}<0.8$). A reason for a low $h$ is that the primary CMB data tightly constrain $\\Omega_{\\rm b}h^2$ while we need a higher $\\Omega_{\\rm b}h$ to make the SZ effect larger. We can accomplish this by increasing $\\Omega_{\\rm b}$ while reducing $h$ slightly within the HST-$h$ prior. A model with large red tilt ($n_{\\rm s}<0.95$) can not produce a large amount of SZ power.  The possible presence of radio point sources that partially ``fill in'' the SZ effect from clusters \\citep{holder02} further enhances the possible discrepancy. A reliable determination of $\\sigma_8$ significantly lower than 1 would be very difficult to reconcile with an SZ interpretation of the measured high-$l$ power.  Uncertainty in estimates for the SZ power spectrum is not yet fully understood, making a more detailed interpretation of the excess power complicated. A major issue is that different simulations have not yet converged, even for adiabatic simulations. Although our current knowledge of missing physics such as gas cooling, star formation, or energy feedback is still limited, these effects appear to not be very important \\citep{white02} except on very small angular scales ($<1'$ or $l>10^4$) where other effects like non-sphericity or merging of halos may also play a role. Nevertheless, the current differences among analytic models, simulations, and estimates of the effects of missing physics are at the level of a factor of two in $C_l$, while the discrepancy between these predictions and the CBI data is about a factor of five.  Given the strong dependence of $C_l$ on $\\sigma_8$, we argue that the discrepancy is due to the difference between a low $\\sigma_8\\la 0.8$ favored by the primary CMB and a high $\\sigma_8\\ga 1$ favored by the SZ effect \\citep{bond02,komatsu02}. \\cite{lahav02} and \\cite{melchiorri02} have found similarly low $\\sigma_8$ from the primary CMB data with the same prior on $h$. {\\em If the excess power is really due to the SZ effect}, then this discrepancy is suggesting that there are some missing components in our analysis. Multi-band SZ observations covering several frequencies will be required to verify the apparent discrepancy.  What is missing in our analysis? The chain that we have used has a strong HST prior on $h$. Lower values of $h$ reduce the discrepancy, but $h \\la 0.4$ would be required to explain the entire difference. Additional components such as massive neutrinos would make the discrepancy worse by driving $\\sigma_8$ to even lower values. Allowing tensor modes will have competing effects of reducing the overall normalization of scalar modes at large scales but also allowing a blue tilt (higher $n_{\\rm s}$), leaving the effects on cluster scales largely unchanged. The effect of allowing a general equation of state for the dark energy will be to slightly enhance the SZ fluctuations for a fixed value of $\\sigma_8$ \\citep{komatsu02}, but to significantly reduce the CMB-preferred value of $\\sigma_8$. With the SZ fluctuation power going as $\\sigma_8^7$, the latter effect will dominate and $w>-1$ will generally reduce the expected fluctuation power. The addition of isocurvature fluctuations may help to reconcile the discrepancy, while it significantly expands the allowed range of many cosmological parameters \\citep{trotta01}.  A very early reionization of the universe ($z_{\\rm re}>20$) will increase the CMB-preferred value of $\\sigma_8$, helping to reduce the discrepancy. We find a broad peak in the distribution of models in the chain around $\\sigma_8e^{-\\tau}\\sim 0.8$, where $\\tau$ is the Thomson-scattering optical depth of the universe; $\\tau\\ga 0.22$ or $z_{\\rm re}\\ga 20$ would comfortably allow $\\sigma_8\\ga 1.0$. The discrepancy thus disappears if the universe was reionized early. The current CMB data cannot break the degeneracy between $\\sigma_8$ and $\\tau$, and allow this area of parameter space, although $\\tau \\ga 0.3$ ($z_{\\rm re} \\ga 30$) appears to be ruled out. CMB-polarization experiments on large angular scales (e.g., MAP or Planck) should be able to break this degeneracy, and detect the signature of reionization \\citep{zaldarriaga97,eisenstein99,kaplinghat02}.  We have presented an example of the ease and power of MCMC methods in applying CMB constraints to calculations that include non-trivial dependence on cosmological parameters. It would be worthwhile to investigate the effects of LSS priors, which should make the discrepancy worse by making $\\sigma_8$ smaller \\citep{lewis02,bond02}. With the strong preference of the high-$l$ measurements for high values of $\\sigma_8$ and $\\Omega_{\\rm b} h$, it would require running a new chain that includes the CBI and BIMA data points, as importance sampling is unreliable in the tails of the current distribution.  Measurements of CMB anisotropy will continue to improve. MCMC methods provide a natural way to incorporate strong constraints on cosmological parameters from CMB experiments into other cosmological studies. Using CMB information to understand the effects of cosmology will allow better understandings of systematic errors in the measurements and insight into important astrophysical processes."
    },
    "0207/astro-ph0207515_arXiv.txt": {
        "abstract": "We present high-resolution ($R \\simeq 18,000$), high signal-to-noise 2 $\\mu$m spectra of two luminous, X-ray flaring Class I protostars in the $\\rho$ Ophiuchi cloud acquired with the NIRSPEC spectrograph of the Keck II telescope. We present the first spectrum of a highly veiled, strongly accreting protostar which shows photospheric absorption features and demonstrates the stellar nature of its central core. We find the spectrum of the luminous ($L_{bol}$ = 10 $L_\\odot$) protostellar source, YLW 15, to be stellar-like with numerous atomic and molecular absorption features, indicative of a K5 IV/V spectral type and a continuum veiling $r_k = 3.0$. Its derived stellar luminosity (3~$L_\\odot$) and stellar radius (3.1~$R_\\odot$) are consistent with those of a 0.5 $M_\\odot$ pre-main-sequence star. However, 70\\% of its bolometric luminosity is due to mass accretion, whose rate we estimate to be $1.6 \\times 10^{-6} M_\\odot$ yr$^{-1}$ onto the protostellar core. We determine that excess infrared emission produced by the circumstellar accretion disk, the inner infalling envelope, and accretion shocks at the surface of the stellar core of YLW 15 all contribute significantly to its near-IR continuum veiling. Its projected rotation velocity $v$ sin $i$ = 50 km s$^{-1}$ is comparable to those of flat-spectrum protostars but considerably higher than those of classical T Tauri stars in the $\\rho$ Oph cloud. The protostar may be magnetically coupled to its circumstellar disk at a radius of 2 $R_*$. It is also plausible that this protostar can shed over half its angular momentum and evolve into a more slowly rotating classical T Tauri star by remaining coupled to its circumstellar disk (at increasing radius) as its accretion rate drops by an order of magnitude during the rapid transition between the Class I and Class II phases of evolution.  The spectrum of WL 6 does not show any photospheric absorption features, and we estimate that its continuum veiling is $r_k \\geq 4.6$. Its low bolometric luminosity (2 $L_{\\odot}$) and high veiling dictate that its central protostar is very low mass, M $\\sim$ 0.1 M$_\\odot$. ",
        "introduction": "The discovery of accreting protostars has been a triumph for understanding low-mass star formation. Infrared and sub-millimeter data have revealed that these objects are deeply embedded, cold, and luminous. The spectral energy distributions (SEDs) of such Class I protostars gradually rise from the near-IR, peak in the far-IR ($\\lambda \\sim 100 \\mu$m), and fall off in the sub-mm and mm wavelength regimes \\citep[e.g.,][]{ALS87}. Observations and models have indicated that protostars consist of three primary components: a massive infalling envelope, a circumstellar accretion disk, and an embryonic stellar core.  With a size of $\\sim$ 10$^4$ AU, the infalling envelope is the largest component of a protostellar object. It extends to within a few AU of the embryonic stellar core and absorbs and reprocesses all the visible (and UV) as well as much of the near-infrared radiation emitted by the central protostar. The infalling envelope contributes the bulk of the observed far-infrared to millimeter-wave luminosity and a significant portion of the detected near- and mid-infrared emission.  The circumstellar accretion disk can be as much as a few hundred AU in size and it contributes a significant portion of the observed mid- and near-infrared luminosity of the protostar. During the protostellar phase of evolution, the embryonic stellar core at its center grows primarily through accretion of material from the circumstellar disk, which itself acquires its mass directly from the infalling envelope.  The embryonic stellar core at the center of the protostar is thought to be a stellar-like object a few solar radii in extent \\citep[e.g.,][]{SST80}. It is the source of the bulk of the protostellar luminosity which is generated by both a strong accretion shock at the stellar surface and the radiation of internal stellar energy. This internal energy is provided by its self-gravity as well as the release of gravitational potential energy during the protostar's assembly from accreting and infalling material. The core's accretion-dominated luminosity is radiated primarily at UV and visual wavelengths, but most, if not all of this radiation is absorbed by dust in the infalling envelope which renders the central protostar invisible at these wavelengths. As a result, little is known about the nature of the embryonic stellar cores at the hearts of protostellar objects.  The observational technique of high-resolution near-IR spectroscopy has the potential to reveal the nature of these heavily embedded central objects. Indeed, we have recently been able to demonstrate the stellar nature of the central objects in flat-spectrum protostars and thus discern their spectral types, effective temperatures, rotation velocities, luminosity classes, and masses using this technique \\citep[hereafter Paper I] {PaperI}.  These objects are in a later evolutionary state than Class I protostars; they are still surrounded by substantial circumstellar disks and envelopes but their accretion luminosities (and rates) have decreased below Class I levels.  However, our previous high-resolution IR spectroscopic study of Class I young stellar objects (YSOs)  did not reveal any spectral absorption features in these protostars \\citep[hereafter Paper II]{PaperII}.  \\cite{LR99} did detect absorption features in the moderate-resolution near-IR spectrum of at least one Class I YSO (YLW 16A / IRS 44), but the moderate continuum veiling of that object ($r_k = 1$) suggests that it is intrinsically similar to a flat-spectrum YSO. Also, \\citet{KBTB98} detected TiO absorptions in moderate-resolution optical spectra of three Class I YSOs in the Tau-Aur dark clouds.  However, these are all low luminosity objects, $L_*$ $\\simeq$ $L_{bol} \\lesssim 0.5$ $L_\\odot$, indicating relatively low levels of accretion.  Thus there is very little or no existing information concerning the nature and basic physical characteristics of the central stellar cores of protostars undergoing significant mass accretion.  Observational determination of the effective temperatures and radii (or surface gravities) of these deeply embedded central objects would provide critical constraints for theoretical protostellar models which predict the values of these parameters as a protostar evolves \\citep{SST80}. Empirical knowledge concerning the initial angular momentum of these objects as well as the evolution of angular momentum in the earliest phases of young stellar evolution is also highly desirable to test protostellar theory. Many classical T Tauri stars (CTTSs) embedded in the $\\rho$ Oph cloud and elsewhere are known to rotate slowly, $v$~sin $i$~$<$~20~km~s$^{-1}$. However, we have found that flat-spectrum protostars in this cloud have similar radii but rotate much more quickly, $v$~sin~$i$~$\\simeq$~40 km s$^{-1}$ (Paper I). Do heavily-accreting Class I protostars in Ophiuchus rotate as fast or faster? If so, how is stellar angular momentum reduced so quickly by the classical T Tauri pre-main-sequence (PMS) evolutionary phase?  We have undertaken a new, very sensitive near-IR spectroscopic study of Class I protostars with the 10 m Keck II telescope in order to address these issues. The Keck observations are considerably more sensitive than our previous deep infrared spectral survey made with the 3-m NASA IRTF. In this paper we report our findings for two of these objects in the $\\rho$ Oph dark cloud, YLW 15 (IRS 43) and WL 6. These are 2 of the 3 $\\rho$ Oph Class I protostars detected in hard X-rays by the ASCA satellite, and both showed strong, variable, and flaring emission \\citep{KKTY97}. We describe our data acquisition and reduction in \\S 2, analyze the spectra in \\S 3, discuss the results in \\S 4, and summarize our conclusions in \\S 5. ",
        "conclusions": "  1. We have detected numerous late-type absorption features in the photosphere of YLW 15. These features include Si (2.0923 $\\mu$m and 2.2069 $\\mu$m), Mg (2.1066 $\\mu$m), Al (2.1099 $\\mu$m), Na (2.2062, 2.2090 $\\mu$m), Sc (2.2058, 2.2071 $\\mu$m), Ti (primarily 2.2217 and 2.2240 $\\mu$m), and Mg (2.2814 $\\mu$m) lines as well as the CO $v = 0 - 2$ band. All of these features are consistent with a single K5 spectral type ($T_{eff} = 4300$ K) with dwarf-like surface gravity (luminosity class IV/V).  2.  YLW 15 has a low-mass central protostar ($M \\simeq 0.5 M_\\odot$) with effective temperature, surface gravity, and {\\it stellar} luminosity similar to that of a T Tauri star despite its high bolometric luminosity ($L_{bol} \\simeq 10$ $L_\\odot$).  We find that the stellar luminosity of YLW 15 is 3 $L_\\odot$ and its accretion luminosity is 7 $L_\\odot$. Assuming the magnetic accretion model of \\citet{SNOWRL94}, we derive the disk mass accretion rate of YLW 15 to be $\\dot{M_D}= 2.3 \\times 10^{-6} M_\\odot$ yr$^{-1}$ and its stellar mass accretion rate to be $\\dot{M_*}= 1.6 \\times 10^{-6} M_\\odot$ yr$^{-1}$. These rates are consistent with the mass accretion rate of the ambient cloud gas and the stellar mass of YLW 15 given the lifetimes of Class I protostars.  3. YLW 15 is highly veiled in the near-IR, $r_k \\geq 3$.  We explore the origin of this veiling and evaluate the contributions of the stellar surface accretion, circumstellar disk, and circumstellar envelope components. We estimate that the circumstellar disk and envelope contribute equally to the veiling each providing $\\sim$ 45\\% of the needed excess emission. We find a significant amount ($r_k \\sim 0.33$) of veiling is provided by the accretion shocks at the stellar surface. Although this contribution is relatively small ($\\sim$ 11\\%) compared to the total, it is significantly greater than that produced by similar processes in typical CTTSs. We determine that emission from the circumstellar envelope is inadequate to veil absorption features formed in the inner dust-free circumstellar disk. We conclude that there are no strong absorption lines formed in the warm inner disk of YLW 15.  4. The projected rotation velocity of YLW 15 is $v$ sin $i$ = 50 km s$^{-1}$, much faster than that of a typical T Tauri star in the $\\rho$ Oph cloud. However, this rotation rate does not appear to be correlated with the X-ray flarings and variability observed in multi-epoch X-ray observations of this source.  We do find that it is plausible that YLW 15 is magnetically coupled to its accretion disk at a radius of 2 $R_*$ in the magnetic accretion model of \\citet{SNOWRL94}. This suggests -- and we verify the plausibility -- that stellar angular momentum decreases by over a factor of 2 in only $\\sim 10^5$ yr between the protostellar and T Tauri evolutionary phases.  5. The near-IR spectrum of WL 6 does not show any obvious absorption features, and we estimate that the veiling of this YSO is $r_k > 4.6$. This high veiling and its low luminosity ($L_{bol} \\simeq 2 L_\\odot$) dictate that its central protostar is likely to have low stellar luminosity ($L_* \\simeq 0.6 L_\\odot$) and to be low mass ($M \\sim 0.1 M_\\odot$)."
    },
    "0207/astro-ph0207500.txt": {
        "abstract": "The methods of abundance determinations in \\hii\\ regions and planetary nebulae are described,  with emphasis on the underlying assumptions and inherent problems. Recent results on abundances in Galactic \\hii\\ regions and in Galactic and extragalactic Planetary Nebulae are reviewed. ",
        "introduction": "\\subsection{Empirical methods}  These are methods in which no check is made for the consistency of the derived abundances with the observed properties of the nebulae.  They can be schematically subdivided into direct methods and statistical methods.  \\subsubsection{Direct methods}  The abundance ratio of two ions is obtained from the observed intensity ratio of lines emitted by these ions.  For example, \\Opp/\\Hp\\ can be derived from \\begin{equation} {\\rm O ^{++}/ H^{+}} = \\frac{{\\rm[ O \\,{\\sc{iii}}]}\\,\\lambda 5007/{\\rm H}\\beta} {j_{{\\rm [ O \\,{\\sc iii}]}(T_{e},n)}/j_{{\\rm H}\\beta (T_{e})}}, \\end{equation} where $j_{\\rm{[O}\\,{\\sc iii}]}(T_{e},n)$ is the emission coefficient of the \\Oiii\\ line, which is dependent on \\Te\\ and $n$ (assumed uniform in the nebula).  \\Te\\ can be derived using the ratio of the two lines \\Oiiit\\ and \\Oiii, which have very different excitation potentials. Other line ratios can also be used as temperature indicators in nebulae, such as \\rNii\\ and \\rSiii. The Balmer and Paschen jumps, the radio continuum and radio recombination lines also allow to estimate the electron temperature, but the measurements are more difficult.   The density is usually derived from intensity ratios of two lines of the same ion which have the same excitation energy but different collisional deexcitation rates.  The most common such ratio is \\rSii. Far infrared lines can also be used to determine densities. Each line pair is sensitive in a given density range (about 2 to 3 decades), which can be ranked as follows (Rubin \\etal\\ 1994): [N\\,{\\sc ii}] $\\lambda$122$\\mu$/205$\\mu$, [O\\,{\\sc iii}] $\\lambda$52$\\mu$/88$\\mu$, [S\\,{\\sc ii}] $\\lambda$6731/6717, [O\\,{\\sc ii}] $\\lambda$3726/3729, [S\\,{\\sc iii}] $\\lambda$18.7$\\mu$/33.6$\\mu$, [A\\,{\\sc iv}] $\\lambda$4740/4711, [Ne\\,{\\sc iii}] $\\lambda$15.5$\\mu$/36.0$\\mu$, [A\\,{\\sc iii}] $\\lambda$8.99$\\mu$/21.8$\\mu$, C\\,{\\sc iii}] $\\lambda$1909/1907.  The electron density can also be measured by the ratio of high order hydrogen recombination lines.   Plasma diagnostic diagrams combining all the information from temperature- and density-sensitive line ratios can also be constructed for a given nebula (e.g.  Aller \\& Czyzak 1983), plotting for each pair of diagnostic lines the curve in the (\\Te,$n$) plane that corresponds to the observed value.  The curves usually do not intersect in one point, due to measurement errors and to the fact that the nebula is not homogeneous (and also to possible uncertainties in the atomic data) and provide a visual estimate of the uncertainty in the adopted values of \\Te\\ and $n$.   The total abundance of a given element relative to hydrogen is given by the sum of abundances of all its ions.  In practise, not all the ions present in a nebula are generally observed.  The only favourable case is that of oxygen which in \\hii\\ regions is readily determined from: \\begin{equation} \\rm{O}/\\rm{H} = \\rm{O}^{+}/\\rm{H}^{+} + \\rm{O}^{++}/\\rm{H}^{+}. \\end{equation} Note that even if \\Oi\\ is observed, it should not be included in the determination of the oxygen abundance, since the reference hydrogen line is emitted by H$^{+}$, while O$^{0}$ is tied to H$^{0}$.  In almost all other cases (except in some cases when multiwavelength data are available), one must correct for unseen ions using ionization correction factors.  A common way to do this in the 70' and 80' and even later was to rely on ionization potential considerations, which led to such simple expressions as: \\begin{equation} \\rm{N}/\\rm{O} = \\rm{N}^{+}/\\rm{O}^{+}, \\end{equation} \\begin{equation} \\rm{Ne}/\\rm{O} = \\rm{Ne}^{++}/\\rm{O}^{++}, \\end{equation} \\begin{equation} \\rm{C}/\\rm{O} = \\rm{C}^{++}/\\rm{O}^{++}. \\end{equation}   In high excitation PNe where \\heii\\ lines are seen, oxygen can be present in ionization stages higher than O$^{++}$.  A popular ionization correction scheme for oxygen (e.g.  Torres-Peimbert \\& Peimbert 1977) was: \\begin{equation} \\frac{\\rm{O}}{\\rm{H}} = \\frac {(\\rm{He}^{+}+\\rm{He}^{++})}{\\rm{He}^{+}} \\frac{(\\rm{O}^{+}+\\rm{O}^{++})}{\\rm{H}^{+}}. \\end{equation} Expressions (2.29 -- 2.31) are based on the similarity the ionization potentials of \\Cp, \\Np, \\Op, \\Nep. Expression (2.32) is based on the fact that the ionization potentials of \\Hep\\ and \\Opp\\ are identical.  However, photoionization models show that such simple relations do not necessarily hold. For example, the charge transfer reaction \\Opp\\ + \\Ho\\ $\\rightarrow$ \\Op\\ + \\Hp\\ being much more efficient than the \\Nepp\\ + \\Ho\\ $\\rightarrow$ \\Nep\\ + \\Hp\\ one, \\Nepp\\ is more recombined than \\Opp\\ in the outer parts of nebulae and in zones of low ionization parameter.  Also, while it is true that no \\Oppp\\ ions can be found outside the \\Hepp\\ Str\\\"{o}mgren sphere, since the photons able to ionize \\Opp\\ are absorbed by \\Hep, \\Opp\\ ions can well be present inside the \\Hepp\\ zone.  Ionization correction factors based on grids of photoionization models of nebulae are therefore more reliable. Complete sets of ionization correction factors have been published by Mathis \\& Rosa (1991) for \\hii\\ regions and Kingsburgh \\& Barlow (1994) for planetary nebulae, or can be computed from grid of photoionization models such as those of Stasi\\'{n}ska (1990), Gruenwald \\& Viegas (1992) for single star \\hii\\ regions, Stasi\\'{n}ska et al.  (2001) for giant  \\hii\\ regions, Stasi\\'{n}ska et al.  (1998) for PNe.  However, it must be kept in mind that ionization correction factors from model grids may be risky too, both because the atomic physics is not well known yet (see Sect. 3.1) and because the density structure of real nebulae is more complicated than that of idealized models. The most robust relation seems to be $\\rm{N}/\\rm{O} = \\rm{N}^{+}/\\rm{O}^{+}$ (but see Stasi\\'{n}ska \\& Schaerer 1997). Such a circumstance is fortunate, given the importance of the N/O ratio both in \\hii\\ regions (as a constraint for chemical evolution studies) and in PNe (as a clue on PNe progenitors).  In spite of uncertainties, ionization correction factors often provide more accurate abundances than summing up ionic abundances obtained combining different techniques in the optical, ultraviolet and infrared domains.  Note that there is no robust empirical way to correct for neutral helium to derive the total helium abundance. The reason is that the relative populations of helium and hydrogen ions mostly depend on the energy distribution of the ionizing radiation field, while those of ions from heavy elements are also a function of the gas density distribution.   In summary, direct methods for abundance determinations are simple, powerful, and provide reasonable results (provided one keeps in mind the uncertainties involved, which will be developed in the next sections). Until recently, abundances were mostly derived from collisionally excited optical lines. This is still the case, but the importance of infrared data is growing, especially since the ISO mission. IR line intensities have the advantage of being almost independent of temperature. They arise from a larger variety of ions than optical lines. They allow to probe regions highly obscured by dust. However, they suffer from beamsize and calibration problems which are far more difficult to overcome than in the case of optical spectra. Abundance determinations using recombination lines of heavy elements have regained interest these last years. They require high signal-to-noise spectroscopy since the strengths of recombination lines from heavy elements are typically 0.1\\% of those of hydrogen Balmer lines. They will be discussed more thoroughly in the next sections, since they unexpectedly pose one of the major problems in nebular astrophysics.   \\subsubsection{Strong line or statistical methods}  When the electron temperature cannot be determined, for example because the observations do not cover the appropriate spectral range or because temperature sensitive lines such as \\Oiiit\\ cannot be observed, one has to go for statistical methods or ``strong line methods''.  These methods have first been introduced by Pagel et al. (1979) to derive metallicities in giant extragalactic \\hii\\ regions. They have since then being reconsidered and recalibrated by many authors, among which Skillman (1989), McGaugh (1991, 1994), Pilyugin (2000, 2001).  Pagel et al. (1979) proposed to use the 4 strongest lines of O and H : \\Ha, \\Hb, \\Oii\\ and \\Oiii. From Sect. 1, the main parameters governing the relative intensities of the emission lines in a nebula are : $<$\\Tstar $>$, the mean effective temperature of the ionization source, the gas density distribution (parametrized by $U$ in the case of homogeneous spheres), and the metallicity, represented by O/H. Luckily oxygen is at the same time the main coolant in nebulae, and the element whose abundance is most straightforwardly related to the chemical evolution of galaxies. The spectra must be corrected for reddening, which is done by comparing the observed \\Ha/\\Hb\\ ratio with the case B recombination value at a typical $T_{e}$\u00a0 and assuming a reddening law (see Sect. 3.3). Therefore two independent line ratios, \\Oii/\\Hb\\ and \\Oiii/\\Hb, remain to determine three quantities. Statistical methods rely on the assumption that $<$\\Tstar $>$ (and possibly $U$) are closely linked to the metallicity, and that it is the metallicity which drives the observed line ratios.  Basing on available photoionization model grids, Pagel et al. showed that (\\Oii\\ + \\Oiii)/\\Hb, later called O$_{23}$,  could be used as an indicator of  O/H at metallicities above half-solar. Skillman (1989) later argued that this ratio could also be used in the low metallicity regime, in cases when the observations did not have sufficient signal-to-noise to measure the \\Oiiit\\ line intensity. McGaugh (1994) improved the method and proposed to use both \\Oiii/\\Oii\\ and O$_{23}$ to determine simultaneously O/H and $U$ (his method should perhaps be called the O$_{23+}$ method).  For the reasons explained above, the same value of (\\Oii\\ + \\Oiii)/\\Hb\\  can correspond to either a high or a low value of the metallicity. A useful discriminator is \\Nii/\\Oii, since it is an empirical fact that  \\Nii/\\Oii\\ increases with O/H (McGaugh 1994).  The expected accuracy of statistical methods is typically 0.2 -- 0.3 dex, the method being particularly insensitive in the turnover region at O/H around  $3~10^{-4}$.  On the low metallicity side, the method can easily be calibrated with data on metal-poor extragalactic \\hii\\ regions where the \\Oiiit\\ line can be measured.  Recently, Pilyugin (2000) has done this using the large set of excellent quality observations of blue compact galaxies by Izotov and coworkers (actually, the strong line method proposed by Pilyugin differs somewhat from the O$_{23}$ method, but it relies on similar principles). He showed that the method works extremely well at low metallicities (with an accuracy of about 0.04~dex). This is a priori surprising, since giant \\hii\\ regions are powered by clusters of stars that were formed almost coevally.  The most massive stars die gradually, inducing a softening of the ionizing radiation field on timescales of several Myr, which should affect the O$_{23}$ ratio, as shown by McGaugh (1991) or Stasi\\'{n}ska (1998).  As discussed by Stasi\\'{n}ska et al. (2001), data on \\hii\\ regions in blue compact dwarf galaxies are probably biased towards the most recent starbursts, and the dispersion in $<$\\Tstar$>$  is not as large as could be expected a priori. Another possibility, advocated by Bresolin et al. (1999) in their study of giant \\hii\\ regions in spiral galaxies is that some mechanism must disrupt the \\hii\\ regions after a few Myr. Of course, the  O$_{23}$ method is expected of much lower accuracy when applied to \\hii\\ regions ionized by only a few stars, since in that case the ionizing radiation field varies strongly from object to object.  On the high metallicity side (O/H larger than about $5~10^{-4}$), the situation is much more complex.  In this regime, there is so far no direct determination of O/H to allow a calibration of the O$_{23}$ method since the \\Oiiit\\ line is too weak to be measured (at least with 4~m class telescopes).  The calibrations rely purely on models but it is not known how well these models represent real \\hii\\ regions.  Besides, at these abundances, the  \\Oii\\ and \\Oiii\\ line intensities are extremely sensitive to any change in the nebular properties (Oey \\& Kennicutt 1993, Henry 1993, Shields \\& Kennicutt 1995). Note that the calibration proposed by Pilyugin (2001) of his related X$_{23}$  method in the high metallicity regime actually refers to O/H ratios that are lower than  5 $10^{-4}$.   Other methods have been proposed as substitutes to the O$_{23}$ method. The S$_{23}$ method, proposed by V\\'{\\i}lchez \\& Esteban (1996) and D\\'{\\i}az \\& P\\'{e}rez-Montero (2000) relies on the same principles as the O$_{23}$ method, but uses  ([S~{\\sc ii}] $\\lambda$6716, $\\lambda$6731 + [S~{\\sc iii}] $\\lambda$9069, $\\lambda$9532)/\\Hb\\ (S$_{23}$) instead of  (\\Oii\\ + \\Oiii)/\\Hb. One advantage over the O$_{23}$ method is that the relevant line ratios are less affected by reddening. Besides, the excitation levels of the \\Sii\\ and \\Siii\\ lines are lower than those of the \\Oii\\ \u00a0and \\Oiii\\ lines, so that S$_{23}$ increases with metallicity in a wider range of metallicities than O$_{23}$ (the turnover region for  S$_{23}$ is expected at O/H around $10^{-3}$). Unfortunately, \\Siii\\ is more difficult to observe than \\Oiii. Oey \\& Shields (2000) argue that the  S$_{23}$ method is more sensitive to $U$ than claimed by Diaz \\& Perez-Montero (2000). This would require futher checks, but in any case, the S$_{23}$ method could be refined into an S$_{23+}$ method in the same way as the O$_{23}$ was refined into the O$_{23+}$ method.  Stevenson et al. (1993) proposed to use [Ar~{\\sc iii}] $\\lambda$7136] / [S~{\\sc iii}] $\\lambda$9532 as an indicator of the electron temperature in metal-rich \\hii\\ regions, and therefore of their metallicity. This method relies on the idea that the Ar/S ratio is not expected to vary significantly from object to object, and that the \\Arpp\\ and \\Spp\\ zones should be coextensive. However, photoionization models show that, because of the strong temperature gradients expected at high metallicity, this method could lead to spurious results.  Alloin et al. (1979)  proposed to use \\Oiii/\\Nii\\ \u00a0as a statistical metallicity indicator. While this line ratio depends on an additional parameter, namely N/O, the accuracy of this method turns out to be similar to that of statistical methods mentioned above. More recently, Storchi-Bergman et al. (1994), van Zee et al. (1998) and Denicolo et al. (2001) advocated for the use of the \\Nii/\\Hb\\ ratio (N$_{2}$) as metallicity indicator. Similarly to \\Nii/\\Oiii, this ratio  shows to be correlated with O/H over the entire range of observed metallicities in giant \\hii\\ regions. The reason why, contrary to the  O$_{23}$ ratio, it increases with O/H even at high metallicity is due to a conjunction of \\Nii/\\Hb\\ being less dependent on \\Te\\ than O$_{23}$, N/O being observed to increase with O/H in giant \\hii\\ regions (at high metallicity at least) and $U$ tending to decrease with metallicity. The advantage of this ratio is that it is independent of reddening and of flux calibration, and is only weakly affected by underlying stellar absorption in the case of observations encompassing old stellar populations. This makes it extremely valuable for ranking metallicities of galaxies up to redshifts about 2.5.  As mentioned above, statistical methods for abundance determinations assume that the nebulae under study form a one parameter family. This is why they work reasonably well in giant \\hii\\ regions. They are not expected to work in planetary nebulae, where the effective temperatures range between 20\\,000~K and 200\\,000~K. Still, it has been shown empirically that there is an upper envelope in the \\Oiii/\\Hb\\ vs. O/H relation (Richer 1993), probably corresponding to PNe with the hottest central stars. The existence of such an envelope can be used to obtain lower limits of O/H in PNe located in distant galaxies.  \\subsection{Model fitting}  \\subsubsection{Philosophy of model fitting}   A widely spread opinion is that photoionization model fitting provides the most accurate abundances.  This would be true if the constraints were sufficiently numerous (not only on emission line ratios, but also on the stellar content and on the nebular gas distribution) and if the model fit were perfect (with a photoionization code treating correctly all the relevant physical processes and using accurate atomic data). These conditions are  never met in practise, and it is therefore worth thinking, before embarking on a detailed photoionization modelling, what is the aim one is pursueing.  Two opposite situations may arise when trying to fit  observations with a model.  The first one occurs  when the number of strong constraints is not sufficient, especially when no direct \\Te\\ indicator is available. Then various models may be equally well compatible with the observations. For example, from a photoionization model analysis Ratag et al. (1997) derive an O/H ratio of 2.2~$10^{-4}$ for the PN M 2-5. However, if one explores the range of acceptable photoionization models one finds two families of solutions (see Stasi\\'{n}ska 2002). The first has O/H $\\simeq$ 2.4~$10^{-4}$, the second has O/H $\\simeq$ 1.2~$10^{-3}$! The reason for such a double solution is simply the behaviour of \\Oiii/\\Hb\\ or \\Oii/\\Hb\\ with metallicity, as explained in Sect. 1.3. Note that both families of models reproduce not only the observed line ratios  (including upper limits on unobserved lines) but also the nebular size and total \\Hb\\ flux.  The other situation is when, on the contrary, one cannot find any solution that reproduces at the same time the \\rOiii\\ line ratio and the constraints of the distribution of the gas and ionizing star(s) (e.g. Pe\\~{n}a et al. 1998, Luridiana et al. 1999, Stasi\\'{n}ska \\& Schaerer 1999). The model that best reproduces the strong oxygen lines has a different value of O/H than would be derived using an empirical electron-temperature based  method.   The difference between the two can amount to factors as large as 2 (Luridiana et al. 1999). It is difficult to say a priori which of the two values of O/H -- if any -- is the correct one.  The situation where the number of strong constraints is large and everything is satisfactorily fitted with a photoionization model is extremely rare. One such example is the case of the two PNe in the Sgr B2 galaxy, for which high signal-to-noise integrated spectra are available providing several electron temperature and density indicators with accuracy of a few \\%. Dudziak et al. (2000) reproduced the 33  (resp. 27) independent observables (including imagery and photometry) with two-density component models having 18 (resp. 14) free parameters for Wray 16-423 (resp. He 2-436). Still, the models are not really unique. The authors make the point that they can reproduce the present observations with a range of values for C/H and  \\Tstar. Yet, the derived abundances are not significantly different from those obtained from the same observational data by Walsh et al. (1997) using the  empirical  method. The only notable difference is for sulfur whose abundance from the models is larger by 50\\%, and for nitrogen whose abundance from the models is larger by a factor of 2.8 in the case of He 2-436. This apparent discrepancy for the nitrogen abundance actually disappears if realistic error bars are considered for the direct abundance determinations (rather than the error bars quoted in the papers). Indeed, the fact that the nebular gas is rather dense, with different density indicators pointing at densities from 3~$10^{3}$\\cmcub\\ up to over $10^{5}$\\cmcub\\ introduces important uncertainties in the temperature derived from \\rNii\\ due to collisional deexcitation. It must be noted that realistic error bars on abundances derived from model fitting are extremely difficult to obtain, since this would imply the construction of a tremendous number of models, all fitting the data within the observational errors.  To summarize, abundances are not necessarily better determined from model fitting. However, model fitting, if done with a sufficient number of constraints, provides ionization correction factors relevant for the object under study that should be more accurate than simple formulae derived from grids of photoionization models.  This could be called a ``hybrid method'' to derive abundances. Such a method was for example used by  Aller \\& Czyzak (1983) and Aller \\& Keyes (1987) to derive the abundances in a large sample of Galactic planetary nebulae, and is still being used by Aller and his coworkers. It must however be kept in mind that if photoionization models do not reproduce the temperature sensitive line ratios, this actually points to a problem that has to be solved before one can claim to have obtained reliable abundances.  Ab initio photoionization models are sometimes used to estimate uncertainties that can be expected in abundance determinations from empirical methods. For example Alexander \\& Balick (1997) and Gruenwald \\& Viegas (1998) explored the validity of traditional ionization correction factors in the case of spatially resolved observations. A complete discussion of uncertainties should also take into account uncertainties in the atomic data and the effect of a simplified representation of reality by photoionization models.   \\subsubsection{Photoionization codes}   Photoionization codes are built to take into account all the major physical processes that govern the ionization and temperature structure of nebulae.  In addition to photoionization, recombination, free-free radiation, collisional excitation they consider collisional ionization (this is important only in regions of coronal temperatures), charge exchange reactions, which are actually a non negligible cause of recombination for heavy elements, especially if the physical conditions are such that the population of residual hydrogen atoms in the ionized gas exceeds $10^{-3}$.  Some codes are designed to study nebulae that are not in equilibrium and they may include such processes as mechanical heating and expansion cooling.   Most nebular studies use static photoionization codes, which assume that the gas is in ionization and thermal equilibrium. The most popular one is CLOUDY developed by Ferland and co-workers, for which an extensive documentation is available and which is widely in use (see Ferland 1998, and http://www.pa.uky.edu/~gary/cloudy/ for the latest release). Several dozens of independent photoionization codes suited for the study of PNe and \\hii\\ regions have been constructed over the years. Some of them have been intercompared at several workshops (P\\'{e}quignot 1986, Ferland et al. 1996 and Ferland \\& Savin 2001). The codes mainly differ in the numerical treatment of the transfer of the ionizing photons produced in the nebula: on the spot reabsorption, outward-only approximation (most codes presently), full treatment (either with classical techniques as in Rubin 1968 or Harrington 1968  or with Monte-Carlo techniques as in Och et al. 1998). They also differ in their capacity of handling different geometries. Most codes are built in plane parallel or spherical approximations, but a few are built in 3D (Gruenwald et al. 1997, Och et al. 1998). While 3D codes are better suited to represent the density distribution in real nebulae, their use is hampered by the fact that the number of free parameters is extremely large. Presently, simpler codes are usually sufficient to pinpoint difficulties in fitting observed nebulae  within our present knowledge of the physical processes occuring in them and to settle error bars on abundance determinations.   When the timescale of stellar evolution becomes comparable to the timescale of recombination processes, the assumption of ionization equilibrium is no more valid.  This for example occurs in PNe with massive ( $>$0.64\\Ms) nuclei, whose temperature and luminosity drop in a few hundred years while they evolve towards the white dwarf stage.  In that case, the real ionization state of the gas is higher than would be predicted by a static photoionization model, and a recombining halo can appear.  To deal with such situations, one needs time dependent photoionization codes, such as those of Tylenda (1979), or Marten \\& Szczerba (1997).   The nebular gas is actually shaped by the dynamical effect of the stellar winds from the ionizing stars. This induces shocks that produce strong collisional heating at the ionization front or at the interface between the main nebular shell of swept-up gas and the hot stellar wind bubble. On the other hand, expansion contributes to the cooling of the nebular gas. Several codes have been designed to treat simultaneously the hydrodynamical equations and the microphysical processes either in 1D (e.g. Schmidt-Voigt \\& K\\\"{o}ppen 1987a and b , Marten \\& Sch\\\"{o}nberner 1991,  Frank \\& Mellema 1994a, Rodriguez-Gaspar \\& Tenorio-Tagle 1998) or in 2D  (Frank \\& Mellema 1994b, Mellema \\& Frank 1995, Mellema 1995). It may be that some of the problems found with static codes will find their solution with a proper dynamical description.  However, so far, for computational reasons, the microphysics and transfer of radiation is introduced in a more simplified way in these codes.  Also, it is much more difficult to investigate a given problem with such codes, since the present state of an object is the result of its entire history, which has to be modelled ab initio. ",
        "conclusions": ""
    },
    "0207/astro-ph0207179_arXiv.txt": {
        "abstract": "\\begin{itemize} \\item We have determined the elemental abundances of Fe, Si, S, Ar, Ca, Ne, Mg, and Ni in the intra-cluster medium (ICM) using all the clusters in the archives of the \\textsl{ASCA} X-ray telescope. \\item The calcium and argon abundances are very low and are not consistent with the abundances of the two well determined $\\alpha$ elements, silicon and sulfur. \\item The results do not show a clear preference for metal enrichment by solely Type Ia supernovae or Type II. \\item Trends in the abundances as a function of temperature (mass) suggest that different processes for enrichment and distribution of metals are important on different size scales. \\end{itemize} ",
        "introduction": "The intra-cluster medium of galaxy clusters is the repository of all the metals produced by the stars in member galaxies.  The determination of the elemental abundances in clusters provides an integrated measurement of metal production throughout the history of the cluster.  The measurement of elemental abundances in clusters is in many ways more straightforward than in other objects.  Clusters are optically thin to X-rays and the spectra do not suffer from the complicating effects of radiative transport and dust common in photospheric measurements and galactic H\\,II regions.  Compilations of X-ray cluster abundances exist (White 2000), but only for the iron-dominated overall metal abundance in clusters.  Our compilation is the first large catalog of intermediate element abundances in clusters able to differentiate between several of the $\\alpha$ elements (Si, S, Ar, Ca), and the iron peak elements (Fe, Ni) observable by X-ray instruments.  The results from this analysis are average elemental abundances from ensembles of clusters with similar properties.  They provide a more general view in contrast with detailed spatially resolved abundance measurements for individual clusters from \\textsl{Chandra} and \\textsl{XMM}.  \\begin{figure}[t] \\resizebox{\\textwidth}{!}{\\rotatebox{90}{\\includegraphics{Si_S_ratios.ps}}} \\caption{This plot compares our [Si/Fe] and [S/Fe] ratios to those measured in stars.  Stellar data is from Timmes, Woosley \\& Weaver (1995); the upper gray bar is stellar [S/Fe] data, and the lower gray bar is stellar [Si/Fe] data.} \\label{starscompare} \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207386_arXiv.txt": {
        "abstract": "{\\small We have accumulated multiwavelength lightcurves for eight black hole X-ray binaries which have been observed to enter a supposed ``soft X-ray transient'' outburst, but which in fact remained in the low/hard state throughout the outburst.  Comparison of the lightcurve morphologies, spectral behaviour, properties of the QPOs and the radio jet provides the first study of such objects as a subclass of X-ray transients (XRTs).  However, rather than assuming that these hard state XRTs are different from ``canonical'' soft XRTs, we prefer to consider the possibility that a new analysis of both soft and hard state XRTs in a spectral context will provide a model capable of explaining the outburst mechanisms for the majority of black hole X-ray binaries.} ",
        "introduction": "In the low/hard state (LHS) the X-ray spectrum is dominated by a power law component; this hard X-ray emission is thought to be produced in a Comptonizing corona.  Corresponding X-ray power spectra for sources in the LHS show a high level of low frequency noise, a broken power law, and at least one quasi-periodic oscillation (QPO).  The LHS is also characterized by a powerful, low intensity jet emitting synchrotron radiation in radio (and often higher) frequencies \\cite{Fen01}.  The X-ray lightcurves of the eight LHS XRTs we consider -- V404~Cyg, A~1524-62, 4U~1543-475, GRO~J0422+32, GRO~J1719-24, GRS~1737-31, GS~1354-64, and XTE~J1118+480 -- exhibit very different morphologies, and the canonical ``FRED''-type lightcurve is not predominant.  The optical lightcurves appear generally, although not consistently, correlated with the X-rays.  Where radio coverage is available it appears that both the main outburst and secondary maxima of the X-ray lightcurves tend to be associated with radio ejections \\cite{Fen01}. Despite the similarities in X-ray and broad-band spectral behaviour of these sources, the most notable feature of these lightcurves is their inconsistency. ",
        "conclusions": "For these LHS transient outbursts, we find the following characteristics: \\begin{itemize} \\item{The X-ray and multi-wavelength lightcurves have very different morphologies.  The relationships between the emission at various wavelengths likewise differs from source to source.} \\item{A low frequency QPO is observed which increases in frequency during the outburst but is not directly correlated with the X-ray luminosity.} \\item{The broad-band spectra of the LHS transients are very similar and do not vary substantially during different epochs.} \\item{In addition to the radio signature of the jet, it may be possible to fit the broad-band spectra in the LHS with a synchrotron spectrum.} \\end{itemize}  To model LHS outbursts, the Disc Instability Model (DIM) should be able to reproduce non-FRED lightcurves and also must consider the production of the power law hard X-ray emission from the corona. Models of LHS outbursts should include the jet, which is likely to be a much more significant contributor to the X-ray luminosity and broad-band emission than has been previously assumed.  Finally, in recent ``canonical'' soft XRT outbursts (\\eg XTE~J1859+226, XTE~J1550-564) the sources have been observed to pass through the LHS on the rise from quiescence to the (V)HS.  It is important that the power requirements of the initial LHS and its associated jet are incorporated into future attempts to model transient outbursts with the DIM, which has not been done to date."
    },
    "0207/astro-ph0207665_arXiv.txt": {
        "abstract": "We analyse the stability of a magnetized medium consisting of a neutral fluid and a fluid of charged particles, coupled to each other through a drag force and exposed to differential body forces (for example, as the result of radiation forces on one phase).  We consider a uniform equilibrium and simple model input physics, but do not arbitrarily restrict the relative orientations of the magnetic field, slip velocity and wave vector of the disturbance.  We find several instabilities and classify these in terms of wave resonances.  We briefly apply our results to the structure of SiO maser regions appearing in the winds from late-type stars. ",
        "introduction": "Multiphase flows are a widespread and important phenomenon in astrophysics.  The difference between heating and cooling rates in different components of astrophysical gases often leads to the formation of a multicomponent medium, in which several phases with widely separate temperatures coexist near to pressure equilibrium. Effective multiphase behaviour can also result from differential coupling of distinct particle species to local magnetic fields or radiation driving forces.  To first order, these differential forces will lead to drift velocities between the different components of the fluid, limited by the effect of frictional terms.  However, this means that there is local source of free energy in the flow.  Radiation pressure on dust, for example, plays a major role in many models of the acceleration of winds from highly evolved, low-mass stars \\cite[\\eg{}]{macgs92}.  The dust streams through the neutral gas and transmits momentum to it through collisions.  The streaming of dust through neutral gas has also received attention in many other astrophysical contexts, including the evolution of dust bounded H{\\sc\\,ii} regions \\cite{cochran} and the radiation-driven implosion of dense globules \\cite{sanford}.  Radiation pressure on dust may levitate interstellar clouds above the disc of the Milky Way \\cite{franco}: in such clouds dust particles will stream through the neutral gas.  No magnetic field was included in any of these studies. Hartquist \\& Havnes~\\shortcite{hartquist} identified conditions under which dust grains are well-coupled to the magnetic field when the grains are driven by radiation pressure.  In many circumstances the dust, gas phase ions, and electrons may be treated as a single fluid. Except for studies of the Wardle instability of shocks in dusty media \\cite{wardle90,stone97,macls97}, investigations of instabilities in weakly ionized astrophysical media on lengthscales short compared to the Jeans length driven by fluids streaming relative to one another have been more limited.  There is a considerable literature on large-scale instabilities driven by the self-gravity of multifluid media \\cite[\\eg{}]{mous76,naka76,huba90,bz95,bals96,zw98,kn00,mams01}. These papers differ in aspects such as the number of different flow components assumed, the precise nature of the inter-species coupling and the inclusion of processes such as the self-gravity of the flow and large-scale gradients in flow properties.  Many of these papers recover a large scale instability, first described by Mouschovias~\\shortcite{mous76} and Nakano~\\shortcite{naka76}, in which the diffusion of magnetic field out of a self-gravitating clump reduces magnetic support, leading eventually to collapse.  In some, rapidly growing, small lengthscale instabilities are found \\cite{huba90,kn00,mams01}, but, to date, rather restricted classes of relative orientation of magnetic field, mean flow and wave vector have been assumed.  We also note that there are many non-astronomical examples of interspersed multiphase flows, such as clouds, fluidized beds and microbial suspensions, which have been studied extensively.  For example, Childress \\& Spiegel~\\shortcite{cs75} find buoyant instabilities, similar to those we discuss here, in systems of finite extent in both astrophysical and terrestrial contexts.  The past work on the Wardle instability and molecular cloud support has treated inhomogeneous media, in which the streaming is induced, for instance, by impulsive acceleration or by large scale variations of the magnetic field.  While astrophysical flows are necessarily inhomogeneous in the large, these variations can serve to obscure the mechanisms of small-scale instability.  Given this wide variety of physical mechanism and equilibrium structure, in this paper we outline a general analysis for a simplified physical model, in which a charged magnetized fluid streams through a neutral fluid as a result of differential body forces.  This model might most directly be related to flows with differential radiative forces on the fluids, but can be applied more widely.  By assuming uniform initial conditions and treating the modes which we find as distributions, we can study the stability of short wavelength modes in general, without needing to treat the specific global features which are important for longer wavelengths.  Our analysis complements the previous work described above, by giving stability criteria for wave-vectors of arbitrary orientation and all initial angles between the body forces and magnetic fields, albeit for rather simpler input physics.  In Section~\\ref{s:basic}, we present the basic two-fluid equations and derive the dispersion relation for linear waves.  In Section~\\ref{s:numeric}, we present numerical solutions of the dispersion relation.  For small wavelengths, we find that `resonances' (where distinct modes have similar phase velocities) are important in understanding the stability properties, and discuss a graphical method of locating these resonances in general geometries.  We then, in Section~\\ref{s:stability}, analyse the stability of the solutions of the dispersion relation, proceeding from general analysis to specific analytic stability criteria for short and long wavelengths.  These criteria compare well with the numerical results in the previous section, and confirm their generality.  In Section~\\ref{s:wind} we apply the long wavelength results to the properties of SiO maser spots in late-type stars.  Finally, in Section~\\ref{s:conclusion}, we summarize our results. ",
        "conclusions": "\\label{s:conclusion} We have seen that the slip between different fluids in a multifluid medium can drive instabilities.  At short wavelengths, several modes of instability result from resonances between waves propagating in the different fluids; these instabilities remain for longer wavelengths but they become increasingly inter-coupled and less easy to characterise.  In particular, neutral shear/ionized slow-mode resonance instabilities will grow in almost all cases where there are finite drift velocities.  The maximum growth rates we find are of order the characteristic ion/neutral collision rates in gas without internal slip, which are typically \\begin{equation} \\nu|{in} = 1.3\\times10^{-10} n|n{\\rm\\,s^{-1}} \\end{equation} \\cite{oster,bz95}, rather than inversely proportional to density as is the case for instabilities derived from leakage of magnetic flux as a result of ambipolar diffusion.  This corresponds to a characteristic lengthscale of $2.5\\times10^{-4} (v/1{\\rm\\,km\\,s^{-1}})n|n^{-1}{\\rm\\,pc}$, where we note that the most unstable wavelengths may be one hundredth of this (Table~\\ref{t:max1}).  The instabilities appear to be similar to the two-stream instability of plasma flows \\cite[\\eg{}]{melrose,bingea00}.  In Section~\\ref{s:wind}, we apply our analysis in its long-wavelength limit to the overall properties of SiO maser spots in late-type stellar winds.  We find that if the SiO spots are not to be subject to violent instabilities, the slip velocity between the phases in these regions must lie outside certain limits, or the magnetic field in these spots must be very strong.  Our short wavelengths results also have important implications for many other astrophysical systems.  In molecular clouds, we expect that the ionized sound speed is small and that the magnetic field is perpendicular to the inter-component slip velocity, but that this slip velocity is not small.  This means that we always can find directions such that inequality~(\\ref{cond}) is satisfied.  It seems likely, therefore, that molecular clouds are generically unstable to the growth of slow-mode waves.  In the present analysis we can only conjecture the non-linear endpoint of these instabilities, but the two obvious possibilities -- fractionation of the phases as a `slugged' flow, with consequent rapid loss of magnetic field support, or the limiting of the wave spectrum at finite amplitude -- each have clear practical and observational consequences for the ecology of the interstellar medium.  Our present study is substantially simplified.  This has allowed us to derive some rather general results.  In future work, we will model the structure of the systems of interest more completely, including the spatial structure of the background flow, more interacting phases (e.g., treating electrons, ions, neutrals and a spectrum of sizes of dust particle as independent species), variation of the frictional constants with state, and the detailed nonlinear evolution of the instabilities."
    },
    "0207/astro-ph0207103_arXiv.txt": {
        "abstract": "We discuss the current implementation of the ALI method into our HYDrodynamical RAdiation (Hydra) code for rapidly expanding, low density envelopes commonly found in core collapse and thermonuclear supernovae, novae and WR stars. Due to the low densities, non-thermal excitation by high energy photons (e.g. from radioactive decays) and the time dependence of the problem, significant departures from LTE are common throughout the envelope even at large optical depths.  ALI is instrumental for both the coupling of the statistical equations and the hydrodynamical equations with the radiation transport (RT). We employ several concepts to improve the stability, and convergence rate/ control including the {\\sl concept of leading elements}, the use of net rates, level locking, reconstruction of global photon redistribution functions, equivalent-2-level approach, and predictive corrector methods. For appropriate conditions, the solution of the time-dependent rate equations can be reduced to the time-independent problem plus the (analytic) solution of an ODE For the 3-D problem, we solve the radiation transport via the moment equations. To construct the Eddington tensor elements, we use a Monte Carlo scheme to determine the deviation of the solution for the RT equation from the diffusion approximation (ALI of second kind).  At the example of a subluminous, thermonuclear supernova (SN99by), we show an analysis of the light curves, flux and polarization spectra and discuss the  limitations of our approach. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207429_arXiv.txt": {
        "abstract": "s{ SS\\,433 is an X-ray binary emitting persistent relativistic double sided jets that expand into the surrounding W50 radio nebula. The SS\\,433\\,/\\,W50 system is then an excellent laboratory for studying relativistic jet interaction with the surrounding interstellar medium. In this context, part of W50 nebula has been mapped with ISOCAM at 15 micron, where the large scale X-ray jets are observed. I will show the results, particularly on the W50 western lobe, on two emitting regions detected in IR with IRAS, and observed in millimeter wavelength (CO(1-0) transition). It is uncertain whether these regions are due to heated dust by either young stars or the relativistic jet, or to synchrotron emission from shock re-acceleration regions as in some extragalactic jets from AGN. In this latter case, INTEGRAL might detect soft gamma-ray emission.} ",
        "introduction": "\\indent SS\\,433 is an X-ray binary composed of probably a massive star and probably a neutron star, but the nature of none of them has been confirmed. The binary system emits compact relativistic jets observed at subarcsecond scales in radio, and which characteristic is to show a precession movement. The Doppler shifted lines observed in the optical spectrum enable to calculate the precession model parameters, resulting in relativistic ($v=0.26$\\,c) jets covering a cone with a half opening angle of $\\theta=19.8^\\circ$, which axis has an inclinaison angle of $i=78.82^\\circ$ to the line of sight, and a precession period of about 162.5 days, the binary period being close to 13.08 days (Margon\\,\\cite{Margon} 1984). The feature of these precessing jets is that during a small time interval of the precession, the eastern jet which is approaching most of the time recedes, and the western jet approaches instead of receding.\\\\  \\indent This binary is the center of the ``sea-shell'' radio nebula W50 shown at 20\\,cm (from Dubner {\\em et al.}\\cite{Dubner} 1998) in the figure~\\ref{figW50cone}. This $2^\\circ \\times 1^\\circ$ nebula ($\\sim 120$\\,pc $\\times$ 60\\,pc at a distance of 3.5\\,kpc) has a circular central shape considered as a supernova remnant, with two extensions called lobes or ``wings'' resulting in this sea-shell or goose beak form. The eastern wing exhibits a clear helical pattern which mirrors at large scales the precession of the jets from SS\\,433. The western wing, smaller and brighter, appears to interact with a denser medium. So different ambient conditions may result in different acceleration and emission mechanisms occurring inside the remnant. On figure~\\ref{figW50cone}, the projection of the precessing movement on the sky plane is superimposed on the 20\\,cm image of the nebula showing that these wings are well constrained inside the precession cone. Thus W50 structure reveals the connection between the subarcsecond relativistic jets from SS\\,433 and the extended nebula over $\\sim 5$ orders of magnitude in scale. \\begin{figure}[!ht] \\centerline{\\psfig{figure=figW50modelePA.ps,width=12cm}} \\caption{W50 at 20\\,cm with the projected precession cone on the sky plane. The cone axis has a position angle of $100^\\circ$ according to Hjellming and Johnston (1981). The two radio lobes lie within this cone.} \\label{figW50cone} \\end{figure} ",
        "conclusions": "\\indent I have shown indications of SS\\,433 jet interaction with the medium of W50 western lobe: the radio lobes lie within the relativistic jets precession cone, the observed IR hotspots and extended emission are coincident with molecular clouds and aligned with the jet axis, and there is a possible link with the X-ray emission. The mid-IR emission could be either thermal or not (see discussion in section~\\ref{sectIRmm}).  The prospects to answer this question are new observations: with XMM-Newton and Chandra to determine the X-ray emission nature (thermal or not, or both), in near-IR to find possible embedded stars, in millimeter to complete the CO cloud mapping and find traces of shocks, and finally with INTEGRAL, SS\\,433 being part of the guaranted time, a detection in soft $\\gamma$-ray would prove that knots are sites for particle acceleration. Any results about SS\\,433\\,/\\,W50 should be regarded in the context of the quasar\\,/\\,microquasar analogy."
    },
    "0207/astro-ph0207335_arXiv.txt": {
        "abstract": "{ We present and discuss polarimetric observations performed with the VLT--UT3 (Melipal) on the afterglow of GRB~011211, $\\sim35$~hours after the burst onset. The observations yielded a 3-$\\sigma$ upper limit of $P<2.7\\%$. We discuss this result in combination with the lightcurve evolution, that may show a break approximately at the time of our observation. We show that our upper limit is consistent with the currently favored beamed fireball geometry, especially if the line of sight was not too close to the edge of the cone. ",
        "introduction": "\\label{sec:int}  It is now generally believed that the afterglow ubiquitously observed in GRBs is produced by synchrotron radiation (see, e.g., Piran \\cite{Pi99}) as a beamed relativistic fireball is decelerated by the impact with the ambient medium ({\\mes} \\& Rees \\cite{MR97}). This interpretation is confirmed by the observation of power-law decaying lightcurves (Wijers et al. \\cite{WRM97}) showing a break at $t \\sim 1$-30~days (Frail et al. \\cite{Fr01}), of power-law spectral energy distributions (Wijers \\& Galama \\cite{WG99}; Panaitescu \\& Kumar \\cite{PK01}) and of linear polarization (Covino et al. \\cite{Co99}; Wijers et al. \\cite{Wi99}; Rol et al. \\cite{Rol00}).  The derivation of the fireball opening angle from the time of breaks in the afterglow lightcurves is crucial to derive the energy budget of GRBs (Frail et al. \\cite{Fr01}). It is nevertheless a matter of open debate whether the breaks are due to collimation or to different hydrodynamical transitions (Moderski et al. \\cite{Mo00}; in 't Zand et al. \\cite{iZ01}). The presence of polarization, and in particular its evolution (Ghisellini \\& Lazzati \\cite{GL99}, hereafter GL99; Sari \\cite{Sari99}) is an alternative and unbiased way to prove that the fireball is beamed and allows to constrain the orientation of the jet with respect to the line of sight to the observer (GL99; Bj\\\"ornsson \\& Lindfors \\cite{BL00}).  Before the observation presented here, 4 GRBs have been observed in polarimetric mode at various wavelengths, yielding two positive measurements and two upper limits.  The first measurement was performed on the afterglow of GRB~990123 in the $R$ band, yielding an upper limit $P < 2.3\\%$ ($95\\%$ confidence level, Hjorth et al. \\cite{Hj99}). The first detection of linear polarization was obtained by Covino et al.  (\\cite{Co99}) on GRB~990510. Observations in the $R$ band at $t\\sim18.5$~hours after the burst yielded $P = (1.7\\pm0.2)\\%$. The detection was confirmed by Wijers et al. (\\cite{Wi99}), who obtained $P = (1.6\\pm0.2)\\%$ at $t \\sim 20$~hours, a value consistent with that of Covino et al. (\\cite{Co99}). Multiple measurements of polarization at three different epochs were performed on GRB~990712 (Rol et al. \\cite{Rol00}). While the position angle did not vary significantly (but the data are also consistent with a 45$^\\circ$ variation), a marginal detection of fluctuation of the polarized fraction was obtained, the second measurement ($P = (1.2\\pm0.4)\\%$ at $t \\sim 16.7$~hours) being smaller than the other two ($P = (2.9\\pm0.4)\\%$ and $P = (2.2\\pm0.7)\\%$ at $t \\sim 10.6$~hours and $t \\sim 34.7$~hours, respectively).  Finally, an attempt to measure near infrared (NIR) polarization in the afterglow of GRB~000301C yielded only a weak $P < 30\\%$ constraint\\footnote{Several other attempts to measure linear polarization of afterglows in the NIR were performed by the same collaboration, but it turned out that for all these bursts an optical--IR afterglow was not detected.}  (Stecklum et al. \\cite{St01}).  As a general rule, some degree of asymmetry is necessary in order to observe polarization. Two general models have been proposed to explain some degree of linear polarization in the framework of synchrotron emission. Gruzinov \\& Waxman (\\cite{GW99}) discuss how ordered magnetic field domains can diffuse in the fireball, predicting $ P\\sim 10\\%$. GL99 (and, independently, Sari \\cite{Sari99}) considered a geometrical setup in which a beamed fireball observed slightly off-axis provides the necessary degree of anisotropy (see also Sect.~\\ref{sec:model}). Variable polarization up to $10\\%$ is predicted. ",
        "conclusions": "\\label{sec:con} We have observed in polarimetric mode the optical afterglow of GRB~011211; our result is a 3-$\\sigma$ upper limit of $P < 2.7\\%$. This is consistent with previous measurements performed on other GRBs. Unfortunately a clear achromatic jet break is not observed in the burst lightcurve, and this does not allow us to perform a clear comparison with the currently favored theoretical models for the production of polarization in beamed fireballs. We can nevertheless deduce that, if the ratio of the observing angle to the jet opening angle was less than 2/3, our measurement would be consistent with the models. This result holds true if a break was present at $t\\sim2$~days (Holland et al. \\cite{Ho02}) or if it was at a much later time."
    },
    "0207/astro-ph0207273_arXiv.txt": {
        "abstract": "{ Further analysis of X-ray spectroscopy results (Willingale et al. 2002) recently obtained from the MOS CCD cameras on-board XMM-Newton provides a detailed description of the hot and cool X-ray emitting plasma in Cas A. Measurement of the Doppler broadening of the X-ray emission lines is consistent with the expected ion velocities, $\\sim1500$ km\\,s$^{-1}$ along the line of sight, in the post shock plasma. Assuming a distance of 3.4 kpc, a constant total pressure throughout the remnant and combining the X-ray observations with optical measurements we estimate the total remnant mass as 10 $M_{\\sun}$ and the total thermal energy as $7\\times10^{43}$ J. We derive the differential mass distribution as a function of ionisation age for the hot and cool X-ray emitting components. This distribution is consistent with a hot component dominated by swept up mass heated by the primary shock and a cool component which are ablated clumpy ejecta material which were and are still being heated by interaction with the preheated swept up material. We calculate a balanced mass and energy budget for the supernova explosion giving a grand total of $1.0\\times10^{44}$ J in an ejected mass; approximately $\\sim0.4$ $M_{\\sun}$ of the ejecta were diffuse with an initial rms velocity $\\sim1.5{\\times}10^{4}$~km\\,s$^{-1}$ while the remaining $\\sim1.8$ $M_{\\sun}$ were clumpy with an initial rms velocity of $\\sim2400$~km\\,s$^{-1}$. Using the Doppler velocity measurements of the X-ray spectral lines we can project the mass into spherical coordinates about the remnant. This provides quantitative evidence for mass and energy beaming in the supernova explosion. The mass and energy occupy less than 4.5 sr ($<$40\\% of the available solid angle) around the remnant and 64\\% of the mass occurs in two {\\em jets} within 45 degrees of a jet axis. We calculate a swept up mass of 7.9 $M_{\\sun}$ in the emitting plasma and estimate that the total mass lost from the progenitor prior to the explosion could be as high as $\\sim20$ $M_{\\sun}$. We suggest that the progenitor was a Wolf-Rayet star that formed a dense nebular shell before the supernova explosion. This shell underwent heating by the primary shock which was energized by the fast diffuse ejecta.} ",
        "introduction": "If we can measure the total mass, the temperature and the bulk velocity of material in a young SNR we can estimate the total energy released by the SN explosion. Coupling this with Doppler measurements we can deproject the mass and energy from the plane of the sky into an angular distribution around the centre of the SN. Here we present further analysis of XMM-Newton data (Willingale et al. 2002) that provides a quantitative assessment of the mass and energy distribution around Cas A. There is a growing body of evidence that the core collapse of massive stars is an asymmetric process. Spectra of supernovae are polarized, neutron stars produced in supernovae have high velocities, mixing of high-Z radioactive material from the core with hydrogen-rich outer layer of ejecta is very rapid, high velocity bullets have been observed in the Vela SNR (Aschenbach et al. 1995) and Cas A itself (Markert et al. 1983, Willingale et al. 2002) is composed of two oppositely directed jets. The analysis presented here confirms the non-spherical nature of the Cas A SNR and also provides details about the ionization state of the X-ray emitting plasma and the total energy and mass budget of the SN explosion. ",
        "conclusions": "We have assumed pressure equilibrium between the hot and cool plasma components to give an estimate of the filling factors within the shell volume. X-ray spectroscopy at higher spectral and spatial resolution could be used to test this assumption. Well resolved emission lines from individual knots would be associated with either the hot or cool component and observation of Doppler broadening of such lines would give us a direct measurement of the ion velocities for individual ion species in the ejecta and swept up material.  Possible supernova core collapse geometries are shown in Fig. \\ref{fig7}. \\begin{figure}[!htb] \\centering \\includegraphics[width=8cm]{submitfig7.eps} \\caption{Supernova collapse geometries} \\label{fig7} \\end{figure} The distribution of X-ray emitting mass around Cas A indicates that the original explosion was not symmetric but somewhere between an axial jet and equatorial plane geometry. The confinement to within $\\pm30$\\degr of the equatorial plane as shown in Fig. \\ref{fig6} is rather striking and the other panel in Fig. \\ref{fig6} clearly demonstrates the enhancement of the emission around the poles in the axial coordinate system. It is noteworthy that spherical collapse can be modelled in one dimension, and the axial or equatorial symmetry can be modelled using just two dimensions but the combination of axial and equatorial would require a full three dimensional treatment. It may be that the processes responsible for what we observe will only be revealed by such three dimensional modelling. The apparent asymmetry of the explosion geometry introduces the possiblity of significant shear within the expanding material during or just after the explosion. This may be the root cause of the turbulence and the clumpiness of the mass distribution in the remnant rather than hydrodynamic instabilities in the dense shell formed much later after a significant mass of surrounding material has been swept up.  The total kinetic energy derived for the ejecta is consistent with the canonical value of $10^{44}$ J. However measurements suggest that the ejected mass was rather large and the rms ejection velocity was correspondingly modest. Cas A was most emphatically a mass dominated rather than radiation dominated supernova explosion. This is in stark contrast with, for example, the Crab Nebula in which no significant ejected mass or energy from the original explosion has been identified, see for example Hester et al. (1995). Collimated or jet-induced hypernovae have been suggested as a possible solution to the energy budget problem posed by gamma ray bursts seen from cosmological distances, Wang \\& Wheeler (1998). However in these cases we are looking for collimation in a radiation dominated explosion. There is no reason to suppose that the degree of mass collimation seen in Cas A is connected with radiation collimation inferred in gamma ray burst events."
    },
    "0207/astro-ph0207104.txt": {
        "abstract": "We review the theory of measuring spectral lines in emission/absorption observations and apply it to a new survey of the 21-cm line against 79 continuum sources.  We develop an observing technique and least-squares procedure to determine the opacity profile, the expected emission profile, and their uncertainty profiles.  We discuss the radiative transfer for the two-component interstellar HI gas and use Gaussian components, separate ones for the Warm and Cold neutral media (WNM and CNM), as a practical implementation of a simple but physically correct model that successfully treats both simple and complicated profiles.  Our Gaussians provide CNM spin temperatures, upper limits on kinetic temperatures for both CNM and WNM from the line widths, column densities, and velocities; we discuss these astrophysical aspects in Paper II. ",
        "introduction": "In February 1999 we used the Arecibo\\footnote{The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under a cooperative agreement with the National Science Foundation.} telescope to begin a series of Zeeman-splitting measurements of the 21-cm line in absorption against continuum radio sources. Zeeman-splitting measurements require high sensitivity and a by-product of this survey is a set of sensitive emission/absorption line data for 79 sources from which spin temperatures and other information can be gleaned.  In \\S \\ref{stokespractice} we discuss Arecibo's instrumental effects and introduce a least-squares technique to account for both them and for angular derivatives of the HI emission. In \\S \\ref{spintempderivation} we discuss the radiative transfer of the two-component (warm and cold) HI and define our technique of Gaussian fitting as a practical means to treat radiative transfer in a physically correct but simple model. \\S \\ref{gaussiancomps} discusses the practical implementation of the Gaussian fitting process and the associated difficulties and uncertainties. \\S \\ref{slopevsgauss} compares our method with a previous method for dealing with the radiative transfer, the ``slope method''. \\S \\ref{summary} is a brief summary of the paper. Heiles (2001a) presented a preliminary report of the astrophysical implications of our Gaussian components on the WNM and CNM; Paper II (Heiles \\& Troland 2002) presents the complete discussion. ",
        "conclusions": "\\label{summary}  This paper discusses the observation and reduction techniques of our large survey of the 21-cm line in emission and absorption. We use Gaussian components and a simple but physically correct model to treat the radiative transfer issues. The major topics are as follows. \\begin{enumerate}  \\item\t\\S \\ref{stokespractice} presents the theory of extracting the opacity and expected profiles from the on- and off-source spectra. We apply this theory to the Arecibo data, which are characterized by several effects common to most telescopes but amplified at Arecibo because of its large sidelobes. The most serious instrumental effect is the impossibility of getting a true off-source spectrum; we develop an observing and reduction technique that not only solves this problem but also provides reliable estimates of uncertainty for the derived opacity and expected profiles. Our results compare well with older data that are correct, but some older data are incorrect. Surprisingly, stray radiation has little influence on Arecibo's emission profiles (\\S \\ref{gaussianprocess}.  \\item\t\\S \\ref{spintempderivation} discusses the radiative transfer of the 21-cm line for the real case in which some of the gas is Warm Neutral Medium (WNM) and some the Cold Neutral Medium (CNM). We present a simple, physically correct model for this radiative transfer for which we decompose the observed profiles into Gaussian components; the CNM components are ordered along the line of sight so that some absorb the emission of others, and the ensemble of CNM clouds is placed an arbitrary fractional distance along the line of sight through the WNM.  Because of our inclusion of radiative transfer, we derive spin temperatures that are much lower than those from previous work. Our temperatures are comparable to those derived for the Magellanic clouds using the ``slope method''. The slope method is another simple, physically correct model for the radiative transfer and works well for simple profiles, but not multicomponent opacity profiles (\\S \\ref{slopevsgauss}).  \\item Fitting Gaussians to spectra is a subjective and nonunique process. \\S \\ref{gaussiancomps} devotes considerable discussion to our method and process, with many illustrative examples to clarify our subjective biases. For the opacity spectra, we generally fit the minimum number of Gaussians required to reproduce them to within the uncertainties, and for many sources the number of blended Gaussians is small. For the expected emission profiles we fit a ``reasonable'' number of additional WNM components (\\S \\ref{gaussianprocess}). We discuss the effect of fitting either too few or too many Gaussians to a line profile and conclude that the derived spin temperatures are not very much affected.  Some optical observers fit many Gaussians to reproduce line shapes exactly. We argue that this procedure is not necessarily correct because lines are always nonthermally broadened, in which case lines are not necessarily Gaussians.  \\item Paper II provides a detailed discussion of the WNM and CNM properties, together with other astrophysical implications.  \\end{enumerate}"
    },
    "0207/astro-ph0207267_arXiv.txt": {
        "abstract": "HST imaging of M22 has allowed, for the first time, a detailed and uniform mapping of mass segregation in a globular cluster. Luminosity and mass functions from the turnoff down to the mid to lower main sequence are presented for M22 in annular bins from the centre of the cluster out to five core radii. Within the core, a significant enhancement is seen in the proportion of 0.5-0.8 $M_{\\sun}$ stars compared with their numbers outside the core. Numerical modelling of the spatial mass spectrum of M22 shows that the observed degree of mass segregation can be accounted for by relaxation processes within the cluster. The global cluster mass function for M22 is flatter than the Salpeter IMF and cannot be represented by a single power law. ",
        "introduction": "As in many areas of astronomy, the advent of the Hubble Space Telescope has revolutionised the study of globular clusters. Primarily because of crowding, ground-based observations of the central regions of globular clusters are limited to brighter stars, at or above the main sequence turnoff. HST allows access to the study of stellar populations below the turnoff including main sequence stars and white dwarfs.  Main-sequence stars below the turnoff in globular clusters (typically $m < 0.8 M_{\\sun}$) have evolved little from their initial zero-age main-sequence (ZAMS) state. Thus, mass functions derived from globular cluster luminosity functions can be used as indicators of a stellar initial mass function (IMF). Most notably in recent years, several groups have used HST WFPC2 photometry to probe mass and luminosity functions for several globular clusters down to the hydrogen burning limit. For example, \\citet{Paresce2000} have documented the turnover in the luminosity function at $\\sim 0.3 M_{\\sun}$ for a sample of twelve Galactic globular clusters. In NGC 6397 \\citet{King1998} found that the mass function increases slowly for masses down to 0.1 $M_{\\sun}$ and then drops rapidly.  Although individual globular cluster main sequence stars are little evolved from the ZAMS, the main sequence itself has been subject to modification by cluster dynamical effects. These include not only intra-cluster effects such as relaxation due to two-body interactions but also tidal interactions between a globular cluster and its Galactic environment. Relaxation of globular clusters has been studied in detail through dynamical equilibrium models \\citep{King1966,Gunn1979} and through direct numerical n-body simulations \\citep{Aarseth1999}. A comprehensive review of globular cluster dynamics is given by \\citet{Meylan1997}. Briefly, two-body interactions tend to transfer kinetic energy outward from the core and produce mass segregation, a depletion of the relative fraction of low mass stars in the central regions relative to their proportions outside the core. Only since the mid-1990's has this effect been reliably observed in globular cluster cores, for example in 47 Tuc \\citep{Paresce1995}, NGC 6752 \\citep{Shara1995} and NGC 6397 \\citep{King1995}. (Note that the core of a globular cluster is usually parameterised by the core radius, $r_{c}$, defined by \\citet{King1962} as the scale factor in his empirical formula for the surface density profile.) The most important external dynamic effect is disk shocking, which tends to strip the lightest stars out of a globular cluster during orbital crossings of the Galactic plane. To best avoid both internal and external dynamical modifications, the stellar luminosity functions in globular clusters should be obtained at radii close to the half-light radius of the cluster \\citep{Lee1991,Paresce2000}.  A further complication in deriving a global IMF is the presence of binary main-sequence stars in a globular cluster. Near-equal-mass binary stars appear on a color-magnitude diagram in a main sequence displaced upwards by 0.75 mag \\citep{Elson1998}. In only a few cases, for example NGC 6752 \\citep{Rubenstein1997}, has the photometry been sufficiently precise to resolve this binary main sequence. Normally, the presence of binary stars will contaminate a main-sequence luminosity function, particularly in the core of a cluster where, due to mass segregation effects, the binary fraction is highest. In 47 Tuc, \\citet{Albrow2001} found the fraction of binary stars to be around 13\\% in the innermost 4 $r_{c}$, with some evidence that this fraction was highest ($\\sim 20\\%$) within 1 $r_{c}$, dropping to $\\sim 8\\%$ at 2.5 $r_{c}$. Such a dropoff was also noted by \\citet{Rubenstein1997} in NGC 6752. For globular clusters showing at least a moderate degree of central concentration, $log(r_{tidal}/r_{c}) \\gtrsim 1.5$, the half-light radius is generally at least several times $r_{c}$ so luminosity functions derived at the half-light radius should be reasonably free from binary contamination.  In this paper we derive the luminosity and mass functions for M22 (NGC 6656), a globular cluster located about one third of the way between the Sun and the Galactic bulge. Our observations (taken as part of another program) are not particularly deep but cover a large spatial area from the center out to several $r_{c}$. Our focus is thus on determining the degree of mass segregation in the middle to upper main sequence rather than on probing the lowest mass stars. From four fields that we subdivide into concentric annular radial bins, we determine how the luminosity and mass functions change with radius in this cluster. Sections 2 and 3 discuss the data and their reduction. In section 4 and 5 we consider the derivation of the luminosity and mass functions. In section 6 we compare these results with a dynamical model for the cluster. ",
        "conclusions": "Extensive HST imaging of M22 has been used to determine the luminosity function for this globular cluster at a number of different radii from the cluster center. Using the \\citet{Baraffe1997} stellar isochrones, we have transformed these luminosity functions into mass functions. The proportion of higher-mass stars was found to be significantly enhanced within one core radius of the center of the cluster compared to regions outside the core. This is the first time that such a detailed mapping of mass segregation from the mid main sequence to the turnoff has been performed for a globular cluster.  Numerical simulation of the radial mass spectrum of M22 using multi-mass King-Michie models has shown that the degree of mass segregation found is well predicted by the standard theory of cluster relaxation."
    },
    "0207/astro-ph0207051_arXiv.txt": {
        "abstract": "We use the Gurvits, Kellermann, \\& Frey compact radio source angular size versus redshift data to place constraints on cosmological model parameters in models with and without a constant or time-variable cosmological constant. The resulting constraints are consistent with but weaker than those determined using current supernova apparent magnitude versus redshift data. ",
        "introduction": "Cosmological models now under consideration have a number of adjustable parameters. A simple way to determine whether a model provides a useful approximation to reality is to use many different cosmological tests to set constraints on cosmological-model-parameter values and to check if these constraints are mutually consistent (see, e.g., Maor et al. 2002; Wasserman 2002).  During the past few years much attention has been focussed on the Type Ia supernova apparent magnitude versus redshift test (see, e.g., Riess et al. 1998; Perlmutter et al. 1999; Podariu \\& Ratra 2000; Waga \\& Frieman 2000; Gott et al. 2001; Leibundgut 2001).\\footnote{ The proposed SNAP satellite should provide much tighter constraints on cosmological parameters from this test (see, e.g., Podariu, Nugent, \\& Ratra 2001a; Weller \\& Albrecht 2002; Wang \\& Lovelace 2001; Gerke \\& Efstathiou 2002; Eriksson \\& Amanullah 2002).} This cosmological test indicates that the energy density of the current universe is dominated by a cosmological constant, $\\Lambda$, or a term in the material stress-energy tensor that only varies slowly with time and space and so behaves like $\\Lambda$.  In conjunction with dynamical estimates which indicate a low non-relativistic matter density parameter $\\Omega_0$ (see, e.g., Peebles 1993), cosmic microwave background anisotropy measurements also suggest the presence of $\\Lambda$ or a $\\Lambda$-like term (see, e.g., Podariu et al. 2001b; Wang, Tegmark, \\& Zaldarriaga 2002; Baccigalupi et al. 2002; Durrer, Novosyadlyj, \\& Apunevych 2001; Scott et al. 2002; Mason et al. 2002).  However, the observed rate of multiple images of radio sources or quasars, produced by gravitational lensing by foreground galaxies, appears to favor a smaller value for $\\Lambda$ (see, e.g., Ratra \\& Quillen 1992; Helbig et al. 1999; Waga \\& Frieman 2000; Ng \\& Wiltshire 2001) than is indicated by the observations mentioned above. It is therefore of interest to examine the entrails of other cosmological tests.  In this paper we consider the redshift-angular size test, using the Gurvits, Kellermann, \\& Frey (1999) compact radio source measurements. The redshift-angular size relation is measured, for structures a few orders of magnitude larger than those considered by Gurvits et al. (1999), by Buchalter et al. (1998) for quasars, and by Guerra, Daly, \\& Wan (2000) for radio galaxies; we do not use these data sets in our analysis here. Vishwakarma (2001), Lima \\& Alcaniz (2002), and references therein, use the Gurvits et al.(1999) data to set constraints on cosmological parameters; our results are consistent with, but as discussed next extend, their analyzes.\\footnote{ The Gurvits et al. (1999) data augments that of Kellermann (1993). Stelmach (1994), Stepanas \\& Saha (1995), Jackson \\& Dodgson (1996), and Kayser (1995) discuss the Kellermann (1993) data.}  Cosmological applications of the redshift-angular size test require knowledge of the linear size of the ``standard rod\" used. Some earlier analyzes of the Gurvits et al. (1999) data appear to assume that this linear size will be determined using additional data, and so quote limits on cosmological parameters (such as $\\Omega_0$ or the cosmological constant density parameter $\\omegal$) for a range of values of this linear size. Here we note that it is best to treat this linear size as a ``nuisance\" parameter (for the cosmologically relevant part of this test), that is also determined by the redshift-angular size data, and so marginalize over it (using a prior to incorporate other information about it, if needed).\\footnote{ The situation here is similar to that for the redshift-magnitude test (e.g., Riess et al. 1998; Perlmutter et al. 1999) where one must marginalize over the magnitude of the standard candle used, treating it as a nuisance parameter. In fact, Gurvits et al. (1999) determine the linear size from the redshift-angular size data by using the model of Gurvits (1994).}  In $\\S$ 2 we summarize our computation. Results are presented and discussed in $\\S$ 3. We conclude in $\\S$ 4. ",
        "conclusions": "Constraints on cosmological model parameters derived from the redshift-angular size compact radio source data of Gurvits et al. (1999) are consistent with but less constraining than those derived from the redshift-magnitude Type Ia supernova data of Riess et al. (1998) and Perlmutter et al. (1999).  Higher quality redshift-angular size data will more significantly constrain cosmological models, and in combination with high quality redshift-magnitude data will provide a check of conventional general relativity on cosmological length scales.  \\bigskip  We are indebted to L. Gurvits for providing the binned redshift-angular size data. We acknowledge helpful discussions with J. Alcaniz, R. Daly, J. Lima, and J. Peebles, and support from NSF CAREER grant AST-9875031."
    },
    "0207/astro-ph0207098_arXiv.txt": {
        "abstract": "Much of the interstellar medium in disk galaxies is in the form of neutral atomic hydrogen, H~I.  This gas can be in thermal equilibrium at relatively low temperatures, $T\\al 300$ K (the cold neutral medium, or CNM) or at temperatures somewhat less than $10^4$ K (the warm neutral medium, or WNM). These two phases can coexist over a narrow range of pressures, $\\pmin\\leq P\\leq\\pmax$. We determine $\\pmin$ and $\\pmax$ in the plane of the Galaxy as a function of Galactocentric radius $R$ using recent determinations of the gas heating rate and the gas phase abundances of interstellar gas. We provide an analytic approximation for $P_{\\rm min}$ as a function of metallicity, far-ultraviolet radiation field, and the ionization rate of atomic hydrogen. Our analytic results show that the existence of $\\pmin$, or the possibility of a two-phase equilibrium, generally requires that ${\\rm H^+}$ exceed ${\\rm C^+}$ in abundance at $\\pmin$. The abundance of ${\\rm H^+}$ is set by EUV/soft X-ray photoionization and by recombination with negatively charged PAHs. In order to assess whether thermal or pressure equilibrium is a realistic assumption, we define a parameter $\\Upsilon\\equiv t_{\\rm cool}/t_{\\rm shk}$ where $t_{\\rm cool}$ is the gas cooling time and $t_{\\rm shk}$ is the characteristic shock time or ``time between shocks in a turbulent medium''. For $\\Upsilon < 1$ gas has time to reach thermal balance between supernovae induced shocks. We find that this condition is satisfied in the Galactic disk, and thus the two-phase description of the interstellar H~I is approximately valid even in the presence of interstellar turbulence. Observationally, the mean density $\\langle n_{\\rm H\\, I} \\rangle$ is often better determined than the local density, and we cast our results in terms of $\\langle n_{\\rm H\\, I} \\rangle$ as well. Over most of the disk of the Galaxy, the H~I must be in two phases: the weight of the H~I in the gravitational potential of the Galaxy is large enough to generate thermal pressures exceeding $\\pmin$, so that turbulent pressure fluctuations can produce cold gas that is thermally stable; and the mean density of the H~I is too low for the gas to be {\\em all} CNM. Our models predict the presence of CNM gas to $R\\simeq 16-18$ kpc, somewhat farther than previous estimates.  An estimate for the typical thermal pressure in the Galactic plane for 3~kpc$\\al R\\al 18$~kpc is $\\pth/k\\simeq 1.4\\times 10^4\\exp(-R/5.5$~kpc) K cm\\eee. At the solar circle, this gives $\\pth/k \\simeq 3000$ K cm\\eee. We show that this pressure is consistent with the ${\\rm C~I^*/C~I_{tot}}$ ratio observed by \\cite{jen01} and the CNM temperature found by \\cite{hei02}.  We also examine the potential impact of turbulent heating on our results and provide parameterized expressions for the heating rate as a function of Galactic radius. Although the uncertainties are large, our models predict that including turbulent heating does not significantly change our results and that thermal pressures remain above $P_{\\rm min}$ to $R\\simeq 18$ kpc. ",
        "introduction": "The interstellar medium (ISM) has a complex thermal and ionization structure. Much of the neutral atomic gas is observed to be either warm neutral medium (WNM) with $T\\sim 10^4$ K or cold neutral medium (CNM) with $T\\sim 100$ K \\citep{kul87,dic90}. Some of the warm gas is partially ionized, the warm ionized medium (WIM), which also has $T\\sim 10^4$ K \\citep{mck77,rey83,haf99}. A small mass fraction of the gas is in the form of hot ionized medium (HIM) with $T\\sim 10^6$ K \\citep{cox74,mck77}. Inside the solar circle, about half the interstellar gas is molecular \\citep{sco87,bro88,bron00}.  A significant simplification occurs if one focuses on the neutral atomic gas,  the CNM and WNM. Some decades ago, \\cite{fie69} demonstrated that the CNM and WNM could coexist in pressure equilibrium, so that the neutral atomic gas could be considered to be a two-phase medium.  They assumed that cosmic rays dominate the heating, but it was subsequently realized that UV starlight dominates the heating due to photoelectric emission from the dust grains in the gas \\citep{wat72}. Using the photoelectric heating rates calculated by \\cite{bak94}, Wolfire et al.\\ (1995; hereafter WHMTB) investigated the thermal balance of the WNM and CNM phases in the local ISM and showed that the two-phase model is in good agreement with a wide variety of data on the ISM in the solar vicinity.  What is the evidence for a two-phase medium elsewhere in the Galaxy? In the inner Galaxy, \\cite{gar89} found that there is H~I emission (which can originate from both CNM and WNM) at all velocities allowed by Galactic rotation.  On the other hand, they found that absorption (which originates only from CNM at their sensitivity) is somewhat less pervasive, particularly within 2 kpc of the Galactic Center. \\cite{liszt93} suggested that the H~I absorption in the inner Galaxy at $R >  2$ kpc is much higher than that reported by \\cite{gar89}. \\cite{kol01} recently repeated  the earlier H~I absorption  study and confirmed the presence of cold gas in the inner Galaxy with an absorption coefficient at $R=5$ kpc approximately 5 times higher than reported by \\cite{gar89}. In the outer Galaxy, the presence of WNM is reasonably well established \\citep{kul87}, whereas that of a widely distributed CNM is less so. Carilli, Dwarakanath, \\& Goss (1998) have measured the temperature of the WNM in absorption features seen towards Cygnus A at distances of 9 kpc and 12 kpc (and $z$ height of $\\sim 1$ kpc) using the Westerbork radio telescope. They find gas temperatures of $\\sim 6000$ K and $\\sim 4800$ K respectively, which are consistent with the low pressure and low UV field models of WHMTB for atomic gas above the plane. Several high-velocity absorption components have been observed in H~I \\citep{col88} and Na~I \\citep{semb94} that arise from CNM clouds at Galactic radii $R\\al 14$ kpc. \\cite{kol01} show H~I absorption to $R\\al 17$ kpc. Molecular clouds,  which presumably form from the CNM phase, are traced to at least $R\\sim 20$ kpc \\citep{wou89,hey01} with an extremely distant H~II region and molecular cloud complex at $R=28$ kpc \\citep{dig94}. \\cite{wou90} and \\cite{bron00} show that the molecular surface density can be fit by a radial exponential in the outer Galaxy to  $\\sim 18$ kpc \\cite[see also][]{wil97}.   In their study of the Perseus arm, \\cite{hey98}, however, find that the molecular gas disk is effectively truncated at $R\\sim 13.5$ kpc.  These two results could be consistent if molecular gas extends to greater radii in directions other than those studied by \\cite{hey98} or if isolated molecular clouds extend to distances much greater than the molecular surface density can be reliably measured from CO surveys.  Thus direct observations of cold H~I or molecular gas extend to at least $R\\sim 18$ kpc. Star formation provides an indirect test for the presence of CNM, since the gas that forms stars presumably goes through the stage of being cold H~I.  In the Galaxy, near-infrared sources, IRAS sources, and H~II regions \\citep{wou90,rud96, kor00,snell02} are seen out to $R\\sim 17-20$ kpc, suggesting that CNM extends out to at least that distance.  What can be learned from observations of H~I in other galaxies?  Two-phase atomic gas has been observed using 21 cm absorption techniques in several extragalactic systems including M31 \\citep{dic93, braun92}, M33 \\citep{dic93}, the LMC \\citep{meb97}, and the SMC \\citep{dic00}.  \\cite{braun97}, using the VLA,  examined the neutral hydrogen emission in 11 nearby spirals. By associating high brightness, narrow emission components with cold gas, he finds that the fraction of cold gas remains relatively constant until the B band surface brightness falls to $\\mu_B\\sim 25$ mag arcsec$^{-2}$, i.e., the $R_{25}$ radius.  At larger radii, the fraction drops, although in some systems more than 10\\% of the H~I is in the form of cold gas out to $(1.5-2) \\times R_{25}$. Since the  extinction-corrected radius of the Galaxy is $R_{25} = 12.25$ kpc \\citep{devau83}, cold gas in the Galaxy could extend out to $1.5-2\\times R_{25}$, or $R\\sim 18.4-24.5$ kpc. \\cite{sel99} interpret the line width of H~I in the outer parts of galaxies as being due to CNM that is being stirred by the magnetorotational instability. Evidence for recent star formation in the outer disk of M31 is presented by \\cite{cuil01}, who find a population of B stars at $1.7 \\times R_{25}$, pointing towards the presence of cold gas in the outer parts of galaxies.  The fact that the neutral atomic gas in the Galactic ISM is in  two phases is a powerful result, since two phases can coexist only over a narrow range of pressure, $\\pmax > P >\\pmin$ with $\\pmax\\al 3\\pmin$ \\citep[see \\S~\\ref{sec:Results}]{fie69}. It is thus possible to estimate the thermal pressure of the H~I with reasonable accuracy---when it is in two phases---from knowing the gas phase abundances, the dust properties, and the intensity of the radiation field. We used this property to study gas in the Galactic halo and constrain the properties of the High Velocity Clouds in \\cite{wol95b}.  \\cite{cor88} have argued that achieving the condition for a two-phase equilibrium  is a necessary step in initiating star formation in young galaxies, while \\citet{par88,par89} has suggested that two-phase equilibria play a key role in regulating the rate of star formation in disk galaxies.  The primary goal of this paper is to predict the average thermal pressure of the ISM as a function of position in the Galaxy using the two-phase criteria. To do this,  we shall extend the models of WHMTB to the inner and outer Galaxy.  In light of the observational evidence that cold gas exists in the outer Galaxy, we shall carry out our model calculations for Galactic radii between 3 kpc and 18 kpc. Knowing the thermal pressure allows one to predict the intensities of the dominant cooling lines of the gas, such as that of C~II 158 \\micron, and examine the heating and cooling routes which determine the energy budget. Locally, the thermal pressure in the ISM is measured through ultraviolet absorption line studies \\citep{jen83,jen01}.   In the near future, telescopes such as ALMA, SOFIA, and Herschel will provide additional measurements of the thermal pressure and dominant cooling lines throughout the Galaxy and in other galaxies. These will test our model for the gas thermal balance and check the importance of thermal instability.  We also can calculate whether the ISM could exist as pure WNM at various positions in the Galaxy by comparing the weight of the H~I layer with $\\pmax$. The problem of determining the phase structure of the H~I in the outer Galaxy has been considered previously by \\cite{elm94}, who find a transition to mainly WNM at $R\\ag 12-14$ kpc. Our results are compared with theirs in \\S~\\ref{sub:prev}.  Although the focus of this paper is on the determination of the thermal pressure in Galactic H~I, it is well known that the thermal pressure is only a small part of the total pressure in the gas; in particular, the turbulent pressure is considerably greater than the thermal pressure \\citep{bou90}.  In \\S~\\ref{sec:tur}, we discuss the relation between the turbulent pressure and the thermal pressure and determine the conditions under which it makes sense to consider multi-phase equilibria in a turbulent medium.  We also discuss in Appendix~\\ref{appen:turbheating} the dissipation of turbulent energy in the ISM and its potential effects on  our results. In \\S~\\ref{sec:gasdust} we discuss the distribution of gas and dust in the Galaxy between these radii, together with the abundances we have adopted.  The heating and ionization in the gas are governed by energetic photons and particles, which are discussed in \\S~\\ref{EnergeticPhotons}. The thermal and chemical processes in our model are slightly modified from those discussed by WHMTB; the differences are briefly described in \\S~\\ref{sec:thermal}. The results of our calculations are presented in \\S~\\ref{sec:Results}. We then construct a simple analytic model of a two-phase equilibrium that shows how the properties of the equilibrium scale with the input parameters (\\S~\\ref{sec:toy}). We compare our model with local and extragalactic observations in \\S~\\ref{sec:comp} and discuss our results in \\S~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Model Assumptions}  Our model for the gas heating and ionization in the Galactic disk is based on a cosmic-ray rate that is constrained by comparisons of observations to  chemical models \\cite[e.g.,][]{dish86,fed96,tak00}, direct observations of the local FUV field and soft X-ray field,  and a theoretical estimate of the EUV intensity. The soft X-ray and EUV intensity is partly derived from the calculations of Slavin et al.\\ (2000), who find that radiation from cooling supernova remnants can produce the fractional ionization seen in clouds at high latitude. The intensity used in this paper is an extension of their calculation to the Galactic plane.  In addition, a stellar EUV component from Slavin et al.\\ (1998) is added so that the total ionizing photon flux from the Galactic disk matches the recombination rate derived from H$\\alpha$ observations. An important component of the EUV and soft X-ray radiation transfer is the opacity produced by the WNM gas component. We have assumed that the opacity is provided by an ensemble of WNM clouds each of column density $N_{\\rm cl}({\\rm H~I})$.  The radiation fields at the WNM/CNM interface and cloud interior are found by passing the incident radiation through additional columns of $N_{\\rm cl}({\\rm H~I})$, and $1\\times 10^{20}$ cm$^{-2}$, respectively.  Our confidence in the adopted local parameters is strengthened by our successful modeling of the densities and temperatures in the local WNM and CNM gas. Furthermore, we obtain a good fit to the [C II] cooling rate per hydrogen as derived by UV line absorption and IR line emission studies. The good match between theory and observations further indicates that we have included all the relevant heating and cooling terms in our model, an argument that is especially strong for the CNM phase in which [C~II] dominates the cooling; any additional heating sources in our model would emerge as excess [C~II] emission. This case is not as strong for the WNM in which the [C~II] emission contributes only  $\\sim 14$\\% of the cooling and is a factor 16 weaker per hydrogen than in CNM. Additional heating terms do not strongly affect the fine-structure line emission due to the strong temperature regulation by Ly$\\alpha$ cooling.  To obtain the FUV field in other regions of the Galaxy we have relied on observations of the gas surface density, metallicity, and OB star distribution, along with a numerical integration for the mean intensity. To check our assumptions, in \\S~\\ref{sub:IRcheck}, we compared the infrared emission produced by dust grains heated by the interstellar radiation field with observations taken by the COBE satellite. The processes that determine the X-ray and cosmic ray distributions are certainly more complicated than our models allow. For example, the soft X-ray flux depends on the temperature and emission measure in the  hot ionized gas component and the optical depth to the emission regions. The temperatures and emission measures depend on the metallicity, and the optical depth depends on the structure of the interstellar medium. We note, however, that Slavin et al.\\ (2000) find that the intensity of X-rays produced by SNR emission does not depend sensitively on the ambient density of the preshock gas. We have carried through our analysis by adopting a simple approach in which we use plausible arguments for scaling the soft X-ray and cosmic ray rates to other regions of the Galaxy based on the distribution of production sources (OB stars) and destruction sinks (various gas components).  Note that we have chosen to extend the OB star distribution with a constant scale length $H_R^{\\rm OB} = 3.5$ kpc out to $R=18$ kpc. Had we adopted the OB star distribution of \\cite{bron00}, $P_{\\rm min}$ and $P_{\\rm max}$ would have been lower and it would have been easier to form CNM in the outer Galaxy. Since the actual OB star distribution in the outer Galaxy is if anything below the one we have adopted, our conclusion that CNM must exist in the outer Galaxy is strengthened.  Although the distribution of total H~I surface density is constrained by observations, the separate column densities of CNM and WNM gas are not well determined away from the solar neighborhood. In calculating the opacity for EUV and soft X-ray photons, we have assumed that the ratio of WNM to CNM surface densities and scale heights are the  solar neighborhood values \\citep{dic90} throughout the Galaxy. This in turn implies that the ratio of CNM to WNM volume filling factors in the midplane is held constant with Galactic radius. A somewhat different  prescription is given by \\cite{fer98} in which the volume fraction of WNM increases in the outer Galaxy. Note that if we allowed the WNM fraction to increase in the outer Galaxy, then the opacity to soft X-ray and EUV radiation would increase as well, thereby reducing $P_{\\rm min}$ and $P_{\\rm max}$ and making conditions less favorable for the existence of WNM. Another difference in our H~I distributions is that \\cite{fer98} used a constant surface density of $\\Sigma_{\\rm H\\,I} = 5$ $M_\\odot$ pc$^{-2}$ in the outer Galaxy to $R=20$ kpc where our distribution drops below 5 $M_\\odot$ pc$^{-2}$ beyond $R >  15$ kpc.  For constant surface density, at $R=18$ kpc the pressure $P_{\\rm WNM'}$ is a factor $\\sim 2$ higher than for our H~I distribution and would make it more likely that CNM gas can exist. In a future paper we shall compare the calculated [C II] emission with the observational data in an effort to independently derive the volume fractions of CNM and WNM gas. This will be particularly telling for the outer Galaxy, where the relative absence of photodissociation regions and H~II regions should permit a clean distinction between WNM and CNM without the confusion of predominantly molecular or ionized gas.  \\subsection{Galactic Distribution of Two-Phase ISM} \\label{sub:disc-distb}  As discussed in \\S~\\ref{sub:ISM2phase}, in Figures~\\ref{fig:PWNM}$a$ through \\ref{fig:PWNM}$c$ we plot $P_{\\rm min}$ and $P_{\\rm max}$ as a function of position in the Galactic midplane for 3 values of the WNM cloud column $N_{\\rm cl}$. We also plot the thermal pressure $P_{\\rm WNM'}$ in the Galactic midplane that would result if all of the H~I layer were WNM gas supported by thermal pressure. For the case in which there is no turbulence and the thermal pressure dominates the pressure, regions in which $P_{\\rm WNM'} > P_{\\rm max}$ {\\em must} have CNM gas. This is because only CNM gas can exist at these pressures, and mass will be converted from the WNM phase to the CNM until the pressure drops below $P_{\\rm max}$, where a two-phase medium can exist. Our figures show that this condition is satisfied over much of the Galactic disk.  $P_{\\rm WNM'}$ is calculated assuming all the diffuse gas is WNM and that thermal pressure dominates and determines the vertical scale height, which we calculate in this case locally to be $\\sim 80$ pc. However, the observed half-height of the ``WNM component'' seen by \\cite{dic90} is $\\sim 265$ pc. This result demonstrates that nonthermal pressure (due to turbulent motions, magnetic fields, and cosmic rays) dominates. We take an analogous approach to analyze the turbulent case. Assuming again that all the H~I gas is WNM, we compare the thermal pressure to $\\pmin$ and $\\pmax$. We use the {\\em observed} $\\langle n_{\\rm H\\, I} \\rangle$  (which includes the effects of turbulence and the greater scale heights which lowers $\\langle n_{\\rm H\\, I} \\rangle$) to estimate a lower limit $\\langle P_{\\rm WNM} \\rangle$ on the thermal pressure $P_{\\rm th,\\, WNM} = \\langle P_{\\rm WNM} \\rangle/f_{\\rm H\\, I}$ where $f_{\\rm H\\, I}$ is the volume filling factor of the WNM (the rest is HIM). Since $\\langle P_{\\rm WNM} \\rangle$ exceeds $\\pmax$ in the outer ($8\\,\\, {\\rm kpc} \\al R \\al 16$ kpc) Galaxy, CNM {\\em must} exist in these regions. Since $\\langle P_{\\rm WNM} \\rangle$ exceeds $\\pmin$ (and turbulence likely drives the local pressures above $\\pmax$ occasionally), and since $f_{\\rm H\\, I} < 1$, we conclude CNM very likely exists at $3\\,\\, {\\rm kpc} \\al R \\al 18$ kpc.  It is difficult, however, from our theoretical models, to rule out an interstellar medium with only HIM and CNM  (and no WNM) in which the intercloud medium is filled with HIM that maintains a pressure  $P > \\pmax$ on the CNM clouds. However, in such a scenario, the volume filling factor of the HIM must be  nearly unity. If the HIM does not fill the intercloud medium, CNM would partially convert to WNM to fill the vacuum, and the pressure in the pervasive WNM would adjust such that $\\pmax > P > \\pmin$ (see Parravano et al.\\ 2002 for further discussion). We also note that the observation of H~I 21 cm emission and absorption throughout the Galaxy strongly suggests the presence of a pervasive WNM.  Assuming that $P$ lies between $\\pmin$ and $\\pmax$, we can use our models to predict the average thermal pressure from equation (\\ref{eq:pthave}), and the results are given in Table~\\ref{tbl:physicalcond}. An approximate analytic fit to these results for our fiducial column density of $N_{\\rm cl}=10^{19}$ cm\\ee\\ is \\beq P_{\\rm th,\\, ave}/k =1.4\\times 10^4\\exp(-R_k/5.5)~~~~~{\\rm K~cm^{-3}}. \\eeq We find that at fixed Galactic radius, the pressure does not change by more that a factor $\\sim 3$ over our range of cloud columns. For our fiducial column, the pressure drops from about 8,200 K cm\\eee\\ at 3 kpc to 3100 K cm\\eee\\ at the solar circle and to 600 K cm\\eee\\ at 18 kpc.  The drop in the thermal pressure from 3 kpc to the solar circle (a factor 2.7) closely matches the drop in the magnetic pressure inferred from radio observations: Beck (2001) estimates that $B$ drops from about 10 $\\mu$G at 3 kpc to 6 $\\mu$G locally, corresponding to a pressure drop by a factor 2.8. (Note that these values for the field are larger than the rms field; as Beck points out, his values of $B$ are $\\langle B^{3.9}\\rangle^{1/3.9}$.)  We can now test the validity of our assumption that the turbulence parameter $\\Upsilon\\equiv t_{\\rm cool}/ t_{\\rm shk}\\al 1$, so that it is meaningful to discuss a two-phase medium. Our results show that locally $\\Upsilon \\approx 0.1$ for CNM and $\\Upsilon \\approx 0.3$ for WNM at a pressure of $P_{\\rm th}/k = 3000$ K cm$^{-3}$ (see eq.~[\\ref{eq:upsilon}]). As a function of Galactic radius we find that $t_{\\rm cool}\\propto T/(n\\Lambda)\\propto\\exp(R_k/2.94)$. Ignoring the weak variation of $t_{\\rm shk}$ due to the variation in $\\Sigma_{\\rm WNM}$ and in $n_{\\rm WNM}^{-0.1}$, we have $t_{\\rm shk}\\propto\\dot\\varsigma_{\\rm SN}^{-1} \\propto \\exp(R_k/3.5)$.  As a result, we find $\\Upsilon \\propto \\exp(R_k/18.4)$, and we conclude that the turbulence parameter is $\\al 1$  and weakly dependent on  radius throughout the Galactic disk.  In Appendix~\\ref{appen:turbheating} we discuss the potential role of turbulence in heating the WNM and CNM phases. Using the admittedly uncertain turbulent heating rate as a function of Galactic radius given by equation~(\\ref{eq:gturb1}), at $R=17$ kpc  we find that $P_{\\rm min}$, $P_{\\rm th,\\, ave}$, and $P_{\\rm max}$ are about a factor 2 higher than for the non-turbulent heating case. However, since $\\langle P_{\\rm WNM} \\rangle$ remains above $P_{\\rm min}$ to $R = 18$ kpc, turbulent heating does not change our conclusion that turbulent fluctuations will produce cold gas that is thermally stable in the outer Galaxy. The rate of turbulent heating does not exceed the rate of photoelectric heating out to $R\\sim 18$ kpc, and from equation (\\ref{eq:upsilonturb}) we conclude that our assumption of thermal balance remains approximately valid.  We conclude by summarizing our most important results. We have shown that both observational evidence and our theoretical models presented here indicate that the thermal pressure in the ISM of the Galaxy lies in the relatively narrow range between $\\pmin$ and $\\pmax$ for 3 kpc $< R < 18$ kpc. We have calculated $\\pmin (R)$, $\\pmax (R)$ and an estimate of the thermal pressure $P_{\\rm th,\\,ave}(R)$ in the Galaxy. We have shown that CNM gas must exist out to 18 kpc. We present phase diagrams for several galactocentric radii and for several cases of varying opacity to EUV and soft X-ray flux. Understanding the neutral phases of the ISM and their dependence on the radiation field is an important step in understanding the formation of molecular clouds and the global star formation rates in a galaxy.  Acknowledgments.  We thank L. Blitz, J. Dickey, C. Heiles, A. Lazarian, H. Liszt, E. Ostriker, and R. Snell for helpful comments, W. Dehnen for providing his code to calculate the Galactic potential, J. Slavin for providing the stellar EUV and SNR X-ray spectra, and T. Sodroski, and N. Odegard for the COBE Galactic longitude profile. We also thank the referee Don Cox for his insightful comments. MGW is supported in part by a NASA LTSA grant NAG5-9271. The research of CFM is supported in part by NSF grant AST-0098365. The research of DJH is supported by NASA RTOP 344-04-10-02, which funds the Center for Star Formation Studies, a consortium of researchers from NASA Ames, University of California at Berkeley, and University of California at Santa Cruz.  \\newpage  \\appendix"
    },
    "0207/gr-qc0207016_arXiv.txt": {
        "abstract": "Binary systems of rapidly spinning compact objects, such as black holes or neutron stars, are prime targets for gravitational wave astronomers. The dynamics of these systems can be very complicated due to spin-orbit and spin-spin couplings. Contradictory results have been presented as to the nature of the dynamics.  Here we confirm that the dynamics - as described by the second post-Newtonian approximation to general relativity - is chaotic, despite claims to the contrary. When dissipation due to higher order radiation reaction terms are included, the chaos is dampened. However, the inspiral-to-plunge transition that occurs toward the end of the orbital evolution does retain an imprint of the chaotic behaviour. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207321_arXiv.txt": {
        "abstract": "Cold and warm absorbers have beeen detected in all types of active galaxies (AGN) from low to high redshift.  This gas, located in the black hole region of AGN, is thought to play an important role in AGN unification scenarios, in explaining the X-ray background, in black hole growth and AGN evolution.  High-resolution spectroscopy with {\\sl Chandra} and {\\sl XMM-Newton} has recently revealed the signatures of warm absorbers in the form of many narrow absorption lines from highly ionized material. The richness in spectral features will provide a wealth of information on the physical processes in the central region of the few X-ray brightest, most nearby Seyfert galaxies. The long-term goal is to obtain this information for a much larger number of objects, particularly at higher redshift. This will be possible with the future X-ray observatory {\\sl XEUS}.  We provide a review of the observations of dusty and dust-free warm and cold absorbers at low and high redshift, including most recent results and exciting questions still open.  Emphasis is on the science issues that we will be able to address with {\\sl XEUS} for the first time, particularly at high redshift, including: (i) determination of metal abundances of X-ray (cold) absorbers by detection of metal absorption edges, (ii) analysis of the composition of dust mixed with cold and ionized gas (K-edges of metals in cold dust and cold gas will be resolvable from each other for the first time), (iii) measurement of the velocity field of the gas, (iv) utilization of these results to investigate the {\\em evolution} of gas and dust in AGN from high to low redshift: the evolution of abundances, dust content, ionization state, amount and velocity of gas, and its role in feeding the black hole.  We emphasize the importance of iron absorption measurements with {\\sl XEUS} at high redshift for two key issues of cosmology: the early star formation history of the universe, and the measurement of cosmological parameters. As an example, we discuss recent {\\sl XMM-Newton} observations of the high-redshift BAL quasar APM 08279+5255. ",
        "introduction": "Neutral (`cold') or ionized (`warm') gaseous material is ubiquitous in the AGN/SMBH environment, and therefore of utmost importance in understanding the AGN phenomenon, the evolution of active galaxies, their link with starburst galaxies and ULIRGs, and the X-ray background. X-ray absorption and emission features provide valuable diagnostics of the physical conditions in the X-ray gas and, in particular, allow to measure elemental abundances at high redshift, with profound consequences for our understanding of the star formation history in the early universe.  Here, we provide a short overview of previous X-ray observations of absorption in AGN, and discuss how exciting questions still open can be addressed with the {\\sl XEUS} observatory. We apologize in advance for incompleteness in citations due to space limitations. ",
        "conclusions": ""
    },
    "0207/astro-ph0207117_arXiv.txt": {
        "abstract": "{We present a data set of images of the gravitationally lensed quasar Q2237+0305, that was obtained at the Apache Point Observatory (APO) between June 1995 and January 1998. Although the images were taken under variable, often poor seeing conditions and with coarse pixel sampling, photometry is possible for the two brighter quasar images A and B with the help of exact quasar image positions from HST observations. We obtain a light curve with 73 data points for each of the images A and B. There is evidence for a long (${\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}}}$\\,100 day) brightness peak in image A in 1996 with an amplitude of about 0.4 to 0.5 mag (relative to 1995), which indicates that microlensing has been taking place in the lensing galaxy. Image B does not vary much over the course of the observation period. The long, smooth variation of the light curve is similar to the results from the OGLE monitoring of the system~\\citep{Wozniak00}. ",
        "introduction": "The quadruple quasar Q2237+0305 was discovered during the Center for Astrophysics (CfA) redshift survey \\citep{Huchra85}. In high-resolution images of the system, four quasar images at a redshift of $z = 1.695$ are seen in a cross-like geometry around the core of a barred spiral galaxy with a redshift of $z=0.0394$.  Due to its geometry this system is known as the Einstein cross \\citep{Schneider88,Yee88}.  The quasar images have a separation of about $0.9$~arcsec from the galaxy centre.  Very quickly after its discovery it was realized that this system is an ideal case for microlensing studies. The challenge in this system is to measure the brightness of the four images individually with high accuracy. In observations with seeing of larger than about an arcsecond, it is very difficult to disentangle the four quasar images, the galaxy core, and other features of the galaxy (such as the bar-like structure that is situated across the galaxy centre, see \\citealt{Yee88}, or \\citealt{Schmidt98}).  The record of semi-regular observations of this system began with the announcement of the first microlensing event by \\citet{Irwin89}. \\citet{Corrigan91} published the ``initial light curve'', that was later augmented by other individual and systematical observations \\citep{Pen93,Houde94}. In \\citet{Wambsganss92} it was emphasized that without frequent sampling, it would be difficult to extract useful information from the microlensing observations. The sample of early observations with good seeing is rather heterogeneous regarding the filters chosen, so that the interpretation is made difficult since the filter differences have to be calibrated out.  \\citet{Oestensen96} presented five years of observations of Q2237+0305 from the Nordic Optical Telescope (NOT). In all four images, microlensing variations had then been detected. A rather striking drop of about one magnitude in 1992 (with very large error bars) within $\\approx 20$ days had been found by \\citet{Pen93} on the basis of observations made at the Apache Point Observatory. In addition to this photometric evidence, \\cite{Lewis98} found spectroscopic signatures for microlensing of the broad line region of this quasar.  Recently, the OGLE team has presented a light curve \\citep{Wozniak00,Wozniak00b} covering about 1200 days between 1997 and 2001.  Updated versions of this photometric data set can be looked at at {\\tt www.astro.princeton.edu/\\linebreak[0]\\~{}ogle/\\linebreak[0]ogle2/\\linebreak[0]huchra.html}. This data set provided a major step forward, and allows qualitatively new approaches in the analysis of the light curves. The OGLE light curves are very densely sampled and show amazing brightness variations in all four quasar images with high amplitudes of more than one magnitude.  Especially image C shows a dramatic brightness peak of about 1.2 mag in 1999 that was resolved by the OGLE data in beautiful detail.  In Sect. 2 our data set is described. In Sect. 3 we explain and describe the details of the data reduction. In Sect. 4 we present our results. We conclude in Sect. 5 with a discussion. ",
        "conclusions": "\\label{discussion}  We have presented monitoring data of Q2237+0305 over a period from 1995 to early 1998. Although the error bars on individual data points are relatively large, the coverage of 73 nights clearly enables us to see some significant trends in the microlensing behaviour. In agreement with the results from the OGLE group \\citep{Wozniak00,Wozniak00b}, we find that Q2237+\\linebreak[0]0305 shows large magnitude variations of several tenths of a magnitude on timescales of less than hundred days.  It has become evident that the microlensing variations in Q2237+\\linebreak[0]0305 can happen rather smoothly over time spans of several months to years. The \\citet{Irwin89} event (also in its interpretation by \\citet{Racine92} as the first half of a double peak) seems to be an example of a short-duration microlensing variation. The \\citet{Pen93} drop is an example of a high-magnitude short-term process, but with very low signal-to-noise at that time.  All in all, it is fair to say that the observations of Q2237+0305 have started to resolve microlensing variations in great detail, and they look very much like the variations that were predicted already in the early 80s for the microlensing effect. By filling up the gaps in the light curves, long-term monitoring programs pave the way for the statistical analysis of microlensing observations using, e.g., higher-order statistics of difference light curves \\citep[e.g., ][]{Kofman97}.  The OGLE light curve started where the data discussed in this paper stopped and thus shows the continuation of our light curve. At Apache Point, and at several other observatories (Maidanak Observatory, NOT), data are still being taken. A new detector is being used at APO that makes it possible to obtain much more accurate magnitude measurements for the quasar images (see, for example, the images at {\\tt www.astro.princeton.edu/\\linebreak[0]\\~{}elt/2237.html})."
    },
    "0207/astro-ph0207671_arXiv.txt": {
        "abstract": "In the framework of the study of extragalactic radio sources, we will focus on the importance of the spatial resolution at different wavelengths, and of the combination of observations at different frequency bands. In particular, a substantial step forward in this field is now provided by the new generation X-ray telescopes which are able to image radio sources in between 0.1--10 keV with a spatial resolution comparable with that of the radio telescopes (VLA) and of the optical telescopes. After a brief description of some basic aspects of acceleration mechanisms and of the radiative processes at work in the extragalactic radio sources, we will focus on a number of recent radio, optical and X--ray observations with arcsec resolution, and discuss the deriving constraints on the physics of these sources. ",
        "introduction": "This contribution is focussed on some basic concepts regarding the study of the non--thermal emission from extragalactic radio sources based on a broad band, multifrequency approach. The principal difficulty in this study arises by the different sensitivity and spatial resolution of the instruments in different bands. Radio telescopes easily reach sub--arcsec spatial resolutions and can image very faint sources by relatively short observations. On the other hand, optical telescopes generally have only arcsec spatial resolution so that combined radio -- optical studies are limited by the lower spatial resolution of the optical telescopes. Although, the problem is alleviated by making use of optical HST observations, it still remains to a certain degree. In order to extend this approach to higher energies, X--ray telescopes with arcsec resolution (at least) and good sensitivity are required. Neverthless the poor spatial resolution of the past X--ray observatories, pioneeristic studies in this direction have been attempted in the last 10-20 years with combined radio (and optical) and {\\it Einstein} or ROSAT X--ray observations. \\begin{table}[htb] \\begin{center} \\caption{X-ray Observatories} \\begin{tabular}{lllllll} \\hline \\hline Satellite & Instrument & Flux$\\rightarrow$1 cts/s & Energy Band & Resolution \\\\ &     & (erg/s/cm$^{2}$) & (keV)  & (arcsec) \\\\ \\hline Einstein & HRI & 1.6E-10 & 0.5-4.0   &   4         \\\\ & IPC & 2.8E-11 & 0.5-4.0   &             \\\\ \\hline ROSAT    & HRI & 3.7E-11 & 0.5-2.4 &   3*         \\\\ & PSPC& 1.4E-11 & 0.5-2.4 &  20         \\\\ \\hline ASCA     & SIS & 3.3E-11 & 0.4-12  & 180         \\\\ & GIS & 3.8E-11 & 0.4-12  & 180         \\\\ \\hline BeppoSAX & MECS& 8.1E-11 & 1.3-10  &    60         \\\\ & PDS & 9.5E-11   & 13-80 &             \\\\ \\hline Chandra  & HRC-I& 2.8E-11  & 0.4-10& 0.5         \\\\ & ACIS-I& 1.1E-11 & 0.4-10& 0.5         \\\\ \\hline \\hline \\end{tabular}  {\\bf Notes}: {\\bf Column 3}, Flux$\\rightarrow$ is the flux (erg/s/cm$^2$) necessary to have 1 cts/s in the detector. {\\bf Column 5}: (*) means that the spatial resolution results affected by errors in the aspect solution associated with the wobble of the space craft. \\end{center} \\end{table} More recently, a substantial step forward has been achieved thanks to the advent of the new generation X--ray satellites: {\\it Chandra} and XMM-{\\it Newton}. In Tab.1 we report the main capabilities of a selection of past X--ray observatories compared with those of {\\it Chandra}: the abrupt increase of the spatial resolution combined with the high effective area of {\\it Chandra} represent a `new revolution' in astrophysics (data taken from Cox, 1999). For the first time it is now possible to investigate the X--ray properties of extragalactic radio sources performing spatially resolved spectroscopy over a relatively large energy band (about 0.4--10 keV). {\\it Chandra} allows us to obtain images with $\\leq$arcsec spatial resolution, comparable with the typical resolution of the VLA radio images and of the optical telescopes and thus it is a tool to perform, for the first time, multiwavelength studies from the radio band to the X--rays.  As we will show in this contribution, the high resolution broad band study of non--thermal emission from the extragalactic radio sources allows us to sample the emission due to very different portions of the spectrum of the relativistic electrons or to study the emission due to different emitting mechanisms at work in the observed regions. After the first 3 years of new {\\it Chandra} observations, it is now clear that the scenarios describing these objects should now be partially revisited.  In Sect.3 we discuss some theoretical aspects related to the relativistic plasma in extragalactic radio sources. More specifically, in Sects. 3.1--3.3 we report the basic concepts of particle acceleration and discuss the resulting spectrum of the relativistic electrons in the framework of the standard shock acceleration scenario. In Sect. 3.4 we introduce the non--thermal emitting mechanisms. In particular, we focus on the sampling of the spectrum of the emitting electrons provided by the observations at different frequencies. In Sect. 3.5 we review the basic methods to derive the energetics of the relativistic plasma in the extragalactic radio sources. Finally, in Sect.4 we concentrate on the most recent high resolution multiwavelength studies of extragalactic radio sources. Our `biased' review is mainly focussed on the new {\\it Chandra} observations of radio lobes, radio jets and radio hot spots. Combining the data with the methods showed in Sect. 3, we discuss on the new constraints on the shape of the low and high energy end of the electron spectrum, on the kinematics of radio jets, and on the energetics of extragalactic radio sources. ",
        "conclusions": "In this contribution we have tried to summarize new insights on the physics of extragalactic radio sources from recent studies based on combined observations in different bands. We concentrated ourselves on the case of the non--thermal emission produced in kpc--scale regions (radio lobes, kpc--scale jets, hot spots) and on the spectrum and energetics of the emitting relativistic electrons.  We have shown that the spectrum of the relativistic electrons in radio sources can be approximated with a power law only in a relatively narrow energy range. In general, the shape of the spectrum depends on the acceleration mechanisms active in the emitting regions and on the competition between such mechanisms and both the processes responsible for the energy losses and the spatial diffusion of the relativistic electrons. Measurements of the flux and spectrum produced by different non--thermal emitting processes in different frequency bands can allow to trace the spectrum of the emitting electrons and thus to constrain the physics of the acceleration in these remote regions. Such measurements are possible only now by combining radio, optical and X--ray observations with arcsec resolution.  Here, a `biased' summary of some of the most promising recent findings:  $\\bullet$ {\\it Low energy end of the electron spectrum}: An advance in constraining the energetics associated to the low energy electrons has coming from the new detections of extended X--ray emission from IC scattering of nuclear photons in powerful FR II radio galaxies and quasars. This effect opens a new window on the study of electrons with $\\gamma \\sim 100$ in radio lobes which are invisible in the radio band as, in general, they would emit synchrotron radiation in the 0.01 to 1 Mhz frequency range. We have stressed that, making use of the classical minimum energy formulae, these electrons are not taken into account in the calculation of the energetics of the radio lobes. On the other hand, the first X--ray detections of IC scattering of nuclear photons indicate that the bulk of the energetics of the radio lobes is contained by these low energy electrons. Additional evidences that the spectrum of the relativistic electrons extends down to low energies ($\\gamma < 500$) is provided by the possible detection of SSC optical emission from radio hot spots and by the X--ray emission from powerful radio jets due to boosted IC scattering of CMB photons.   $\\bullet$ {\\it High energy end of the electron spectrum}: The maximum energy of the electrons in extragalactic radio sources depends on the balance between acceleration and losses mechanisms. The discovery of optical synchrotron emission from hot spots in the last decades has pointed out the presence of efficient accelerators active in these regions and able to accelerate electrons up to very high energies ($\\gamma \\geq 10^5$). More recently, the discovery of synchrotron X--ray emission from an increasing number of jets of low power radio galaxies suggest the presence of even more energetic electrons ($\\gamma \\geq 10^{6-7}$) possibly accelerated in low B--field regions. These findings and the recent suggestion about synchrotron X--ray emission also from jets of powerful sources, would indicate the presence of extremely efficient accelerators in the emitting regions. The extremely short acceleration time--scales (down to $\\sim 10^2$ yrs) requested by these findings are actually starting to put interesting limits on the spatial diffusion coefficient of the electrons and thus on the microphysics prevailing in these regions. The recent additional evidence for continuous reacceleration of relativistic electrons in radio jets might finally suggest the presence of energetic turbulence and thus that part of the kinetic power of the jets is dissipated into the developing of plasma instabilities and in the re--acceleration of relativistic particles. This might match with the proposed scenario in which a slow {\\it layer}, where a fraction of the kinetic power is dissipated, surrounds a fast moving {\\it spine}.   $\\bullet$ {\\it Energetics of radio lobes and hot spots}: A crucial point in the study of the extragalactic radio sources is the calculation of the energetics associated to these objects. As we have shown, the advent of {\\it Chandra} makes it possible to extensively apply the IC method and to derive the energy density of the relativistic electrons and of the magnetic field in the emitting regions. This has been done in the case of a number of hot spots and radio lobes : the derived energetics are usually larger (1 to 30 times) than that calculated with the classical minimum energy formulae. This is a direct consequence of the presence of low energy electrons which contain most of the energetics, but it is also due to moderate departures from the minimum energy conditions found in a number of cases. Additional statistics is requested to better address this point. We further claim that deep X--ray follow up of the most unambiguous detections of diffuse IC emission from radio lobes will probably allow to derive the first maps of the magnetic field intensity in the extragalactic radio sources providing unvaluable information on the prevailing physics.   $\\bullet$ {\\it Kinematics of the radio jets}: One of the most interesting findings obtained combining the radio, optical and {\\it Chandra} X--ray data of powerful radio jets is the recent discovery of highly relativistic speeds of the jets up to several tens of kpc of distance from the nucleus. Bulk Lorentz factors $\\Gamma_{\\rm bulk} \\sim$ 3--10 up to these distances limit the radiative efficiency of the jets that should be low in order to preserve the kinetic power of the jet itself. Again, a scenario with a velocity structured jet with a fast {\\it spine} surrounded by a slow {\\it layer} might help to better understand this findings.  \\vskip 0.2cm  \\noindent {\\it Acknowledgements} I am grateful to all my collaborators, in particular to M. Bondi, A. Comastri and G. Setti for help and discussions. I am indebted to F. Mantovani who invited me to give these lectures and to M. Marcha for a careful reading of the manuscript and for useful comments."
    },
    "0207/astro-ph0207392_arXiv.txt": {
        "abstract": "{We explore the dependence of the amplitude of stellar dynamo cycle variability (as seen in the Mount Wilson Ca {\\sc II} HK timeseries data) on other stellar parameters. We find that the fractional cycle amplitude $A_{\\rm cyc}$ (i.e.\\ the ratio of the peak--to--peak variation to the average) decreases somewhat with mean activity, increases with decreasing effective temperature, but is not correlated with inverse Rossby number $Ro^{-1}$. We find that $A_{\\rm cyc}$ increases with the ratio of cycle and rotational frequencies $\\omega_{\\rm cyc}/\\Omega$ along two, nearly parallel branches. ",
        "introduction": "In a recent series of papers (Brandenburg, Saar, \\& Turpin 1998; Saar \\& Brandenburg 1999 [=SB], 2001), we have been exploring relationships between magnetic cycle periods $P_{\\rm cyc}$, rotation periods $P_{\\rm rot}$, and other stellar properties. By combining $P_{\\rm cyc}$ obtained from observations of Ca {\\sc II} emission (e.g., Baliunas et al. 1995 [=Bea95]),  photometry, and $P_{\\rm rot}$ variation (e.g., Lanza \\& Rodono 1999), we found evidence for trends between cycle ($\\omega_{\\rm cyc}$) and rotational ($\\Omega$) frequencies (Saar \\& Brandenburg 2001),  and between $\\omega_{\\rm cyc}/\\Omega$ and both $R'_{\\rm HK}$ and $Ro^{-1}$ (Brandenburg et al. 1998; SB).   Here, $Ro^{-1}$ is the inverse Rossby number and $R'_{\\rm HK}$ is the Ca {\\sc II} HK flux, corrected for photospheric contributions and normalized by the bolometric flux (see Noyes et al. 1984).  Another important, but less studied observational property of stellar dynamos is the cycle {\\it amplitudes} (i.e., the mean amplitude of the cyclic variability as seen in some activity diagnostic). There is a relationship between cycle amplitude and period seen in the Sun, where shorter cycles tend to be followed by stronger ones (e.g., Hathaway, Wilson \\& Reichmann 1994), but there has been little work along these lines in stars. Focusing on inactive stars, Soon et al. (1994) suggested that the fractional peak-to-peak cycle amplitude seen in chromospheric Ca {\\sc II} HK emission,  \\noindent $A_{\\rm cyc} = \\Delta R'_{\\rm HK}/\\langle R'_{\\rm HK}\\rangle$,  \\noindent decreased linearly with $\\log (\\Omega/\\omega_{\\rm cyc})^2$. Studying a larger sample, Baliunas et al. (1996) found that among inactive stars, $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.9}$. Results for active stars were less clear.   Baliunas et al. (1996) analyzed {\\it all} cycles detected by Bea95, independent of their quality. In our earlier work on cycle periods we found it very useful to begin by taking only a critically chosen sample of reliable, high grade cycles (e.g., Bradenburg et al. 1998, SB). In this paper, we therefore take a fresh look at cycle amplitudes, in light of the new results for $P_{\\rm cyc}$, using the same high quality cycle sample. ",
        "conclusions": "We first investigated how $\\Delta R'_{\\rm HK}$ depends on $\\langle R'_{\\rm HK}\\rangle$ itself (Fig. 1). We find $\\Delta R'_{\\rm HK} \\propto \\langle R'_{\\rm HK}\\rangle^{0.77}$ ($\\sigma = 0.17$ dex; fitting the primary $P_{\\rm cyc}$ only). This implies that $A_{\\rm cyc} \\propto \\langle R'_{\\rm HK}\\rangle^{-0.23}$ -- fractional cycle amplitudes {\\it decrease} somewhat with increasing activity.   Some of the scatter about the fit can be explained by an additional dependence on $B-V$ color: $\\Delta R'_{\\rm HK}$ is larger in K stars than G or F at fixed $R'_{\\rm HK}$.  This is seen more clearly by plotting $A_{\\rm cyc}$ vs. $B-V$ color (Fig. 2), which shows an steady increase in  $A_{\\rm cyc}$ from F through G stars, until a maximum is reached by the mid K stars ($B-V \\sim 1.0$). The averages by spectral type are $\\langle A_{\\rm cyc}$(F)$\\rangle = 0.17$, $\\langle A_{\\rm cyc}$(G)$\\rangle = 0.22$, and $\\langle A_{\\rm cyc}$(K)$\\rangle = 0.35$.  This behavior may be the result of a dependence of cycle amplitude on fractional convection zone depth, up to some limiting value in mid-K stars.  Some of the scatter in Fig. 1 is also intrinsic.  If stellar cycles are like the Sun, activity will be restricted in latitude. Stars with different inclinations $i$ will then exhibit different apparent $A_{\\rm cyc}$ (Radick et al. 1998; Knaack et al. 2001). Analysis of the models in Knaack et al. (2001), however, indicates that 77\\% of the measured $A_{\\rm cyc}$ should lie within $\\pm$15\\% of the mean, with only 23\\% (those with $i \\leq 39^\\circ$) will range from -15\\% to -46\\% of $\\langle A_{\\rm cyc}\\rangle$. Furthermore, stars more active than the Sun are expected to have a wider latitude distribution of activity (e.g., Schrijver \\& Title 2001), and thus should show less scatter due to varying $i$. Thus the effects of $i$ should be relatively small, and strongest stars with $R'_{\\rm HK} \\la R'_{\\rm HK}(\\odot)$.       \\begin{figure}[ht] \\psfig{file=drhk_rhk4.ps,height=5.5cm,width=8cm} \\caption{$\\Delta R'_{\\rm HK}$ vs. $\\langle R'_{\\rm HK}\\rangle$ for the stellar sample; symbols indicate F ($\\triangle$), G ($\\circ$, the Sun's is doubled), and K (box) stars, filled symbols are more active stars, size $\\propto$ $P_{\\rm cyc}$ and $A_{\\rm cyc}$ ``quality\",  dotted lines connect secondary $P_{\\rm cyc} (\\times)$ in some stars.  We find $\\Delta R'_{\\rm HK} \\propto \\langle R'_{\\rm HK}\\rangle^{0.77}$ (solid line), implying $A_{\\rm cyc} \\propto \\langle R'_{\\rm HK}\\rangle^{-0.23}$.} \\label{figlabel} \\end{figure}  \\begin{figure}[ht] \\psfig{file=amp_bv4.ps,height=5.5cm,width=8cm} \\caption{$A_{\\rm cyc} = \\Delta R'_{\\rm HK}/\\langle R'_{\\rm HK}\\rangle$ vs. $B-V$ (symbols as in Fig. 1), showing an increasing $A_{\\rm cyc}$ with decreasing effective temperature in F and G stars, reaching a maximum in mid-K stars.} \\label{figlabel} \\end{figure}  In contrast to the dependence on $B-V$, $A_{\\rm cyc}$ shows little dependence on rotation, whether expressed as $\\Omega$, or (see Fig. 3) as $Ro^{-1} = \\tau_C/P_{\\rm rot}$ (where $\\tau_C$ is the convective turnover timescale; Noyes et al. 1984). Since $\\langle R'_{\\rm HK}\\rangle$ shows a clear dependence on $Ro^{-1}$ (Noyes et al. 1984), this result implies $\\Delta R'_{\\rm HK}$ has a similar dependence on rotation.  Indeed, for $Ro^{-1} \\leq 1.5$, $\\Delta R'_{\\rm HK} \\propto Ro^{-1}$  (above this $Ro^{-1}$, there is evidence for cycle amplitude saturation).  Since $A_{\\rm cyc}< 1$, one might infer that only portion of $R'_{\\rm HK}$ derives from a cycling dynamo.  This suggests that the non-cycling (small-scale) component of the dynamo is prominent and thus (from Fig. 3) must have a similar dependence on rotation as the cycling component. However, this is at odds with the lack of any strong rotational dependence for $R'_{\\rm HK}$ in very inactive, ``flat\" activity stars where activity is likely fueled by a small-scale dynamo alone (Saar 1998). Thus, a more likely possibility is that there is significant temporal overlap in $R'_{\\rm HK}$ between cycles, reducing the {\\it apparent} amplitude $A_{\\rm cyc}$.  In this scenario, the cyclic dynamo dominates the rotational dependence of $R'_{\\rm HK}$, despite the apparently small $A_{\\rm cyc}$. Enhanced cycle overlap with increasing $\\langle R'_{\\rm HK}\\rangle$ would also explain the decrease in $A_{\\rm cyc}$ with $\\langle R'_{\\rm HK}\\rangle$ (Fig. 1).   \\begin{figure}[ht] \\psfig{file=amp_ro4.ps,height=5.5cm,width=8cm} \\caption{$A_{\\rm cyc} = \\Delta R'_{\\rm HK}/\\langle\\Delta R'_{\\rm HK}\\rangle$ vs. $Ro^{-1} = \\tau_C/P_{\\rm rot}$ (symbols as in Fig. 1), showing no clear relationship.} \\label{figlabel} \\end{figure}  Study of the dependence of $A_{\\rm cyc}$ on $P_{\\rm cyc}$ with our dataset reveals a complex situation.  Most inactive stars (defined as $\\log R'_{\\rm HK} \\leq -4.75$), combined with a few active ones, trace out a relation similar to that found by Baliunas et al. (1996), while most active ones show a similar, nearly parallel relation offset at higher $A_{\\rm cyc}$. Specifically, 20 stars (14, or 70\\% of them inactive) show $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.66}$ ($\\sigma = 0.082$ dex), and 9 stars (7, or 78\\% of them active) exhibit $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.85}$ ($\\sigma = 0.081$ dex). Only two lower quality detections cannot be assigned to one of these branches. This ``branched\" structure is similar to that seen between $\\omega_{\\rm cyc}/\\Omega$ and $Ro^{-1}$ or $R'_{\\rm HK}$ by Brandenburg et al. (1998), SB, and Saar \\& Brandenburg (2001) (and first noted by Saar \\& Baliunas 1992 and Soon et al. 1993). There are some differences in branch membership, though; for example, active G stars HD 1835, 20630, and 26913 lie on the active (``A\") branch in SB, but the inactive (``I\") branch here (Figs. 4, 5).     \\begin{figure}[ht] \\psfig{file=amp_oo4w_nop2.ps,height=5.5cm,width=8cm} \\caption{$A_{\\rm cyc} = \\Delta R'_{\\rm HK}/\\langle\\Delta R'_{\\rm HK}\\rangle$ vs. $\\omega_{\\rm cyc}/\\Omega$ (symbols as in Fig. 1), showing an inactive (marked ``I\"; $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.66}$) and an active (marked ``A\"; $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.85}$) branch. } \\label{figlabel} \\end{figure}   If we also consider secondary $P_{\\rm cyc}$, we find three lie on these branches, while four do not.  Including the secondary $P_{\\rm cyc}$ in the branch fits  causes little change: for the ``I\" branch, $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.68}$ ($\\sigma = 0.082$ dex) while for the ``A\" branch, $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.88}$ ($\\sigma = 0.076$ dex). Four of the stars show $A_{\\rm cyc}$(primary)/$A_{\\rm cyc}$(secondary) = 2.26$\\pm$0.18, suggesting there may be some ``preferred\" amplitude ratios. All four secondary cycles ``unmatched\" to a branch lie at low $A_{\\rm cyc}$ and high $\\omega_{\\rm cyc}/\\Omega$ and may indicate a third branch -- more data is needed to confirm this.  Thus, it appears that $A_{\\rm cyc}$ and $P_{\\rm cyc}$ are related in stars, just as they are in the Sun. Study of stars indicates the relationship is multivalued, and depends on rotation.    \\begin{figure} \\psfig{file=amp_oo4w.ps,height=5.5cm,width=8cm} \\caption{Same as Fig. 4, but including secondary $P_{\\rm cyc}$. An inactive (marked ``I\"; $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.68}$) and an active (marked ``A\"; $A_{\\rm cyc} \\propto (\\omega_{\\rm cyc}/\\Omega)^{0.88}$) branch are indicated. } \\label{figlabel} \\end{figure}     The present analysis is limited due to the saturation of Ca {\\sc II} emission for $Ro^{-1} \\ga 3$, which likely suppresses and obscures the visibility of cycles for more active stars.  Cycle overlap is also a concern. A logical next step would be to investigate photometric cycle amplitudes, which in active stars are due to spots (rather than the plage/network seen in Ca {\\sc II}). Photometric cycle amplitudes saturate at considerably higher $Ro^{-1}$ (Messina et al. 2001), permitting study of $A_{\\rm cyc}$ in much more active, faster rotating stars. Ideally, some normalized  quantity like the fractional luminosity amplitude $\\Delta L/L$ or the starspot filling factor $f_S$ (e.g., from TiO measurements) should be used."
    },
    "0207/astro-ph0207501_arXiv.txt": {
        "abstract": "{We have developed a Monte Carlo method to compute the luminosity function of galaxies, based on photometric redshifts, which takes into account the non-gaussianity of the probability functions, and the presence of degenerate solutions in redshift.  In this paper we describe the method and the mock tests performed to check its reliability.  The NIR luminosity functions and the redshift distributions are determined for near infrared subsamples on the HDF-N and HDF-S.  The results on the evolution of the NIR LF, the stellar mass function, and the luminosity density, are presented and discussed in view of the implications for the galaxy formation models.  The main results are the lack of substantial evolution of the bright end of the NIR LF and the absence of decline of the luminosity density up to a redshift $z \\sim 2$, implying that most of the stellar population in massive galaxies was already in place at such redshift. ",
        "introduction": "The study of galaxy formation and evolution implies the availability of statistical samples at large look-back times. At large redshifts, though, only star forming galaxies will be entering the samples obtained in the visible bands and, to be able to probe their stellar masses, observations at longer wavelengths are needed.  In the last decade, deep photometric galaxy samples have become available, namely through the observations by HST of the HDF-N (Williams et al.\\ \\cite{williams}) and HDF-S (Casertano et al.\\ \\cite{casertano}), which have been coordinated with complementary observations from the ground at near-infrared (NIR) wavelengths (Dickinson et al.\\cite{dick}; da Costa et al.\\ \\cite{dacosta}).  At the same time, reliable photometric redshift techniques have been developed, allowing estimates of the distances of faint galaxies for which no spectroscopic redshifts can be obtained nowadays, even with the most powerful telescopes (e.g.\\ Connolly et al. \\ \\cite{con}; Wang et al.\\ \\cite{wan}; Giallongo et al.\\ \\cite{gia}; Fern\\'andez-Soto et al.\\ \\cite{fsoto}; Arnouts et al.\\ \\cite{arnouts}; Furusawa et al. \\cite{furu}; Rodighiero et al.\\ \\cite{rodighiero}; Le Borgne \\& Rocca-Volmerange \\cite{leborgne}; Bolzonella et al.\\ \\cite{hyperz} and the references therein).  One of the main issues for photometric redshifts is to study the evolution of galaxies beyond the spectroscopic limits. The relatively high number of objects accessible to photometry per redshift bin allows to enlarge the spectroscopic samples towards the faintest magnitudes, thus increasing the number of objects accessible to statistical studies per redshift bin. Such slicing procedure can be adopted to derive, for instance, redshift distributions, luminosity functions in different bands, or rest-frame colours as a function of absolute magnitudes, among the relevant quantities to compare with the predictions derived from the different models of galaxy formation and evolution. This approach has been recently used to infer the star formation history at high redshift from the UV luminosity density, to analyse the stellar population and the evolutionary properties of distant galaxies (e.g.\\ SubbaRao et al.\\ \\cite{subba}; Gwyn \\& Hartwick \\cite{gwyn}; Sawicki et al.\\ \\cite{sawicki}; Connolly et al.\\ \\cite{con}; Pascarelle et al.\\ \\cite{pascarelle}; Giallongo et al.\\ \\cite{gia}; Fern\\'andez-Soto et al.\\ \\cite{fsoto}; Poli et al.\\ \\cite{poli}), or to derive the evolution of the clustering properties (Arnouts et al.\\ \\cite{arnouts}; Magliocchetti \\& Maddox \\cite{maglio}, Arnouts et al.\\ \\cite{arnouts2}).  We have developed a method to compute luminosity functions (hereafter LFs), based on our public code \\emph{hyperz} to determine photometric redshifts (Bolzonella et al.\\ \\cite{hyperz}).  This original method is a Monte Carlo approach, different from the ones proposed by SubbaRao et al. (\\cite{subba}) and Dye et al. (\\cite{dye}) in the way of accounting for the non-gaussianity of the probability functions, and specially to include degenerate solutions in redshift.  In this paper we present the method and the tests performed on mock catalogues, and we apply it specifically to derive NIR LFs and their evolution on the HDF-N and HDF-S.  The NIR luminosity is directly linked to the total stellar mass, and barely affected by the presence of dust extinction or starbursts. According to Kauffmann \\& Charlot (\\cite{kauff}), the NIR LF and its evolution constitute a powerful test to discriminate between the different scenarios of galaxy formation, i.e.\\ if galaxies were assembled early, according to a monolithic scenario, or recently from mergers.  The theoretical NIR LFs derived by Kauffmann \\& Charlot exhibit a sharp difference between the two models at redshifts $z>1$. The PLE models foresee a constant bright-end for the LF, whereas hierarchical models are expected to undergo a shift towards faint magnitudes with increasing redshift.  In this paper we compare the theoretical predictions with the observations on the HDFs, in order to extend the analysis performed from spectroscopic surveys (Glazebrook et al.\\ \\cite{glaze}, Songaila et al.\\ \\cite{songaila}, Cowie et al.\\ \\cite{cowie2}) up to $z \\sim 2$.  The comparison between the present NIR LF results and a similar study in the optical-UV bands, obtained with the same photometric redshift approach (Bolzonella et al.\\ Paper II, in preparation), could provide with new insights on the galaxy formation scenario.  The plan of this paper is the following. Section~\\ref{catalogue} gives a brief description of the photometric catalogues used to select the samples, their properties being discussed in Sect.~\\ref{sample}.  In Sect.~\\ref{lf} we describe the technique conceived to compute the luminosity functions using the \\emph{hyperz} photometric redshift outputs, and the test of the method through mock catalogues. The results obtained on the HDFs near-infrared LF are presented and discussed in Section~\\ref{resu}, together with the Luminosity Density and the Mass Function derived from them. The implications of the present results on the galaxy formation models are discussed in Sect.~\\ref{discuss}, and the main results are summarized in Sect.~\\ref{sum}.  Throughout this paper we adopt the cosmological parameters $\\Omega_0 = 1$, $\\Omega_\\Lambda = 0$, when not differently specified.  Magnitudes are given in the AB system (Oke \\cite{oke}).  Throughout the paper, the Hubble Space Telescope filters F300W, F450W, F606W and F814W are named $U_{300}$, $B_{450}$, $V_{606}$ and $I_{814}$ respectively. ",
        "conclusions": "\\label{discuss}   LFs are a convenient way to describe the galaxy population and to get hints about the mechanisms of formation and evolution of galaxies.  Two scenarios are in competition to explain the history of galaxies up to the present epoch. The formation of elliptical galaxies is especially intriguing, because despite their old and apparently simple stellar population, their process of formation is far from being understood.  The two models in competition are: \\begin{itemize} \\item The hierarchical scenario, which is based on CDM cosmological models (e.g. Kauffmann et al.\\ \\cite{kauff1}; Cole et al.\\ \\cite{cole}). It assumes that galaxies have formed through merging of sub-galactic structures. In particular, elliptical galaxies are born from the merging of two disks of comparable dimensions; during the interaction, a burst of star formation occurs, consuming the gas present in the progenitors, and then passive evolution follows.  \\item The monolithic scenario, revisited in different ways by many authors, which basically assumes that galaxies formed approximately at the same epoch, with particular details depending on the morphological types (e.g.\\ Rocca-Volmerange \\& Guiderdoni \\cite{rocca}; Pozzetti et al.\\ \\cite{pozz}).  Ellipticals are believed to form at $z_{\\rm form} \\ga 2-3$ and, at the time of their assembly, a burst of star formation occurred, followed by passive evolution of the stellar population. \\end{itemize}  The two scenarios foresee different characteristics as a function of redshift for the progenitors of the blue and red galaxies in the local universe.  Moreover, since most of the merging would take place at $z<2$, the redshift range between $1$ and $2$ is fundamental to discriminate between monolithic and hierarchical scenarios.  Both mechanisms account for the properties of elliptical galaxies up to $z\\simeq1$, e.g.\\ the CM relation (Gladders et al.\\ \\cite{glad}; Kauffmann \\& Charlot \\cite{kauff2}), but, between $z=1$ and $z=2$, the expectations for the photometric properties of all galaxies differ significantly between the two scenarios.  A powerful test to establish what drives the galaxy evolution, was proposed by Kauffmann \\& Charlot (\\cite{kauff}), using the $K$-band luminosity function.  The luminosity in the $K$ filter is directly linked to the mass in stars and is barely affected by the presence of dust extinction, thus making the $K$ photometry a privileged tool to study galaxy formation. Moreover, galaxies with the same stellar mass have nearly the same $K$ magnitude, independently on their star formation history. For these reasons the $K$-band LF can probe if galaxies were assembled early, according to the monolithic scenario, or recently from mergers.  Kauffmann \\& Charlot (\\cite{kauff}) built two PLE models with density parameters $\\Omega_0=1$ and $\\Omega_0=0.2$ and two hierarchical models based on the CDM cosmology, with the same values of $\\Omega_0$.  They computed the evolution of the $K$-band LF at increasing redshifts for the PLE models and for the $\\Omega_0=1$ hierarchical model, each one able to reproduce the local LF.  On the contrary, their low density hierarchical model failed to reproduce the local $K$-band LF and they did not compute the evolution of the LF in this framework.  At redshifts $z>1$ a sharp difference between the two models is predicted, as explained in Kauffmann \\& Charlot.  The PLE models foresee a constant bright-end for the LF, with small differences between the flat and the open cosmology, whereas the hierarchical model considered by Kauffmann \\& Charlot undergoes a shift toward faint magnitudes (see Fig.~\\ref{lfk_kc}).  The development of hierarchical scenarios in open or flat cosmological models with cosmological constant mitigates the discrepancy between these two scenarios (monolithic vs. hierarchical), because the epoch of the major merging moves at higher redshifts (Fontana et al.\\ \\cite{fontana}, Cole et al.\\ \\cite{cole}).  In this case, the PLE model could not be distinguished from a scenario where the galaxy assembly occurs at early epochs, followed by a passive evolution of the stellar population. Our results for the cumulative redshift distribution actually support such scenarios (see Fig.~\\ref{nzcum}).  \\begin{figure} \\centerline{\\psfig{file=lfK_kc2bw_new.ps,width=0.49\\textwidth}} \\caption{Comparison between the theoretical luminosity functions derived by Kauffmann \\& Charlot (\\cite{kauff}) and the present results on the HDFs.  Histograms represent the prediction for the hierarchical, $\\Omega_0=1$ CDM, (dashed line) and PLE scenarios (solid line).  The solid and long dashed continuous lines display respectively the STY LFs fitted to the data in the HDF-N and HDF-S. Circles and triangles show the estimate of the LF with the $C^-$ method in the HDF-N and HDF-S.  In the lower panel we show the model predictions both at $z=1$ and $z=2$, because we estimated the LF using galaxies with redshifts between these two extremes.} \\label{lfk_kc} \\end{figure}  The present results, i.e.\\ the lack of significant evolution in the bright part of the LF from the redshift range $[0,1]$ to $[1,2]$, provide a stringent clue, supporting the idea that massive galaxies were already in place at high redshifts, against the old CDM hierarchical model adopted by Kauffmann \\& Charlot.  The comparison between these theoretical predictions and the observations derived in the present paper can be found in Fig.~\\ref{lfk_kc}.  In the lower redshift bin, our estimates of the LFs present a faint-end slope in agreement with the hierarchical model, whereas at bright magnitudes the small differences between the two models do not support any claim. In the redshift bin $z=[1,2]$ we have compared our estimate with both the predictions of the models at redshift $z=1$ and $2$: the model predictions suitable for our sample should lie between the two.  It is evident from the lower panel of Fig.~\\ref{lfk_kc} that the very bright part of the LFs remains well above the hierarchical model predictions at $z=1$ and $2$, being much closer to the predictions of the monolithic/PLE-like scenario.  On the other hand, the faint end slope seems to be in better agreement with the hierarchical model, even in this redshift range.  Other recent studies on the Hubble fields, based on different selection criteria, seem to indicate that the formation of elliptical galaxies should be placed at $z>2$ (Ben\\'{\\i}tez et al.\\ \\cite{benitez}; Broadhurst \\& Bouwens \\cite{broad}).  In the present paper we do not select galaxies according to morphological types, but the same conclusions apply to the most massive (NIR luminous) galaxies in our sample.  The lack of evolution in the bright end of the LF is in good agreement with the results found from spectroscopic surveys.  Glazebrook et al.\\ (\\cite{glaze}) found no evidence for evolution in the $K$-band LF up to $z \\le 0.5$, and concluded that massive spheroids were in place at $z \\ge 1$ and then evolved passively.  Songaila et al.\\ (\\cite{songaila}) found a lack of significant evolution in their $K$-band sample up to a redshift of $\\sim 1$.  Cowie et al.\\ (\\cite{cowie2}) found little evolution in their sample of red (old) objects to $z \\sim 1$.  The present results extend the previous findings in redshift, up to $z \\sim 2$, with the same conclusions with respect to the evolution of the most massive galaxies.  The comparison between the present LFs in the near-IR and in the optical-UV bands (Bolzonella et al.\\ in preparation) will provide new insights on the galaxy formation scenario."
    },
    "0207/astro-ph0207070_arXiv.txt": {
        "abstract": "Carbon monoxide and ammonia have been detected in the spectrum of Gl 229B at abundances that differ substantially from those obtained from chemical equilibrium.  Vertical mixing in the atmosphere is a mechanism that can drive slowly reacting species out of chemical equilibrium.   We explore the effects of vertical mixing as a function of mixing efficiency and effective temperature on the chemical abundances in the atmospheres of brown dwarfs and on their spectra.  The models compare favorably with the observational evidence and indicate that vertical mixing plays an important role in brown dwarf atmospheres. ",
        "introduction": "The discovery of strong methane bands in the spectrum of Gl 229B indicated at once that it has a very low effective temperature ($\\sim 1000\\,$K) and it was hailed as the first clearly identified brown dwarf (Nakajima et al. 1995; Oppenheimer et al. 1995). This follows from the chemistry of carbon in stellar atmospheres. In cool M dwarfs, carbon is predominantly in the form of carbon monoxide (CO), while in Jovian planets methane (CH$_4$) is the dominant form.  The transition occurs in the brown dwarf regime, at $T \\sim 1100\\,$K for a pressure of 1 bar (Fegley \\& Lodders 1996).  The presence of methane bands is so striking that it prompted the creation of the new spectral class of T dwarfs (Kirkpatrick et al. 1999).  Gl 229B is a T6 dwarf. In this context, the discovery of a strong CO band in the 4.5 -- 5$\\,\\mu$m spectrum of Gl 229B by Noll, Geballe \\& Marley (1997) is quite remarkable.  Analysis of the spectrum reveals an abundance of CO that is over 3 orders of magnitude larger than expected from chemical equilibrium calculations (Noll et al. 1997; Griffith \\& Yelle 1999; Saumon et al. 2000), as anticipated by Fegley \\& Lodders (1996).  A similar situation is observed in the atmosphere of Jupiter where CO is detected while chemical equilibrium overwhelmingly favors CH$_4$ as the reservoir of carbon (Prinn \\& Barshay 1977).  The excess of CO in Gl 229B can be understood if the chemistry of carbon is driven away from equilibrium by vertical mixing on a time scale shorter than the rate of conversion of CO into CH$_4$ (Fegley \\& Lodders 1996; Noll et al 1997; Griffith \\& Yelle 1999).  The two reservoirs of nitrogen in brown dwarfs are molecular nitrogen (N$_2$) and ammonia (NH$_3$). Ammonia is the dominant molecule at low temperatures.  While N$_2$ is invisible in the near infrared, Saumon et al. (2000) detected NH$_3$ features in Gl 229B in the $K$ band and inferred its presence in the $H$ band.  Each detection corresponds to a different level in the atmosphere. NH$_3$ abundances determined separately for each band reveal a value in agreement with chemical equilibrium in the $H$ band and a {\\it depletion} by at least a factor of 4 in the $K$ band.  These results can be interpreted consistently with the CO observation by invoking vertical mixing in the atmosphere on a time scale shorter than the rate of conversion of N$_2$ into NH$_3$.  The modeling of non-equilibrium chemistry of CO and other molecules due to mixing in planetary atmospheres is a mature field (e.g. Barshay \\& Lewis 1978; Fegley \\& Prinn 1985; Fegley \\& Lodders 1994; Lodders \\& Fegley 1994).  Application to brown dwarfs have so far been limited to modeling the excess of CO in Gl 229B (Griffith \\& Yelle 1999) and a detailed study of the kinetics of CNO chemistry (Lodders \\& Fegley 2002). Here we present the first systematic exploration of the effects of mixing on the chemistry of carbon and nitrogen in brown dwarfs atmospheres.  We first summarize the main features of non-equilibrium chemistry of carbon and nitrogen caused by vertical mixing. Using a simple parametrization for the rate of mixing, we then compute the resulting non-equilibrium abundances of CO, CH$_4$, N$_2$, NH$_3$ and H$_2$O as a function of effective temperature and mixing efficiency.  The new abundance profiles lead to significant changes in the emergent spectra of cool brown dwarfs. ",
        "conclusions": "Motivated by observations in the T6 dwarf Gl 229B of a strong CO band where none was expected and of a depletion of NH$_3$, we have explored the effects of vertical mixing in brown dwarf atmospheres on the most prominent molecules visible in their spectra.  Vertical mixing, by convection or eddy turbulence, can occur fast enough to keep abundances of CO, CH$_4$, H$_2$O and NH$_3$ from reaching their values at chemical equilibrium. Increased efficiency of mixing leads to an enrichment of CO and a depletion of CH$_4$, H$_2$O, and NH$_3$.  The model of vertical mixing can explain the CO and NH$_3$ observations in Gl 229B self-consistently.  Furthermore, observations of NH$_3$ in three different bandpasses can test the model prediction of a flat abundance profile throughout the atmosphere (Saumon et al. 2002).  In this study, the efficiency of mixing is treated as a free parameter but it can eventually be calibrated with data.  New observational evidence suggests that the overabundance of CO observed in Gl 229B may be a common feature of late L dwarfs and of T dwarfs.  If this is indeed the case, it becomes essential to include non-equilibrium chemistry in modeling the atmospheres and spectra of cool brown dwarfs.  Because it affects the relative abundances of major constituents of the atmosphere, vertical mixing further complicates the determination of the composition of brown dwarfs.  Finally, vertical mixing can result in a suppression of the $M^\\prime$ flux by up to one magnitude for $\\Teff$ between $\\sim 700$ and 1100$\\,$K (Fig. 4).  This must be taken into account in plans to image extrasolar giant planets by capitalizing on their predicted surperthermal $M$ band flux (Marley et al. 1996, Burrows et al. 1997).  \\begin{figure}[t] \\plotfiddle{figure4.ps}{7.8cm}{0}{67}{67}{-210}{-155} \\caption{Absolute $M^\\prime$ (MKO) magnitude versus effective temperature for L and T dwarfs. The data are from Leggett et al. (2002) and Golimowski et al. (2002). Open squares show 2MASS $0559-14$ as a single star (upper point) and as an equal pair binary (lower point).  Curves show sequences of models with $g=10^5\\,$cm/s$^2$ for different values of the eddy diffusion coefficient.  From top to bottom, $K=0$, 10$^2$, $10^4$, and $10^6\\,$cm$^2$/s, respectively.   $K=0$ corresponds to the equilibrium case.} \\end{figure}"
    },
    "0207/astro-ph0207593_arXiv.txt": {
        "abstract": "Cosmological gamma-ray bursts may be powered by rotating black holes with contemporaneous emission of gravitational radiation from a surrounding torus. We calculate the resulting stochastic background radiation assuming strong cosmological evolution and a uniform black hole mass distribution of $M=$ (4--14)M$_\\odot$. The predicted spectral flux density corresponds to a peak spectral closure density of (1--2$)\\times 10^{-7}$, and has comparable contributions at 450 Hz$\\times\\kappa$ and over 300--450 Hz$\\times\\kappa$ from nearby and distant sources, respectively, where $\\kappa$ refers to an uncertainty factor of order unity in the radius of the torus. For two optimized advanced LIGO-type detectors the proposed gravitational wave background could be detectable within a year of integration. ",
        "introduction": "Cosmological gamma-ray bursts are the most enigmatic transient events in the Universe. They show a bi-modal distribution in durations of short bursts around 0.3 s and long bursts around 30 s \\citep{kou93}. Based on their GRB fluence and time-variability, their inner engines should be compact and highly energetic. Leading candidates for GRB progenitors are collapsars and mergers of black holes and neutron stars. In particular, long bursts have been associated with core collapse of massive stars \\citep{woos93,Mwoos99}, or their hypernova variants, in star-forming regions \\citep{Paczy97,Brown2000}. These, and mergers of compact binaries, are believed to result in black hole plus disk or torus systems --- see van Putten (2001) for a review.  A torus around a rapidly rotating black hole converts spin energy into various channels, notably gravitational and thermal radiation, winds and MeV neutrino emissions \\citep{mvp02}. These emissions last for the life-time of rapid spin of the black hole -- the de-redshifted durations of tens of seconds for long bursts from black hole-torus systems in suspended accretion \\citep{mvpostr}. The single-source spectrum is here described by a horizontal branch in the $\\dot{f}(f)$ diagram \\citep{mvp00} within a frequency range in the vicinity of 1 kHz for a 7-M$_{\\odot}$ black hole. GRB energies appear to have a diversity of about one order of magnitude \\citep{Frail2001,Piran2001}. Black hole-torus systems are expected to have a distribution in black hole mass $M$ consistent with the recently proposed association with soft X-ray transients in the hypernova proposal of \\citet{Brown2000}; i.e., $M\\simeq$ (4--14)M$_{\\odot}$.  In this work, we calculate the gravitational wave (GW) spectra expected from a cosmological distribution of black hole-torus systems, assuming strong cosmological evolution locked to the star-formation rate (SFR). We shall estimate the expected contribution to the stochastic background in GWs from the low- and high-redshift populations. Following \\citet{scm01} and \\citet{Frail2001}, the total event rate is normalized to a local GRB rate of 0.5 yr$^{-1}$ Gpc$^{-3}$ at $z=0$, assuming a ``flat-$\\Lambda$\" cosmology. ",
        "conclusions": "\\noindent We draw several conclusions:  \\noindent 1. Black hole-torus systems associated with GRBs are expected to give a substantial contribution to the stochastic GW background in a frequency window of 300-450 Hz$\\times\\kappa$, where $\\kappa$ denotes an uncertainty factor of order unity in the radius of the torus. If black hole-torus systems exist also as transient sources independent of GRBs, their event rate will be higher with a commensurably more pronounced stochastic GW background.  \\noindent 2. It is instructive to compare the presented results with radiation from rapidly rotating neutron stars. Black hole-torus systems produce spectral flux densities which are similar to the contribution from r-mode instabilities as described in Ferrari et al. (1999b). Here, the large output provided by the spin-energy of the black hole compensates for an event rate which is less than the formation rate of neutron stars by some three orders of magnitude. The frequencies and $DC$ of the former are markedly higher, respectively, lower than those predicted for radiation from neutron star modes, however. The $DC$ of the signals from neutron star r-modes may be as high as $10^{9}$ (Ferrari et al. 1999b) in contrast to our $DC$ estimate of order unity for the signals from GRB sources. This comparision becomes more favorable with recent understanding of r-modes in more detailed and realistic scenarios. The r-modes are driven by a generally weak gravitational radiation-reaction force and, in the perturbative limit, effectively decoupled from other modes \\citep{sch02}. A recent appreciation of various channels for creating viscosity \\citep{owen02,rez01,wu01} renders their saturation energies small, and less so than previously thought \\citep{arra02}.  \\noindent 3. The high output in gravitational radiation powered by the spin-energy of the black hole gives rise to a major contribution to the predicted GW background by a relatively nearby population of sources. This is apparent in the edge at 450 Hz$\\times\\kappa$ in Fig. 4 to the 300--450 Hz$\\times\\kappa$ flat-$\\Lambda$ spectral flux density plateau due to sources at higher redshifts. The comparable contributions by the nearby and distant sources is a robust result, depending only on the assumption that the proposed black hole-torus systems are tightly locked to the SFR.  \\noindent 4. As the local GRB rate (assumed in this work) is about 1 yr$^{-1}$ within a radius of 100 Mpc, the average background spectral flux density should be representative of sources from $z=$ 0.02--5 for an observation time of one year. One can view the background spectral flux density shown in Fig. 4 as the result of averaging the flux from a discrete number distribution over an infinite observation time. We plan to investigate these issues further in connection with the detectability of this predicted background using advanced LIGO-type detectors.  \\noindent 5. The estimated frequency range of about 200--2000 Hz of the proposed background and foreground radiation in gravitational waves from black hole-torus systems in association with GRBs defines a new source with a well-defined event rate in the high-frequency bandwidth of LIGO/VIRGO detectors as well as for some of the current bar detectors. Cross-correlating the output between pairs of detectors promises to be the optimum detection strategy. For an idealized advanced LIGO pair of detectors, the $S/N$ could be high enough for detection of the proposed background in a year of integration. Using present detector technology and locations, the optimum detector pair combination may be a resonant-mass and interferometric detector, such as the LIGO-WA and ALLEGRO pair."
    },
    "0207/astro-ph0207185_arXiv.txt": {
        "abstract": "\\noindent We present the mass functions for different mass estimators for a range of cosmological models.  We pay particular attention to how universal the mass function is, and how it depends on the cosmology, halo identification and mass estimator chosen.  We investigate quantitatively how well we can relate observed masses to theoretical mass functions. ",
        "introduction": "One of the most fundamental predictions of a theory of structure formation is the number density of objects of a given mass, the mass function.  Accurate mass functions are used in a number of areas in cosmology; in studies of galaxy formation, in measures of volumes (e.g.~galaxy lensing) and in attempts to infer the normalization of the power spectrum, the statistics of the initial density field, the density parameter or the equation-of-state of the dark energy from the abundance of rich clusters.  One of the most intriguing aspects of the mass function is that it appears universal, in suitably scaled units, for a wide range of theories. A complete understanding of this phenomenon currently eludes us.  If we are to attempt to use the mass function to infer cosmological parameters from the abundance of objects of some given property, then we need to understand how accurate our theory for the mass function is. This involves understanding how to define the mass of an object in cosmology, a task which is non-trivial as there is no clear boundary between a halo and the surrounding large-scale structure in theories of hierarchical structure formation. The purpose of this paper is to calculate the mass functions for different mass estimators for a range of cosmological models, to see how well these statistics can be computed from semi-analytic theories and to investigate quantitatively how to relate observed masses to theoretical mass functions.  We find that the mass function is only approximately `universal', and only for a few mass estimators. We discuss how accurately one can convert between different mass estimates, so as to relate what can easily be measured to what can easily be predicted. We also discuss the limitations which non-universality of the mass function would place on cosmological parameter estimation were it not to be corrected for.  The outline is as follows: after a review of some background (\\S\\ref{sec:pressschechter}) we present mass functions, derived from N-body simulations (\\S\\ref{sec:simulations}), for a variety of different mass definitions (\\S\\ref{sec:massdef}).  We compare these mass definitions, based on mean density contrast, with the concept of a virialized halo (\\S\\ref{sec:virial}). We investigate how universal the mass function is (\\S\\ref{sec:universal}) for each different estimator and present fitting functions (\\S\\ref{sec:fitting}) to the mass functions for 3 different cosmological models.  We finish by considering the effect of clustering on the mass function (\\S\\ref{sec:clustering}) and summarize our main results (\\S\\ref{sec:conclusions}). ",
        "conclusions": "\\label{sec:conclusions}  The multiplicity function, a measure of the number of halos per comoving volume element per unit mass, is one of the central predictions of a model of structure formation. Dark matter dominated models in which structure evolves hierarchically from gaussian initial conditions predict a mass function which is nearly universal if expressed in the right units.  This is only true for a narrow class of mass estimators, and is specifically not true for the estimators which have been most commonly used up until now.  Given the complicated process by which halos form in hierarchical models, the role of mergers and prevalence of sub-structure, it is highly convenient that the mass function (in scaled units) is so close to universal. We have provided fitting functions to the mass function from N-body simulations for 8 different mass estimators, and shown how one can convert between them. We have found that measuring the mass of a halo using one definition and using a simple average spherical profile, such as the NFW profile, to convert to the `universal' $M_{180b}$ provides a remarkably good method of estimating the mass function, even though individual halos show a large scatter among different mass estimates.  Finally let us remark upon the small non-universality in the multiplicity function, which can lead to misestimates of the true mass function if one uses a fitting function like Eq.~(\\ref{eqn:fnu}). Neglecting the factor $d\\log\\sigma/d\\log M$ in the mass function, making an error of $\\delta n/n$ in the number density per $\\log M$ at mass $M$ translates into an error of $(\\nu^2-1)^{-1} \\delta n/n$ in the normalization $\\delta\\sigma/\\sigma$. So for a typical cluster, with $\\nu\\sim 2-3$, the scatter in the mass function doesn't limit our knowledge of $\\sigma_8$ until the other uncertainties are pushed below ${\\cal O}(5\\%)$. Thus the non-universality of the mass function is not currently a limitation to using the abundance of rich clusters to determine the normalization of the matter power spectrum. The uncertainties become increasingly important when it comes to using the evolution of the mass function as a probe of $\\Omega_{\\rm m}$ or the equation of state, $w$, of the dark energy.  For the latter, errors on $\\sigma_8$ approaching the percent level are required.  For these ambitious measurements it may not be sufficient to use a simple parameterized form for the multiplicity function.  One could either resort to full blown numerical simulations for a grid of models `near' the parameter region of interest or attempt to find a `second variable' which correlates well with the scatter between the simulation results and the P-S predictions. As we approach the level of precision where these effects matter a variety of other effects also become important, including the effects of clustering (see e.g.~\\S\\ref{sec:clustering}, Evrard et al.~\\cite{EvrHubble}, Hu \\& Kravtsov~\\cite{HuKra}) and how to treat merging systems. It will be a challenge for theorists to keep the `theory uncertainty' below the `experimental uncertainty' with the increasingly rapid advances in observations."
    },
    "0207/astro-ph0207650_arXiv.txt": {
        "abstract": "We present an \\xmm\\ observation of the bright, narrow-line, ultrasoft Seyfert 1 galaxy \\ton. The 0.3--10~keV X-ray spectrum is steep and curved, showing a steep slope above 2.5~keV ($\\Gamma \\sim 2.3$) and a smooth, featureless excess of emission at lower energies.  The spectrum can be adequately parameterised using a simple double power-law model.  The source is strongly variable over the course of the observation but shows only weak spectral variability, with the fractional variability amplitude remaining approximately constant over more than a decade in energy. The curved continuum shape and weak spectral variability are discussed in terms of various physical models for the soft X-ray excess emission, including reflection off the surface of an ionised accretion disc, inverse-Compton scattering of soft disc photons by thermal electrons, and Comptonisation by electrons with a hybrid thermal/non-thermal distribution. We emphasise the possibility that the strong soft excess may be produced by dissipation of accretion energy in the hot, upper atmosphere of the putative accretion disc. ",
        "introduction": "\\label{sect:intro}  Many Seyfert 1 galaxies and radio-quiet quasars possess X-ray spectra that steepen below $\\sim 2$~keV. This `soft excess' emission, usually defined as the excess emission over and above the hard X-ray power-law, was first seen in spectra obtained by \\heao\\ (Pravdo \\et\\ 1981), \\exosat\\ (Arnaud \\et\\ 1985; Turner \\& Pounds 1989) and \\einstein\\ (Bechtold \\et\\ 1987).  Its physical origin remains uncertain, as does its connection with the so-called `big blue bump' emission rising through the ultraviolet, although it is often associated with thermal emission from the putative accretion disc (e.g. Arnaud \\et\\ 1985; Czerny \\& Elvis 1987).  Boller, Brandt \\& Fink (1996) and Laor \\et\\ (1997) showed that the objects with the steepest soft X-ray spectra (implying the strongest soft excesses) tend to have relatively narrow optical \\hb\\ lines; many are classified optically as narrow-line Seyfert 1s (NLS1s; Osterbrock \\& Pogge 1985). These `ultrasoft' Seyferts often show other notable properties such as very rapid X-ray variability (e.g. Forster \\& Halpern 1996; Boller \\et\\ 1997; Turner \\et\\ 1999; Leighly 1999a; Brandt \\et\\ 1999) and strong optical \\feii\\ emission (e.g. Lawrence \\et\\ 1997; Vaughan \\et\\ 2001). The X-ray spectral form, X-ray variability and optical broad-line properties seem interconnected.  Here we report the results of a 30~ksec \\xmm\\ observation of Tonantzintla S180 ($z=0.062$) obtained as part of a guaranteed time programme (PI: Th. Boller) to  study the  timing and spectral properties of ultrasoft NLS1s using \\xmm.  \\ton\\ is one of the X-ray brightest ultrasoft Seyferts (Vaughan \\et\\ 1999; Leighly 1999a) and its luminosity is such that it is often classified as a quasar. Previous X-ray observations with \\asca\\ (Turner, George \\& Nandra 1998; Vaughan \\et\\ 1999; Leighly 1999b; Ballantyne, Iwasawa \\& Fabian 2001) and \\sax\\ (Comastri \\et\\ 1998) showed a strong excess of soft emission and tentative evidence for an  ionised iron line. The more recent high-resolution \\chandra\\ LETGS spectrum (Turner \\et\\ 2001b) showed no strong, narrow absorption or emission features and showed the soft X-ray continuum is smooth and featureless.  The rest of the paper is organised as follows. In the following section the observation and basic data reduction procedures are outlined. Section~\\ref{sect:spectra} describes the X-ray spectral fitting results. This is followed by an analysis of the variability properties of \\ton\\ in section~\\ref{sect:timing}.  The implications of these results are discussed in section~\\ref{sect:disco}. ",
        "conclusions": "\\label{sect:disco}  This paper presents an uninterrupted $\\sim 30$~ksec \\xmm\\ observation of the ultrasoft, narrow-line Seyfert 1 galaxy \\ton. The spectrum above 2.5~keV is slightly steeper than the norm for Seyfert 1 galaxies ($\\Gamma \\approx 2.26$ compared to $1.9$) but the iron K$\\alpha$ line emission is only poorly constrained. At lower energies the spectrum steepens, seen as a soft excess over the hard X-ray power-law with a featureless, approximately power-law form.  The shape of the soft excess is similar to that  observed in other objects. For example, Marshall \\et\\ (2002) found the high-resolution soft X-ray spectrum of the ultrasoft Seyfert 1 Mrk~478 (which has a similar X-ray luminosity to \\ton) was a featureless power-law. PKS 0558--504 also showed a similar soft excess shape (O'Brien \\et\\ 2001). The \\chandra\\ HETGS observation of the variable, narrow-line Seyfert 1 galaxy NGC~4051 again showed the soft excess to be a rapidly variable continuum component. However, in this lower-luminosity source it showed significant spectral curvature (Collinge \\et\\ 2001).  The observed flux of \\ton\\ showed rapid variability, changing by $\\sim 50$ per cent during the observation, but showed only weak spectral variability. This implies that, to first order at least, the curved X-ray spectrum is varying as one component. There are indications that this may be a general result,  that Seyferts with steep X-ray spectra show very little spectral variability (at least on short timescales). The \\asca\\ monitoring of Ark~564 and \\ton\\ showed flat RMS spectra (Turner \\et\\ 2001c; Romano \\et\\ 2001; Edelson \\et\\ 2002), as did the recent \\xmm\\ observation of 1H~0707--495 (Boller \\et\\ 2002).  This lack of short timescale spectral variability is somewhat different from the case in ``normal'' Seyfert 1s, which tend to show stronger variability at lower energies (e.g., Nandra \\et\\ 1997; Markowitz \\& Edelson, 2001; Vaughan \\& Edelson, 2001).  The lack of strong spectral features means that the observed form of the soft excess is consistent with a range of models. The data clearly rule out models based on blends of narrow soft X-ray lines (e.g., Turner \\et\\ 1991), but leave a variety of continuum emission mechanisms as plausible alternatives. Reprocessing of primary X-rays by an ionised accretion disc can produce a strong, smooth excess similar to that observed  but has difficulty simultaneously explaining the hard X-ray spectrum (section~\\ref{sect:refl_spec}). The observed lack of spectral variability is also a challenge for this model. The highly correlated, simultaneous variability between 3--10~keV flux (dominated by the primary X-rays) and softer X-ray bands (dominated by the reprocessed emission) requires the reprocessing to produce a near-perfect ``reverberation'' signature with a delay of $\\Delta t <500$~s (section~\\ref{sect:timing}).  Assuming a black hole mass of  $\\gs 10^{7} \\Msun$ (from the Eddington limit) this places the reprocessor within $\\ls 10 r_{\\rm g}$ of the hard X-ray source.  As first pointed out by Bechtold \\et\\ (1987) the soft excess is too hot to be ``bare'' thermal emission from the accretion disc.  The spectral form is consistent with emission from multiple blackbody components, but  the derived temperatures and sizes are inconsistent with those predicted for an accretion disc (unless \\ton\\ is highly super-Eddington).  Allowing for Doppler and gravitational shifts does not significantly alter this result.  Inverse-Compton scattering of soft photons by thermal electrons provides a more physically satisfying explanation for the broad soft excess (sections~\\ref{sect:compt_spec} and \\ref{sect:compt2_spec}), in which case the seed photon source is consistent with thermal accretion disc emission.   The harder power-law emission extending to 10~keV (and presumably beyond) can also be produced by Comptonisation, either in another purely thermal plasma or by non-thermal electrons in a plasma with a  hybrid thermal/non-thermal distribution.  Whether this is dominated by thermal or non-thermal electrons is impossible to tell without higher energy data; fits using thermal and hybrid thermal/non-thermal models are comparable in the \\xmm\\ band, where the predicted spectra both resemble power-laws.  A problem with the models discussed above is that they have difficulty explaining the rapid and energy-independent variability of \\ton\\ (section~\\ref{sect:timing}). These constraints force the hard and soft X-ray producing regions to be in very close causal contact with one another. If at high accretion rates the disc is puffed-up due to radiation pressure, then it is possible that the separation between the hot corona and the optically thick accretion disc can be negligibly small. In this scenario, some fraction of the accretion energy is dissipated in the surface layers of the disc, producing hot electrons which can Comptonise the soft photons from below.  The accretion disc model of Hubeny \\et\\ (2001) demonstrates the formation of a hot upper layer to the accretion disc due to dissipation near the surface (see their figure 10), and simulations of magnetized accretion discs by  Miller \\& Stone (2000) suggest that even higher temperatures may be reached, leading to the formation of a true corona.  Simulations of emergent disc spectra accounting for dissipation in the upper atmosphere will be presented elsewhere (Ross \\et\\ in prep.), but early indications suggest that the observed strong, smooth soft excess can be reproduced.  We also note the lack of a significant narrow, neutral iron emission line at 6.4~keV. The formal (90 per cent) upper limit on the equivalent with of such a line is $60$~eV, similar to that obtained for 1H 0707--495 by Boller \\et\\ (2002). Again this contrasts with the situation in more normal Seyfert 1s, which tend to show a narrow 6.4~keV emission line from distant, cold material (e.g., Yaqoob, George \\& Turner, 2002). However, whether this is due to a systematic difference between ultrasoft and normal Seyfert 1s remains to be seen."
    },
    "0207/astro-ph0207520_arXiv.txt": {
        "abstract": "We predict the fraction of dark halo lenses, that is, the fraction of lens systems produced by the gravitational potential of dark halos, on the basis of a simple parametric model of baryonic compression. The fraction of dark halo lenses primarily contains information on the effect of baryonic compression and the density profile of dark halos, and is expected to be insensitive to cosmological parameters and source population. The model we adopt comprises the galaxy formation probability $p_{\\rm g}(M)$ which describes the global efficiency of baryonic compression and the ratio of circular velocities of galaxies to virial velocities of dark halos $\\gamma_v=v_{\\rm c}/v_{\\rm vir}$ which means how the inner structure of dark halos is modified due to baryonic compression. The model parameters are constrained from the velocity function of galaxies and the distribution of image separations in gravitational lensing, although the degeneracy between model parameters still remains. We show that the fraction of dark halo lenses depends strongly on $\\gamma_v$ and the density profile of dark halos such as inner slope $\\alpha$. This means that the observation of the fraction of dark halos can break the degeneracy between model parameters if the density profile of dark halo lenses is fully settled. On the other hand, by restricting $\\gamma_v$ to physically plausible range we can predict the lower limit of the fraction of dark halo lenses on the basis of our model. Our result indicates that steeper inner cusps of dark halos  ($\\alpha\\gtrsim 1.5$) or too centrally concentrated dark halos are inconsistent with the lack of dark halo lenses in observations. ",
        "introduction": "Strong gravitational lensing offers a powerful probe of the matter distribution in the universe. So far $\\sim60$ lensed quasars are known, and their properties are summarized by CfA/Arizona Space Telescope Lens Survey (CASTLES\\footnote{Kochanek, C.~S., Falco, E.~E., Impey, C., Lehar, J., McLeod, B., \\& Rix, H.-W., http://cfa-www.harvard.edu/castles/}). Most of these lens systems have image separations $\\theta\\lesssim3''$. The cold dark matter (CDM) scenario, however, predicts sufficiently cuspy dark halos \\citep*[e.g.,][]{navarro96,navarro97}, thus dark halos can produce the significant amount of multiple images even at $\\theta\\gtrsim 3''$ \\citep*{wyithe01,keeton01b,takahashi01,li02,oguri02a}. It has been unclear whether the distribution of image separations should be computed based on galaxies \\citep*[using the luminosity function and the density profile of galaxies;][]{turner84,turner90,fukugita91,fukugita92,maoz93,kochanek96,chiba99} or dark halos \\citep*[using the mass function and the density profile of dark halos;][]{narayan88,cen94,wambsganss95,kochanek95,maoz97,wambsganss98,mortlock00,wyithe01,keeton01b,takahashi01,oguri02a}. The distribution of image separations, however, clearly indicates that the use of only dark matter properties cannot match observations. That is, the observed distribution of image separations is never reproduced from the theoretical calculation using the mass function of dark halos and only one population for lensing objects \\citep[e.g.,][]{li02}. The possible solution to explain  the distribution of image separations is the modification of inner structure of dark halos by introducing baryonic cooling \\citep*{keeton98,porciani00,kochanek01b,keeton01a,sarbu01,li02}. In this picture, inside low mass halos baryons are efficiently compressed and form sufficiently concentrated galaxies, while inside larger halos such as group- or cluster-mass halos global baryon cooling hardly occurs and thus the inner structure of dark halos remains unmodified.  Then a question comes to our mind: {\\it what is the fraction of dark halo lenses?} Here the term ``dark halo lenses'' is used to describe lenses which are produced by the gravitational  potential of dark halos. In other words, the lens objects of dark halo lenses are dark halos which have no central galaxies or have central galaxies but they are too small to dominate in gravitational lensing. Dark halo lenses  exhibit characteristic properties such as the small flux ratios and the detectable odd images \\citep{rusin02}, thus can be distinguished from usual galaxy lenses. Since dark halos are expected to have a steep central cusp, the significant amount of dark halo lenses should be observed. But so far no confirmed dark halo lens system is observed in strong gravitational lensing survey. The exception is arc statistics in rich clusters \\citep*{bartelmann98,williams99,meneghetti01,molikawa01,oguri01,oguri02b}, but the known cluster lenses were all found by searching for lenses in detail after identifying a rich cluster. In the surveys which first identify source objects and see whether they are lensed or not, it seems that dark halo lenses have not been observed yet: statistical argument \\citep*{kochanek99} and individual properties \\citep{rusin02} imply that current ambiguous quasar pairs are likely to be binary quasars. Even known lensing systems in clusters, such as Q0957+561 \\citep*{walsh79}, are produced mainly by a galaxy in the cluster. The cluster potential contributes to lensing only as a perturbation. Therefore, it should be checked whether the lack of dark halo lenses in observations really reconciles with the theoretical prediction.  Description of baryonic effects needs detailed models for the star formation and feedback \\citep[e.g.,][]{cole00}. Instead, in this paper we predict the expected fraction of dark halo lenses on the basis of a simple (minimal) parametric model \\citep{kochanek01a}. This model comprises the formation probability of galaxies, $p_{\\rm g}(M)$, and the ratio of circular velocities of galaxies to virial velocities of dark halos, $\\gamma_v$. The former describes the global efficiency of baryonic compression, and the latter models the modification of inner structure of dark halos due to baryonic compression. The model parameters are chosen so as to reproduce the velocity function of galaxies and the distribution of image separations. Although there remains the strong degeneracy between model parameters, this degeneracy can be broken from the observation of the fraction of dark halos if we fix the density profile of dark halos. On the other hand, by restricting a range of $\\gamma_v$ from various theories and observations, we can also derive the lower limit of the fraction of dark halo lenses. Our main finding is that steep inner slopes of dark halos ($\\alpha\\gtrsim1.5$) or too centrally concentrated dark halos are inconsistent with the lack of dark halo lenses in observations, even if various uncertainties are taken into account. Although this constraint on the density profile is not so severe, a large lens sample obtained by e.g., Sloan Digital Sky Survey (SDSS) can put tighter constraints on the density profile of dark halos as well as the model of baryonic compression.  The plan of this paper is as follows. In \\S \\ref{sec:th}, we describe the model of baryonic compression. Section \\ref{sec:const} is devoted to constrain the model parameters, and \\S \\ref{sec:frac} presents our predictions for the fraction of dark halo lenses. Finally, we summarize conclusions in \\S \\ref{sec:conc}. Throughout this paper, we assume the lambda-dominated cosmology $(\\Omega_0, \\lambda_0,h,\\sigma_8)=(0.3,0.7,0.7,1.04)$, where the Hubble constant in units of $100{\\rm km\\,s^{-1}Mpc^{-1}}$ is denoted by $h$. As shown below, however, our results are quite insensitive to a particular choice of cosmological parameters. ",
        "conclusions": "} We have studied the effect of baryonic compression assuming the simple parametric model used by \\citet{kochanek01a}. Our model has following two elements: the galaxy formation probability $p_{\\rm g}(M)$ (eq. [\\ref{gfp}]) which describes the global efficiency of baryonic compression, and the ratio of circular velocities of galaxies to virial velocities of dark halos $\\gamma_v=v_{\\rm c}/v_{\\rm vir}$ which means how the inner structure of dark halos is modified due to baryonic compression.  The model parameters are constrained from the observed velocity function of galaxies and the distribution of image separations in strong gravitational lensing, although the strong degeneracy between model parameters still  remains. By using this model, we predict the fraction of dark halo lenses $f_{\\rm dark}$ (eq. [\\ref{fdark}]). Here  dark halo lenses mean the lens systems which are produce by the gravitational potential of dark halos. The fraction of dark halo lenses is independent of the normalization of total lensing rate, thus is insensitive to cosmological parameters and information of sources such as redshift and flux distribution. Instead the fraction of dark halo lenses is expected to have information on both the effect of baryonic compression and the density profile of dark halos such as the inner density profile for which we modeled $\\rho\\propto r^{-\\alpha}$. We found that $f_{\\rm dark}$ is indeed sensitive to the inner slope $\\alpha$, concentration parameter $c_{\\rm norm}$, and model parameters such as $\\gamma_v$. Therefore, definite predictions of $f_{\\rm dark}$ need correct knowledge of baryonic compression as well as the density profile of dark halos. This also means that we can constrain the model of baryonic compression from the observation of $f_{\\rm dark}$ if the density profile of dark halos is well known.  Although the fraction of dark halo lenses is difficult to predict, we can still derive the lower limit of $f_{\\rm dark}$ by restricting $\\gamma_v\\leq2$ which is inferred from various theories and observations (see \\S \\ref{sec:const:vc}). We found that the steep inner profiles ($\\alpha\\gtrsim1.5$) or too centrally concentrated dark halos are inconsistent with the lack of dark halo lenses in observations. As described above, our result is quite insensitive to cosmological parameters and source population. Therefore, our result is complementary to the result of \\citet{keeton01b} who obtained similar constraint on the density profile of dark halos using the total lensing probability which also depends strongly on cosmological parameters. One of possible systematic effects which may change $f_{\\rm dark}$ is the effect of galaxies in groups or clusters. By using the result of \\citet{keeton00}, we found that this effect changes $f_{\\rm dark}$ by a factor $3/4$ at most. Therefore this effect does not change our main result so much. Other important systematic effect is the existence of empty halos which are neglected in our model. This effect is, however, not important for the lower limit of $f_{\\rm dark}$, because this effect only increases the fraction of dark halo lenses.  One possible criticism of our result is that the baryon compression model we use in this paper is too simple. Many of the previous work, however, have not addressed this problem. For example, most work of gravitational lensing statistics which use the mass function of dark halos assumed that circular velocities of galaxies are the same as virial velocities of dark halos (this corresponds to $\\gamma_v=1$ in our model). But this assumption seems to be invalid \\citep[e.g.,][]{seljak02a}, and since the number of galaxy lenses scales as $\\propto \\gamma_v^4$ the deviation from $\\gamma_v=1$ should not be dismissed. This fact is seen even in our simple model where the connection between galaxy lenses and dark halo lenses sensitively depends on $\\gamma_v$, and is indeed difficult to be determined. The importance of baryonic compression, however, means that we can constrain the model of baryonic compression from observations of the fraction of dark halo lenses. Although the current sample of gravitational lensing may be too small for this purpose, the larger lens sample obtained by e.g., SDSS can strongly constrain the model of baryonic compression as well as the density profile of dark halos."
    },
    "0207/astro-ph0207228_arXiv.txt": {
        "abstract": "Occultation and microlensing are different limits of the same phenomena of one body passing in front of another body.  We derive a general exact analytic expression which describes both microlensing and occultation in the case of spherical bodies with a source of uniform brightness and a non-relativistic foreground body.  We also compute numerically the case of a source with quadratic limb-darkening. In the limit that the gravitational deflection angle is comparable to the angular size of the foreground body, both microlensing and occultation occur as the objects align.  Such events may be used to constrain the size ratio of the lens and source stars, the limb-darkening coefficients of the source star, and the surface gravity of the lens star (if the lens and source distances are known).  Application of these results to microlensing during transits in binaries and giant-star microlensing are discussed.  These results unify the microlensing and occultation limits and should be useful for rapid model fitting of microlensing, eclipse, and ``microccultation'' events. ",
        "introduction": "When two stars (or other bodies) come into close alignment on the sky, the foreground star may either eclipse or microlens the background star.  As the stars align, if the angular size of the foreground star is much larger than its gravitational deflection angle, then the foreground star can eclipse; if the contrary is true then it can magnify.  More precisely, gravitational lensing by a point mass produces two images of a distant object, one interior and one exterior to the Einstein radius in the lens plane, $R_E=[4R_G D_L(D_S-D_L)/D_S]^{1/2}$ where $R_G=GM/c^2$ is the gravitational radius for a lens of mass $M$, and $D_{L,S}$ are the distances to the lens or source.  Both images move toward the Einstein radius as the lens and source approach, so the outer image will be occulted during the approach if the radius of the lens is larger than the Einstein radius.  The inner image, however, starts off near the origin and thus is occulted when the source is far from the lens.  As the lens and source approach, the inner image can become unocculted if the lens is smaller than the Einstein radius (Figure 1). Occultation is most important in microlensing if $R_E \\sim R_L$, where $R_L$ is the radius of the lens star (assumed to be spherical). In Galactic microlensing, typically $R_L \\ll R_E$, so the occultation of the inner image occurs, but is usually rather faint. In special circumstances, such as in eclipsing binaries containing compact objects \\citep{mae73,mar01} or lensing by giant stars, $R_L \\sim R_E$, so the effects of both microlensing and occultation must be included. This ``microccultation'' can show more varied behavior than the usual microlensing or occultation lightcurves and can be used to constrain the surface gravity of the lens star \\citep{bro96}.  \\begin{figure*} \\centerline{\\psfig{file=f1.eps,width=5in}} % \\figcaption[]{(a) Stars at various positions in source plane.  Shaded region shows the area in which the inner image is occulted by the lens star for $r_L<1$, which is the region $\\beta > \\beta_L = 1/r_L -r_L$.  The axes have units of $R_E D_S/D_L$.  (b) Images of the star in the image plane. The arrow is imaged as well for reference.  The dashed line is $R_E$.  (c) Magnification as a function of position - dotted line is without occultation, while solid line includes occultation.  The symbols show the magnification of the source at the positions depicted in (a). (d)-(f) Same as (a)-(c), but for $r_L>1$.  In this case, the inner image is fully occulted, and the shaded region in (d) shows where the outer image is occulted by the lens star.} \\end{figure*}  \\citet{mae73} and \\citet{mar01} have carried out numerical computations of microccultation lightcurves.  Here we present an exact {\\it analytic} solution for the lightcurve of a uniform source which agrees with their work, and we present numerical calculations for limb-darkened sources. \\citet{bro96} and \\citet{boz02} computed lightcurves for lensing events, treating the source as a point source, while the expressions presented here are valid for extended and limb-darkened sources as well. In section 2 we discuss microlensing and occultation of a point source. In section 3 we include the finite size of a uniform source.  In section 4 we include the effects of limb-darkening numerically for microlensing and occultation. In section 5 we apply the results to several astrophysical cases of possible interest, namely white dwarf-main sequence binaries, microlensing in globular clusters, and microlensing by supergiants. In section 6 we summarize. ",
        "conclusions": "We have computed exact formulae for lensing of a uniform extended source by an opaque, spherical lens (with escape velocity much smaller than $c$). The formulae only differ significantly from the usual occultation or microlensing formulae in the limit that $R_L \\sim R_E$, which may be relevant for lensing by white dwarfs in binaries or lensing by giant stars.  Small deviations due to lensing in eclipsing brown-dwarf binaries may be detectable with very precise photometry, which may be another application of the expressions derived here.  A code written in IDL which carries out the calculations presented here can be downloaded from \\url{http://www.astro.washington.edu/agol/}."
    },
    "0207/astro-ph0207534_arXiv.txt": {
        "abstract": "Although the currently favored cold dark matter plus cosmological constant model for structure formation   assumes  an  $n=1$ scale-invariant initial power  spectrum,  most  inflation models  produce at least  mild deviations  from $n=1$.  Because  the lever arm from the CMB normalization to galaxy scales is long, even a small ``tilt'' can have important   implications  for galactic observations.  Here  we calculate    the   COBE-normalized   power    spectra   for   several well-motivated models of  inflation and compute implications for the substructure content and central  densities of  galaxy halos. Using an analytic model, normalized against N-body simulations, we show that  while halos in  the standard  ($n=1$)  model are overdense by a factor of $\\sim 6$ compared to observations, several of our  example inflation+LCDM models  predict halo densities well within the range of  observations, which prefer models  with $n \\sim 0.85$. We go on to use a semi-analytic model  (also normalized against N-body simulations) to follow the merger histories  of galaxy-sized halos and track the orbital decay, disruption, and evolution of the merging substructure. Models with $n \\sim 0.85$  predict  a factor of  $\\sim 3$ fewer  subhalos at a fixed circular velocity than the standard $n = 1$ case.  Although this level  of reduction does   not resolve the ``dwarf  satellite problem'',  it  does  imply  that   the level  of  feedback  required to match the observed number of dwarfs is sensitive to   the initial power  spectrum.  Finally,  the fraction of galaxy-halo mass that is bound up in substructure is consistent with limits imposed by multiply  imaged quasars for all models considered: $f_{sat} > 0.01$ even for an effective tilt of $n \\sim  0.8$. We conclude that, at their current level, lensing constraints of this kind do not provide strong limits on the primordial power spectrum. \\vspace{1pc} ",
        "introduction": "The  cold  dark matter  plus  cosmological  constant  (LCDM) model  of structure formation is the  most successful and popular current theory for  the  origin  of   universal  structure.   The  theory  is  highly predictive, and given that observations constrain the universe to be nearly flat (with $\\Omega_m = 1 - \\Omega_{\\rm \\Lambda}  \\approx 0.3$, and $h \\approx  0.7$) a  crucial unknown  is the  primordial spectrum of density  fluctuations.   It  is   usually  assumed  that  the  initial fluctuation spectrum is scale-invariant, with $P(k) \\propto k^n$, $n = 1$.   This choice is motivated by the fact that slow-roll inflation models predict {\\em  nearly}  scale-invariant spectra.   However, generic models  of inflation do  not predict primordial  power spectra that  are {\\em exactly}  scale-invariant and  almost all  models have some small, often non-negligible,  ``tilt'' ($ n\\ne1$) and ``running'' of the spectral index ($ dn/d\\ln k \\ne 0$).  In this proceeding we report on   our  work  to   model  the small-scale implications of reasonable  initial  power   spectra,  and  demonstrate  that  the substructure  content and central  densities of  dark halos  should be very sensitive to small deviations from scale-invariance.  Our models are semi-analytic, but they are normalized and tested against N-body simulations.  As we discuss below, our results may be important for interpreting galaxy rotation curves, dwarf galaxy counts, and substructure mass fractions as measured from multiply imaged quasars.  In Zentner \\& Bullock \\cite{zb} (hereafter ZB1 and ZB2) we chose several well-motivated and representative inflation models and calculated the implied power spectra to second order in slow-roll using the method of Stewart and Gong \\cite{SG}. Adopting $\\Omega_{\\rm  M} = 0.3$, $\\Omega_{\\rm L} = 0.7$, $\\Omega_{\\rm b}h^2 = 0.020$, and $h = 0.72$ (LCDM), we calculated P(k) using the results of Eisenstein \\& Hu \\cite{EH} and normalized to COBE on the pivot scale, $k_* \\simeq 0.0023$ h$^{-1}$Mpc. In the interest of brevity, we report on only three models here: \\begin{itemize} \\item{IPL: $n = 0.94$, $dn/d\\ln k = -0.001$, \\\\ $\\sigma_8 = 0.83$} \\item{RM: $n = 0.84$, $dn/d\\ln k = -0.004$, \\\\ $\\sigma_8 = 0.65$} \\item{$n=1$: $n = 1.00$, $dn/d\\ln k = 0.000$, \\\\ $\\sigma_8 = 0.95$}. \\end{itemize} The tilt and running are evaluated at $k=k_*$ and $\\sigma_8$ is the rms fluctuation amplitude within spheres of $8  h^{-1}$Mpc implied by the COBE normalization. Observational constraints on $\\sigma_8$ have historically favored a value near unity; however, some recent  estimates  indicate surprisingly  low  values  of  $\\sigma_8$  \\cite{sig8}. Roughly speaking,  $0.60 < \\sigma_8 < 1.2$ spans the range of recently reported values (for $\\Omega_m \\simeq 0.3$). ",
        "conclusions": ""
    },
    "0207/astro-ph0207644_arXiv.txt": {
        "abstract": "We introduce exact analytical solutions of the steady-state hydrodynamic equations of scale-free, self-gravitating gaseous disks with flat rotation curves. We express the velocity field in terms of a stream function and obtain a third-order ordinary differential equation (ODE) for the angular part of the stream function. We present the closed-form solutions of the obtained ODE and construct hydrodynamical counterparts of the power-law and elliptic disks, for which self-consistent stellar dynamical models are known. We show that the kinematics of the Large Magellanic Cloud can well be explained by our findings for scale-free elliptic disks. ",
        "introduction": "Astrophysical disks are either stellar or gaseous. The existence of non-axisymmetric equilibrium states of such systems has been attractive both from theoretical and observational aspects. The search for dynamical equilibria of stellar systems is known as the self-consistency problem where one seeks for a positive phase space distribution function (DF) that produces the density of the model and satisfies the collisionless Boltzmann equation (Binney \\& Tremaine 1987). The construction of the DF can be done by analytical means or numerical exploration based on Schwarzschild's orbit superposition method (Schwarzschild 1979; Kuijken 1993, hereafter K93; Levine \\& Sparke 1998; Teuben 1987; de Zeeuw, Hunter, \\& Schwarzschild 1987; Jalali \\& de Zeeuw 2002, hereafter JZ). The main difficulty linked to the self-consistency problem is the lack of the second integral of motion for most planar models.  Gaseous disks, however, are investigated through solving the hydrodynamic equations of compressible, self-gravitating fluids. There are several works in the literature that deal with the stability of perturbed axisymmetric gaseous disks (Heemskerk, Papaloizou, \\& Savonije 1992; Goodman \\& Evans 1999; Lemos, Kalnajs, \\& Lynden-Bell 1991). The problem of non-axisymmetric systems is somewhat different because our knowledge is limited even about equilibrium states. Following a first-order linear theory, Syer \\& Tremaine (1996, hereafter ST96) constructed simple scale-free disks and approximately solved the steady-state hydrodynamic equations for the components of the velocity vector. They also proposed a nonlinear theory based on the calculation of Fourier coefficients using Newton's method. A similar strategy was adopted by Galli et al. (2001, hereafter G01) in the study of isothermal molecular clouds and the fragmentation hypothesis of multiple stars. Their nonlinear solutions, in contrast with ST96, were generated by direct numerical integration of reduced ordinary differential equations (ODE) and an iterative scheme to match the boundary conditions. These methods of approach seem inefficient when one seeks for global properties of highly non-axisymmetric equilibria and the transient response of fluid disks in the vicinity of such states. Therefore, finding analytical solutions for the hydrodynamic equations is of special interest even in the steady-state conditions.  In this paper we explore exact non-axisymmetric solutions of self-similar gaseous disks with flat rotation curves. In \\S2 we express the governing equations and reduce them to a third-order ODE in \\S3 where we also introduce our {\\it closed-form} analytic solutions. We construct hydrodynamical counterparts of the JZ models in \\S4 and apply our results to the Large Magellanic Cloud (LMC) in \\S5. We show that the kinematics of the LMC can be explained by our model elliptic disks. We end up with conclusions in \\S6. ",
        "conclusions": "Exact equilibrium solutions of self-gravitating systems are rare. Two classes of such solutions were found by Medvedev \\& Narayan (2000) for three-dimensional, axisymmetric isothermal systems. A generalization of these solutions for non-axisymmetric rotating systems was then introduced by Shadmehri \\& Ghanbari (2001). Although the density distribution of their models is non-axisymmetric, the velocity field has only the toroidal component. This makes their models too simple to be used in the study of fully three-dimensional non-axisymmetric distributions of matter.  In a systematic search for non-axisymmetric equilibrium solutions, we started our study from gaseous disks because they inherit many properties of three-dimensional systems. Using the benefit of the existence of a stream function, we were successful in reducing the governing equations of scale-free models to an incomplete linear ODE for the angular part of the stream function. It then became clear that Liouville's transformation can reduce the obtained ODE to a second-order inhomogeneous ODE, which is similar to that of a forced harmonic oscillator. This was striking because it led us to derive closed-form solutions for the radial velocity component, stream function and the pressure $p$. In fluid mechanics, transport equations are written for mass, momentum and energy. While the pressure occurs in governing equations, there is no extra independent equation in order to determine $p$. The problem is rather difficult in the presence of an anisotropic stress tensor, which is encountered in turbulence and Jeans' equations of stellar hydrodynamics. Taking the curl of momentum equations is a trick that temporarily removes the pressure from our calculations of the stream function. But, one should take care of the sign of $p$ once it is eventually extracted from momentum equations. Only a positive pressure distribution is physical because a fluid medium can resist against a compressive stress not a tensional one.  Non-axisymmetric stellar disks have necessarily anisotropic DFs and pressure tensors. This point is the main difference between the physics of a steady-state gaseous disk and its stellar counterpart described by Jeans' equations. Nevertheless, isotropic pressure distribution is not an obstacle for the existence of non-axisymmetric, self-gravitating gaseous disks as we showed in \\S\\ref{sec:examples}.  Our scale-free elliptic disks were used to model the outer part of the LMC disk. The main restriction of our modeling is the isotropy of the stress tensor. We speculate that our results will not be changed in a significant way if the model is stellar with an anisotropic pressure. The other important issue is the perturbation of the Galaxy. The tidal field of the Galaxy induces a force whose $R$-distribution is totally different from that of a logarithmic potential field. However, if the tidal field of the Galaxy was dominant, the rotation curve of the LMC would show substantial deviation from flatness. But, we know that the rotation curve of the LMC is almost flat. Therefore, our isolated models provide acceptable {\\it zeroth-order} solutions. Such a modeling of the LMC is a quick route for interpreting the observational data. It is obvious that taking the tidal effects into account, will destroy the enjoyable scale-free nature of our disks."
    },
    "0207/astro-ph0207191_arXiv.txt": {
        "abstract": "{We present the redshift distribution of a complete sample of 480 galaxies with $K_s<20$ distributed over two independent fields covering a total area of 52 arcmin$^2$. The redshift completeness is 87\\% and 98\\% respectively with spectroscopic and high-quality and tested photometric redshifts. The redshift distribution of field galaxies has a median redshift $z_{med}\\sim 0.80$, with $\\sim$ 32\\% and $\\sim$9\\% of galaxies at $z>1$ and $z>1.5$ respectively. A ``blind'' comparison is made with the predictions of a set of the most recent $\\Lambda$CDM hierarchical merging and pure luminosity evolution (PLE) models. The hierarchical merging models overpredict and underpredict the number of galaxies at low-$z$ and high-$z$ respectively, whereas the PLE models match the median redshift and the low-$z$ distribution, still being able to follow the high-$z$ tail of $N(z)$. We briefly discuss the implications of this comparison and the possible origins of the observed discrepancies. We make the redshift distribution publicly available. ",
        "introduction": "The mass assembly history of galaxies remains one of the critical issues in observational cosmology: did galaxies reach their present stellar mass only recently (say, at $z\\lsim 1$) ? Or were most (massive) galaxies already in place by $z\\sim 1$ ? Spectroscopic surveys of faint galaxies selected in the $K$-band currently offer the best opportunity to answer these questions (Broadhurst et al. 1992). The main advantages with respect to optically selected samples include: the direct sensitivity to the galaxy stellar mass rather than to the ongoing/recent star formation activity (Gavazzi et al. 1996; Madau et al. 1998), the smaller K-correction effects, and the minor influence of dust extinction.  In this framework, we have completed an optical and near-infrared spectroscopic survey down to $K_s<20$ (dubbed ``K20 survey'') using ESO VLT telescopes and instruments, with full survey details being given in Cimatti et al. (2002b, hereafter Paper III; see also {\\tt http://www.arcetri.astro.it/$\\sim$k20/}).  The K20 sample includes 546 objects to $K_s \\leq 20$ (Vega system) over two independent fields (52 arcmin$^{2}$ in total), so to be less affected by the cosmic variance. The spectroscopic redshift completeness is 94\\% and 87\\% for $K_s \\leq 19$ and $K_s \\leq 20$ respectively. This makes the K20 sample the largest and most complete spectroscopic sample of galaxies with $K_s<20$ available to date (see Paper III; cf. Cowie et al. 1996; Cohen et al. 1999). Moreover, a 98\\% redshift completeness is reached for the $K_s \\leq 20$ sample when including the photometric redshifts obtained with the available deep $UBVRIzJK_s$ imaging for those objects without a spectroscopic redshift. If stars and broad-line AGNs are excluded, the total number of galaxies with $K_s\\leq20.0$ and with redshifts is 480.  In two previous papers based on the K20 survey we showed that Extremely Red Objects (EROs, defined by $R-K_s>5$) are nearly equally populated by old passively evolving galaxies and by dusty star-forming systems at $z \\sim 1$ (Cimatti et al. 2002a, Paper I). The number of all (old+dusty) EROs is strongly underpredicted by hierarchical merging models (HMMs), whereas old EROs have a density consistent with PLE models for passive early-type galaxies (Paper I), and are strongly clustered as opposed to dusty EROs (Daddi et al. 2002, Paper II).  In this Letter, we present and discuss the observed redshift distribution, $N(z)$, of all the galaxies in the K20 sample, irrespective of their color, and compare it to the expectations for the case of PLE of galaxies, as well as to the predictions of various HMM renditions. The currently favored cosmological model is adopted, i.e., $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "Early predictions of the expected fraction of galaxies at $z>1$ in a $K<20$ sample indicated respectively $\\approx 60\\%$ and $\\approx 10\\%$ for a PLE case and for a (then) standard $\\Omega_m=1$ CDM model (Kauffmann \\& Charlot 1998). This version of PLE was then ruled out by Fontana et al. (1999). The more recent PLE and HMM used in this paper consistently show that for $z>1$ the difference between the predictions of different scenarios is much less extreme. These results partly from the now favored $\\Lambda$CDM cosmology which pushes most of the merging activity in hierarchical models at earlier times compared to $\\tau$CDM and SCDM models with $\\Omega_m=1$, and partly to different recipes for the star formation modes, which tend to narrow the gap between HMMs and the PLE case (e.g. Somerville et al. 2001; Firth et al. 2002). The disagreement between the observed $N(z)$ and the predictions of the most updated HMMs based on a $\\Lambda$CDM cosmology would then become even stronger in the case of old-fashioned CDM models with $\\Omega_m=1$ because structures form later in a matter-dominated universe, and thus they would predict an even lower fraction of galaxies at high-$z$. In this respect, our results can be seen as additional evidence that the universe is not matter-dominated ($\\Omega_m<1$), and suggest that the HMMs may perform better if $\\Omega_m$ is even lower than the currently favored $\\Omega_m=0.3$.  Nevertheless, the results of the K20 survey indicate that the shape and the median of the observed redshift distribution of $K_s<20$ galaxies are in broad agreement with the expectations of PLE models, while disagree with the predictions of current hierarchical merging models of galaxy formation. This discrepancy refers to all galaxies, irrespective of color or morphology selection, and therefore is more general than the already noted discrepancies with EROs (Daddi et al. 2000; Paper I; Cimatti 2002). The poor performance of HMMs in accounting for the properties of even $z=0\\rightarrow\\sim 1$ early-type galaxies has been emphasized in the past (e.g., Renzini 1999; Renzini \\& Cimatti 1999). Moreover, among low-redshift galaxies there appears to be a clear anti-correlation of the specific star formation rate with galactic mass (Gavazzi et al. 1996; Boselli et al. 2001), the most massive galaxies being ``old'', the low-mass galaxies being instead dominated by young stellar populations. This is just the opposite than expected in the traditional hierarchical merging scenario, where the most massive galaxies are the last to form.  On the other hand, the strong clustering of EROs seems to be rather consistent with the predictions of CDM models of large scale structure evolution (Daddi et al. 2001, Paper II; Firth et al. 2002). Thus, adopting the hierarchical merging $\\Lambda$CDM scenario as the basic framework for structure and galaxy formation, the observed discrepancies may be ascribed to the heuristic algorithms adopted for the star formation processes and their feedback, both within individual galaxies and in their environment. Our results suggest that HMMs should have galaxy formation in a CDM dominated universe to closely mimic the old-fashioned {\\it monolithic collapse} scenario. This requires to enhance merging and star formation in massive haloes at high redshift (say, $z\\gsim 3$), while in the meantime suppressing star formation in low-mass haloes. For instance, Granato et al. (2001) suggested the strong UV radiation feedback from the AGN activity during the era of supermassive black hole formation to be responsible for the suppression of star formation in low-mass haloes, hence imprinting a ``anti-hierarchical'' behavior in the baryonic component. The same effect may well result from the feedback by the starburst activity itself (see also Ferguson \\& Babul 1998).  In summary, the redshift distribution presented in this paper, together with the space density, nature, and clustering properties of the ERO population (Paper I, Paper II) and the redshift evolution of the luminosity and stellar mass functions derived for the K20 sample (Pozzetti et al. 2002, Fontana et al. 2002) provide a new set of observables on the galaxy population in the $z\\sim 1-2$ universe, thus bridging the properties of $z\\sim 0$ galaxies with those of Lyman-break and submm/mm-selected galaxies at $z$ \\gtsima 2-3. While making a step towards the fully empirical mapping of galaxy formation and evolution, this set of observables poses a new challenge for theoretical models to properly reproduce."
    },
    "0207/astro-ph0207472_arXiv.txt": {
        "abstract": "{ The basic form of drifting sub-pulses is that of a periodicity whose phase depends (approximately linearly) on both pulse longitude and pulse number. As such, we argue that the two-dimensional Fourier transform of the longitude-time data (called the Two-Dimensional Fluctuation Spectrum; 2DFS) presents an ideal basis for studies of this phenomenon.  We examine the 2DFS of a pulsar signal synthesized using the parameters of an empirical model for sub-pulse behaviour. We show that the transform concentrates the modulation power to a relatively small area of phase space in the region corresponding to the characteristic frequency of sub-pulses in longitude and pulse number. This property enables the detection of the presence and parameters of drifting sub-pulses with great sensitivity even in data where the noise level far exceeds the instantaneous flux density of individual pulses.  The amplitude of drifting sub-pulses is modulated in time by scintillation and pulse nulling and in longitude by the rotating viewing geometry (with an envelope similar to that of the mean pulse profile). In addition, sub-pulse phase as a function of longitude and pulse number can differ from that of a sinusoid due to variations in the drift rate (often associated with nulling) and through the varying rate of traverse of magnetic azimuth afforded by the sight line.  These deviations from uniform sub-pulse drift manifest in the 2DFS as broadening of the otherwise delta-function response of a uniform sinusoid. We show how these phase and amplitude variations can be extracted from the complex spectrum. ",
        "introduction": "Very soon after the discovery of pulsars, it was noticed that the intensity data sometimes showed ``second periodicities'' within each pulse, and that the relative phase of this periodic signal drifted in an organised manner from pulse to pulse \\citep{dc68}. The striking patterns made by drifting sub-pulses in two-dimensional pulse longitude-time diagrams seemed to be saying something about the fundamentals of radio pulsar emission, but the question was, and still is, `what?' Numerous studies of the phenomenon were conducted in the first few years of pulsar science (e.g. \\citealt{tjh69,sspw70,col70a,bac70b,bac70c,th71,bac73,pag73}), their frequency steadily decreasing as the limitations of existing online (recording) and offline (analysis) equipment were reached. Along with attempts at fitting the data with sub-pulses and tracking their drift, the Fourier technique developed by Backer in the references above (now known as the longitude-resolved fluctuation spectrum, or LRFS) came to be perhaps the most widely used means of examining the average properties of sub-pulse drift.  With the availability of modern digital back-ends and convenient offline computing power, we believe that there is now potential for renewed progress in the field of drifting sub-pulses. We describe in this paper a means of analysis of drifting sub-pulse data using what may be considered an extension of the Fourier work of Backer. The technique was first conceived as a means of detecting and characterising drifting sub-pulses in the large population of moderately weak pulsars that have been discovered since the 1970s when most studies took place. It is also useful, when sufficient signal-to-noise ratio is available, for the investigation of what might be considered the fine details of the phenomenon: deviations from constant amplitude sub-pulses with purely uniform phase evolution. {\\changed We begin with a mathematical description of a drifting sub-pulse signal as a function of pulse longitude and pulse number (Sect.~\\ref{sec:signal}) and examine its two-dimensional Fourier response (Sect.~\\ref{sec:2dfs}). We then describe a technique for measuring the deviations of a given signal from a uniform pure sinusoid (Sect.~\\ref{sec:extract}), examine what is expected under the usual model of drifting sub-pulses (Sect.~\\ref{sec:carousel}), and use this model as a basis for simulations testing the applicability of the methods of analysis described here (Sect.~\\ref{sec:sim}).} ",
        "conclusions": "We have shown that the two-dimensional Fourier spectrum of pulsar longitude-time data is of value in the analysis of drifting sub-pulses. We consider the drifting component of the signal as the convolution of a complex ``modulation envelope'' with a pure two-dimensional periodicity. The modulation envelope describes the variations in average sub-pulse amplitude and phase as a function of longitude and time, and can be decomposed into the product of two one-dimensional envelopes which are functions of time and longitude respectively. This makes the technique well-suited to studying a variety of phenomena associated with drifting sub-pulses, including longitudinal variations in sub-pulse spacing (induced due to viewing geometry or other factors), variations in the drift rate as a function of time (due to the recovery from a null, random phase noise, etc.), comparison of the longitudinal amplitude dependence of the drifting and steady components of the pulsar emission, and variations in the average sub-pulse amplitude as a function of time (due to nulling, carousel rotation, etc.)."
    },
    "0207/astro-ph0207158_arXiv.txt": {
        "abstract": "Observations of the starburst galaxy, M82, have been made with the VLA in its A-configuration at 15~GHz and MERLIN at 5~GHz enabling a spectral analysis of the compact radio structure on a scale of $< 0.1''$ (1.6~pc). Crucial to these observations was the inclusion of the Pie Town VLBA antenna, which increased the resolution of the VLA observations by a factor of $\\sim$2. A number of the weaker sources are shown to have thermal spectra and are identified as H{\\sc ii} regions with emission measures $\\sim$10$^7$~cm$^{-6}$~pc. Some of the sources appear to be optically thick at 5~GHz implying even higher emission measures of $\\sim$10$^8$~cm$^{-6}$~pc. The number of compact radio sources in M82 whose origin has been determined is now 46, of which 30 are supernova related and the remaining 16 are H{\\sc ii} regions. An additional 15 sources are noted, but have yet to be identified, meaning that the total number of compact sources in M82 is at least 61. Also, it is shown that the distribution of H{\\sc ii} regions is correlated with the large-scale ionised gas distribution, but is different from the distribution of supernova remnants. In addition, the brightest H{\\sc ii} region at (B1950) 09$^h$ 51$^m$ 42.21$^s$ +69$^{\\circ}$ 54$'$ 59.2$''$ shows a spectral index gradient across its resolved structure which we attribute to the source becoming optically thick towards its centre. ",
        "introduction": "The radio emission from star-forming galaxies can be either of a thermal (e.g. free-free) or non-thermal (e.g. synchrotron radiation) origin. The non-thermal radiation originates from relativistic electrons which have been accelerated by supernovae resulting from the deaths of massive stars and the thermal emission is from free-free emission from gas which has been ionised by hot, young stars. Therefore, it is obviously of interest to investigate the relative contributions of these processes to the emission which we observe from nearby star-forming galaxies. A separation of the non-thermal and thermal components of radio emission from the archetypal starburst galaxy, M82, has already been carried out by \\citet{Allen99} on angular scales of 2$''$ and they identified several complexes of ionised gas. However, 2$''$ is not sufficient resolution to separate the most compact radio sources from the more diffuse radio structure. In this paper we compare observations of M82 made with the VLA in its A-configuration at 15 GHz, incorporating the VLBA Pie Town antenna, with MERLIN 5 GHz observations. The maps made from these observations have a resolution $\\leq$~0.1$''$ which extends the analysis of the most compact radio structure into the thermal regime. Thus, we are now able to comment on the relative contribution of thermal and non-thermal processes to the radio emission on scales of order 1~pc. For consistency with previous publications, we assume a distance of 3.2~Mpc to M82 throughout the paper (Burbidge, Burbidge \\& Rubin 1964)\\nocite{Burbidge64}.  Spectral information on 26 of the compact radio sources in M82 has already been published \\citep{Wills97,Allen98}. These analyses were prompted by 408~MHz MERLIN observations and the radio continuum spectra of the detected sources were mainly consistent with non-thermal synchrotron emission, occasionally with a low-frequency turnover which was attributed to free-free absorption by ionised gas. However, free-free absorption is likely to be responsible for the non-detection of many of the remaining sources at low frequencies. In addition, \\citet{Wills97} (hereafter W97) identified two sources which were inconsistent with a free-free absorbed synchrotron spectrum, the AGN candidate, 44.01+59.6\\footnote{All positions within M82 are quoted relative to (B1950) 09$^h$ 51$^m$ +69$^{\\circ}$ 54$'$} and a possible H{\\sc ii} region at 40.62+56.0. \\citet{Allen98} (hereafter AK98) came to similar conclusions regarding the majority of the sources, but additionally identified a flat-spectrum ($\\alpha \\sim$ -0.1, where $S \\propto \\nu^{\\alpha}$) source at 42.21+59.2, which they identified as an H{\\sc ii} region. However, over 50 compact radio sources have been identified in the nuclear region of M82, therefore until now there has only been spectral information published on around half of the total number, similar to the situation in another nearby starburst, NGC~2146 \\citep{Tarchi00}. Also, since previous investigations concentrated on the brightest sources at the longer radio wavelengths, there would have been a bias away from the identification of thermal sources such as H{\\sc ii} regions. This is because their electron temperatures are unlikely to exceed 10$^4$~K and hence even if the ionised gas is optically thick, the brightness temperatures cannot exceed this limit (n.b. 10$^4$~K $\\equiv$ 0.48~mJy~beam$^{-1}$ with MERLIN at 5~GHz, 50~mas resolution). Therefore, in order to make robust comparisons between the numbers of compact radio sources of different types in starburst galaxies we must first make more confident identifications of a greater fraction of the sources. Furthermore, with the improved sensitivity of radio observations, it is clear that further populations of compact radio sources will be discovered in more distant galaxies. Hence the need to fully investigate the nature of these sources in nearby galaxies. ",
        "conclusions": "\\label{discuss} The VLA in its A-configuration connected with the Pie Town VLBA antenna via a real-time fibre optic link has been used to image M82 at 15 GHz. These observations have provided the highest resolution images yet made at this frequency and allowed a direct comparison with 5 GHz MERLIN images at a similar resolution. The conclusions may be summarised as follows. \\begin{enumerate} \\item In addition to the 26 compact radio sources which have been previously identified in the nucleus of M82, a further 20 sources have been identified. The final identifications are summarised in Table~\\ref{final_ID_table} (n.b. the controversial sources at 41.95+57.5 and 44.01+59.5 are included under `supernova remnants'). By comparison with the 12.8 $\\mu$m [NeII] map of \\citet{Achtermann95}, it was noted that the locations of the compact H{\\sc ii} regions broadly coincided with the peaks in the large-scale ionised gas distribution. \\begin{table} \\begin{center} \\caption{\\label{final_ID_table}The number of identified compact H{\\sc ii} regions and young supernova remnants in M82.} \\begin{tabular}{|c|c|c|c|}\\hline & Before                 & This work                     & TOTAL\\\\ & this work              &                               & \\\\ \\hline Supernova Remnants & 25                     & 5                             & 30   \\\\ H{\\sc ii} Regions  & 1                      & 15                            & 16   \\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\item Of the 15 new H{\\sc ii} regions identified, all have steep `inverted' spectra at the highest resolution ($\\alpha_5^{15} >$ +0.4). For the simple model of an H{\\sc ii} region in which S $\\propto \\nu^2$ in the optically-thick regime and S $\\propto \\nu^{-0.1}$ in the optically-thin regime this would imply that these sources are becoming optically thick between 5 and 15 GHz. However, \\citet{Olnon75} and \\citet{Panagia75} showed that simple models of ionised nebulae with significant electron density gradients could produce inverted spectra, the slope of which depends on the structure of the region. This may be the case for a number of the sources, but further high-resolution imaging at different frequencies is required to confirm or deny this possibility. \\item The low-resolution spectra of the regions in which the compact sources have been imaged show inverted spectra for about half of the sources and flattening spectra at the highest frequencies for the remainder. It is inferred that at the highest resolution only the most compact and optically-thick components of the ionised gas are being detected. \\item Of the 46 identified sources in M82, $\\sim$ 35$\\%$ are H{\\sc ii} regions and the remaining $\\sim$ 65$\\%$ probably have supernova origins. Therefore, it would appear that M82 is not lacking in compact H{\\sc ii} regions in comparison to other galaxies such as NGC~2146 \\citep{Tarchi00} as much as was previously thought, although it still appears that M82 is in a more advanced starburst stage than NGC~2146 due to the greater number of supernova remnants. \\item Five additional sources were identified as supernova remnants since they have steep, negative spectral indices at the higher frequencies. At the lower frequencies they are often experiencing a turnover in their spectra by 1.4~GHz requiring emission measures for the foreground ionised gas of $\\sim$ 10$^7$ cm$^{-6}$ pc. These supernova remnants are obscured at the lower radio frequencies where previous studies have concentrated and may represent those supernova remnants embedded deeper within M82 and hence behind a greater amount of ionised gas. \\item An additional 14 unidentified sources were added to the analysis and after taking into account the transient source noted by Kronberg, Biermann \\& Schwab (1985)\\nocite{Kronberg85}, the total number of compact sources now stands at 61. The 14 unidentified sources are most likely the older supernova remnants in the sample and hence are not bright enough to have been included in the essentially flux-limited analyses. However, assuming that they are all SNR, a comparison between the H{\\sc ii} region distribution and SNR distribution shows significant differences, although any difference in the overall extent of the distributions is not significant. Hence, the radio data cannot be used to either confirm or deny the propagating star-formation hypotheses of \\citet{Satyapal97} and \\citet{DeGrijs00}. \\end{enumerate} \\subsection*{Acknowledgements} MERLIN is a national facility operated by the University of Manchester on behalf of PPARC. The VLA is operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. ARM acknowledges the receipt of a PPARC postgraduate research grant."
    },
    "0207/astro-ph0207314_arXiv.txt": {
        "abstract": "{Using \\hst\\thanks{Based on observations with the NASA/ESA \\emph{Hubble Space Telescope}, obtained at the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract No. NAS 5--26555.} and ground-based optical and NIR imaging data\\thanks{Obtained at the German--Spanish Astronomical Center, Calar Alto, operated by the Max--Planck--Institute for Astronomy, Heidelberg, jointly with the Spanish National Commission for Astronomy.}$^,$ \\thanks{Obtained at the Kitt Peak National Observatory, operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.}, we investigate whether the blue compact dwarf (BCD) galaxy I Zw 18 possesses an extended low-surface-brightness (LSB) old stellar population underlying its star-forming regions, as is the case in the majority of BCDs. This question is central to the long-standing debate on the evolutionary state of I Zw 18. We show that the exponential intensity decrease observed in the filamentary LSB envelope of the BCD out to $\\ga$18\\arcsec\\ ($\\ga$1.3 kpc assuming a distance of 15 Mpc) is not due to an evolved stellar disc underlying its star-forming regions, but rather, due to extended ionized gas emission. Ionized gas accounts for more than 80\\% of the line-of-sight emission at a galactocentric distance of $\\sim$ 0.65 kpc ($\\sim$ 3 effective radii), and for $\\ga$ 30\\% to 50\\% of the $R$ light of the main body of I Zw 18. Broad-band images reveal, after subtraction of nebular line emission, a relatively smooth stellar host extending slightly beyond the star-forming regions. This unresolved stellar component, though very compact, is not exceptional for intrinsically faint dwarfs with respect to its structural properties. However, being blue over a radius range of $\\sim$ 5 exponential scale lengths and showing little colour contrast to the star-forming regions, it differs strikingly from the red LSB host of standard BCDs. This fact, together with the comparably blue colours of the faint C component, $\\sim$ 1.6 kpc away from the main body of I Zw 18, suggests that the formation of I Zw 18 as a whole has occurred within the last 0.5 Gyr, making it a young BCD candidate. Furthermore, we show that the ionized envelope of I Zw 18 is not exceptional among star-forming dwarf galaxies, neither by its exponential intensity fall-off nor by its scale length. However, contrary to evolved BCDs, the stellar LSB component of I Zw 18 is much more compact than the ionized gas envelope. In the absence of an appreciable underlying stellar population, extended ionized gas emission dominates in the outer parts of I Zw 18, mimicking an exponential stellar disc on optical surface brightness profiles. } ",
        "introduction": "} Since its discovery by Sargent \\& Searle (\\cite{SS70}), I Zw 18 has been looked at as the prototypical blue compact dwarf (BCD) galaxy. Its low oxygen abundance (Searle \\& Sargent \\cite{SS72}), established in numerous subsequent studies (Lequeux et al. \\cite{Leq79}; French \\cite{Fre80}; Kinman \\& Davidson \\cite{KD81}; Pagel et al. \\cite{Pagel92}; Skillman \\& Kennicutt \\cite{SK93}; Martin \\cite{Martin96}; V{\\'\\i}lchez \\& Iglesias-Par\\'amo \\cite{ViIp98}; Izotov \\& Thuan \\cite{IT98a,IT98b}; Izotov et al. \\cite{Yu99}) to be 12~+~log(O/H)~$\\approx$~7.2 makes it the least chemically evolved star-forming galaxy in the local Universe. Whether this is a signature of youth (cf. e.g. Izotov \\& Thuan \\cite{IT98a}) remains, however, a subject of debate.  I Zw 18 was described by Zwicky (\\cite{Fritz66}) as a pair of compact galaxies, subsequently shown to be two compact star-forming (SF) regions within the same galaxy with an angular separation of 5\\farcs8, the brighter northwestern (\\nw) and fainter southeastern (\\se) components. Ground-based and \\hst\\ observations have revealed that both regions are embedded in a low-surface brightness (LSB) filamentary envelope extending out to a radius $\\sim$20\\arcsec\\ (Davidson et al. \\cite{Dav89}; Dufour \\& Hester \\cite{Duf90}; Dufour et al. \\cite{Duf96a}; Martin \\cite{Martin96}; \\\"Ostlin et al. \\cite{Gor96}), and within an extensive \\H1\\ halo with a projected size of 60\\arcsec $\\times$ 45\\arcsec\\ (van Zee et al. \\cite{vZ98}; see also Viallefond et al. \\cite{Via87}).  The spectrophotometric properties of the fainter detached component I Zw 18\\,C, located $\\sim$ 22\\arcsec\\ northwest of the \\nw\\ region are still poorly known. Dufour et al. (\\cite{Duf96a}), Petrosian et al. (\\cite{Petr97}), Izotov \\& Thuan (\\cite{IT98a}), van Zee et al. (\\cite{vZ98}) and Izotov et al. (\\cite{Yu01a}) have shown it to have the same recession velocity as the main body, thus establishing its physical association to I Zw 18. The SF activity of I Zw 18\\,C is weak, its \\ha\\ equivalent width not exceeding $\\sim$60 \\AA\\ along the major axis (Izotov et al. \\cite{Yu01a}). In spite of deep Keck\\,II spectroscopy, Izotov et al. (\\cite{Yu01a}) failed to detect oxygen lines, so its oxygen abundance is not known.  Colour-magnitude diagram (CMD) studies, based on \\hst\\ WFPC2 images, suggest for the main body an age between several 10 Myr (Hunter \\& Thronson \\cite{HT95}; Dufour et al. \\cite{Duf96b}) and $\\sim$~1 Gyr (Aloisi et al. \\cite{Alo99}). \\\"Ostlin (\\cite{Gor00}) argues from \\hst\\ NICMOS $J$ and $H$ images that a fit to the $J$ vs. $J-H$ CMD is best achieved with a stellar population with age as high as 5~Gyr. As for I Zw 18\\,C, Dufour et al. (\\cite{Duf96b}) and Aloisi et al. (\\cite{Alo99}) derive an age of a few hundred Myr. \\begin{figure} \\begin{picture}(16,8.7) \\put(0,0){{\\psfig{figure=sf0.eps,width=8.5cm,angle=0,clip=}}} \\end{picture} \\caption[]{\\hst/WFPC2 archival image of I Zw 18 in the $R$ band. North is at the top and east to the left. The two star forming regions in I~Zw~18 are labelled \\nw\\ and \\se. Regions marked {\\sl loop} and {\\sl H$\\alpha$ arc} have been investigated spectroscopically in Izotov et al.~(\\cite{Yu01a}). The northwestern supershell in I Zw 18 and the detached irregular component I Zw 18\\,C are indicated. The larger circles with radii 14\\arcsec\\ and 22\\arcsec\\ are centered between the \\nw\\ and \\se\\ star forming regions.} \\label{f0} \\end{figure}  Kunth \\& \\\"Ostlin (\\cite{KO00}) conclude from optical and NIR surface photometry studies that I Zw 18 possesses an evolved and spatially extended stellar population underlying its SF regions. Their conclusion is based on a nearly constant relatively red $B-R$ $\\sim$ 0.6 mag and outwards increasing $B-J$ colour of the LSB envelope. Because the latter has a surface-brightness profile that can be well fitted by an exponential law, Kunth \\& \\\"Ostlin (\\cite{KO00}) ascribed the LSB emission to a stellar disc with an age of $\\sim$~5 Gyr.  On the theoretical front, Legrand (\\cite{Leg00}) and Legrand et al. (\\cite{Leg00a}) in an attempt to explain the paucity of nearby SF dwarf galaxies with oxygen abundance 12~+~log(O/H)~$<$~7.2 proposed that these systems form over the Hubble time through a continuous low-level SF process. In this scenario, I Zw 18 would have an extended stellar disc with $m_V$ $\\sim$ 20 mag and a mean surface brightness $\\sim$~28~$V$ \\sbb.  If the presence of a significant old stellar background can be demonstrated then I Zw 18 would be a standard BCD, and the hypothesis of it being a young galaxy must be abandoned. I Zw 18 would then be like the vast majority of BCDs which are evolved dwarf galaxies where star formation is occurring within an extended, circular or elliptical stellar host galaxy (Loose \\& Thuan \\cite{LT86}, hereafter LT86; Kunth et al. \\cite{KMV88}; Papaderos et al. 1996a, hereafter \\cite{P96a}). Such systems, classified iE/nE by LT86, account for $\\sim$ 90\\% of the local BCD population. Their red LSB underlying component dominates the surface brightness and colour distribution in their outer parts and contributes on average $\\sim$ 1/2 of their $B$ luminosity within the 25 $B$ \\sbb\\ isophote (Papaderos et al. \\cite{P96b}).  The view that I Zw 18 has properties similar to old BCDs, with the exception of its low oxygen abundance, is not well established, however. Stars populating the stellar disc postulated by Legrand et al. (\\cite{Leg00a}) and Kunth \\& \\\"Ostlin (\\cite{KO00}) have not been seen in I Zw 18 (Izotov \\& Thuan \\cite{IT01}). By contrast, in CMD studies of Local Group dwarfs (see e.g. Grebel (\\cite{Grebel00}) for a review) and of a few nearby BCDs (cf. Sect. \\ref{S3a}), a census of several thousands resolved stars unambiguously proves the existence of an evolved and spatially extended stellar background. In fact, the observational evidence for a dominant 1 -- 5 Gyr old stellar population in I Zw 18 rests on a sample of only about one dozen red point sources seen both in the optical and NIR ranges by \\hst\\ (Aloisi et al. \\cite{Alo99}; \\\"Ostlin \\cite{Gor00}). Ages quoted from this tiny statistical probe assume no or a uniform extinction, while non-uniform dust absorption in I Zw 18 has been discovered recently (Cannon et al. \\cite{Cannon01}).  The assumption by Kunth \\& \\\"Ostlin (\\cite{KO00}) that ionized gas does not dominate the LSB envelope of I Zw 18, thus its colours need not be corrected for gaseous emission before deriving ages, has been disputed by Papaderos et al. (\\cite{P01}) and Izotov et al. (\\cite{Yu01a}). The latter authors remarked that the mean $B-R$ and $B-J$ colours derived by Kunth \\& \\\"Ostlin (\\cite{KO00}) and Papaderos et al. (\\cite{P01}) in the outer regions of the main body are entirely consistent with pure ionized gas emission. This is also consistent with the results by Hunt et al. (\\cite{Hunt02}) who placed from deep NIR data an upper limit of $\\la$ 15\\% to the $J$ band light fraction of stars older than 500 Myr in I Zw 18.  The large range of ages derived previously, together with the unexplored role of extended gaseous emission as well as dust absorption led us to reexamine the photometric structure of I Zw 18. Our goal was to determine whether there exists a stellar LSB component underlying the filamentary envelope of I Zw 18 and extending well beyond the \\nw\\ and \\se\\ regions. If present, how does its colour and structural properties compare to those of standard BCDs, and what implications can be drawn for the evolutionary state of I Zw 18?  The paper is organized as follows. In Sect. \\ref{S2} we discuss the set of ground-based and \\hst\\ data included in this study and briefly describe the techniques used in the surface photometry analysis. We consider essential to study the photometric structure of I Zw 18 not in isolation but in the context of the main class of evolved BCDs. The photometric structure of these systems is discussed in Sect. \\ref{S3} on the example of the nearby iE BCDs Mkn 178 and VII Zw 403. Section \\ref{S4} focuses on I Zw 18, the distance of which is assumed throughout to be 15 Mpc (Izotov et al. \\cite{Yu01a}). In Sect. \\ref{S4a} we derive surface brightness profiles (SBPs) of its main body on the usual assumption that its LSB emission is predominantly of stellar origin. The properties of the LSB component after subtraction of nebular line emission (Sect. \\ref{S4b}) are studied in Sect. \\ref{S4c}. The photometric structure of I Zw 18\\,C is discussed in Sect. \\ref{S4d}. In Sect. \\ref{dis1} we compare the structural properties of I Zw 18 with those of standard BCDs. In Sect. \\ref{dis2} we investigate whether spatially extended ionized gas emission can mimic the SBP of a red exponential stellar disc. The evolutionary state of I Zw 18 in the light of the present results is discussed in Sect. \\ref{dis3}. Our conclusions are summarized in Sect. \\ref{Conclusions}. ",
        "conclusions": "} The central question to this study is (i) the evolutionary state of I Zw 18 and (ii) whether this galaxy is qualitatively different as compared to standard blue compact dwarf (BCD) galaxies concerning its photometric structure.  We first discuss the photometric structure of two nearby {\\sl bona fide} old BCDs, Mkn 178 and VII Zw 403, to provide a comparison basis for I Zw 18. Both galaxies exemplify that star formation (SF) in BCDs does not extend out to a fortuitous galactocentric distance but, as consistently implied by surface photometry and CMD studies, occurs within the 25 $B$ \\sbb\\ isophotal radius \\e25\\ of the underlying low-surface-brightness (LSB) host. The confinement of star-forming activity to the inner part of a BCD results in an appreciable colour contrast ($\\sim$0.6--1 $B-R$ mag) between the star-forming and LSB component.  Using ground-based $B$, $V$, $R$, $I$ and $J$ and archival \\hst/WFPC2 images we investigate next the photometric structure of I Zw 18. Our surface brightness profiles (SBPs) confirm the exponential intensity fall-off and the moderately red $B-R$ colour ($\\sim$~0.6 mag) found in the outer parts of I Zw 18 by Kunth \\& \\\"Ostlin (\\cite{KO00}) which led them to conclude that I Zw 18 possesses an evolved ($\\sim$~5 Gyr) and spatially extended stellar disc underlying its star-forming regions.  While the $B$ and $R$ SBPs alone suggest {\\sl prima facie} that I Zw 18 resembles most other BCDs in having an underlying old stellar population, inspection of more colours severely challenges this interpretation. The moderately red $B-R$ colour of the LSB envelope cannot be reconciled with the red $V-R$ ($>$~0.4 mag) and is in conflict with the blue $V-I$~$\\sim$~0 mag, $B-V$~$\\sim$~0.1 mag and $B-J$~$\\sim$~0.6 mag. No evolved population model is consistent with these colours, irrespective of star formation history, age or metallicity. This inconsistency can be removed by taking into account strong contamination of the optical light over galactocentric distances as large as $\\sim$~1.3 kpc by ionized gas emission. This procedure is further supported by the large ($>$~1300 \\AA) \\ha\\ equivalent widths observed in the periphery of the BCD, as well as by the good spatial correlation of the large H$\\alpha$ equivalent widths with $V-R$ and $V-I$ colours.  We have used archival \\hst\\ narrow-band [\\ion{O}{iii}]~$\\lambda$5007 and \\ha\\ images to correct broad-band $B$, $V$ and $R$ \\hst/WFPC2 images for nebular line emission. While complete removal of ionized gas continuum and line emission is not possible with the available narrow-band images, subtraction of the most prominent nebular emission lines is sufficient to allow us to explore the hypothesis of a stellar host galaxy in I Zw 18. We find that ionized gas contributes $\\ga$~30\\% to 50\\% of the $R$ emission of I Zw 18. The stellar and ionized gas continuum emission dominates the $R$ light only within the effective radius ($R^*\\approx$~0.2 kpc), whereas its line-of-sight contribution decreases to $<$~20\\% at a galactocentric radius $R^*>$~0.65 kpc. Consequently, the exponential intensity decrease derived in the LSB outer parts (0.6$\\la R^*\\,{\\rm (kpc)}\\la$1.3) of I Zw 18 from uncorrected broad-band images is mainly due to ionized gas emission.  Subtraction of nebular line emission reveals a relatively smooth and very blue stellar LSB envelope (denoted \\lsb), extending not much beyond the star-forming regions of I~Zw~18. Although our surface photometry does not go below $\\sim$~26 $B$ \\sbb, it does allow for a quantitative study of the structural properties and colours of this underlying host galaxy. We find that it is not exceptional when compared to intrinsically faint ultra-compact dwarfs, regarding its central surface brightness ($\\mu_{\\rm E,0}=20.7$ $B$ \\sbb) and exponential scale length ($\\alpha\\approx$~120 pc). This is also the case for I~Zw~18~C, which shows in its outer parts an exponential intensity decrease with $\\alpha\\approx$~100 pc and a marked flattening with respect to the exponential fit for $R^*\\la$~0.3 kpc.  Although it does not stand out in the $\\mu_{\\rm E,0}$~--~$\\alpha$ parameter space, the \\lsb\\ host of I Zw 18, being blue on a radius range of $\\sim$5 exponential scale lengths and showing little colour contrast to the star-forming regions, differs strikingly from the underlying component in standard BCDs. The colours, as determined at its southeastern tip where ionized gas emission is negligible, are consistent with an instantaneous burst or continuous star formation starting less than 0.5 Gyr ago. Since I Zw 18 C shows comparably blue colours at its reddest northwestern tip, we conclude that most of the stellar mass in I Zw 18 has formed within the last 0.5 Gyr.  Finally, we show that the exponential intensity decrease observed in the filamentary envelope of I Zw 18 is typical of ionized gaseous halos in star-forming dwarf galaxies. I Zw 18 is not exceptional among BCDs, neither by the extent nor the exponential scale length of its ionized envelope. However, in the absence of a significant underlying stellar population, extended ionized gas emission dominates the light in the periphery of I Zw 18, mimicking the SBP of a relatively red ($B-R\\sim$0.6 mag) old stellar disc."
    },
    "0207/astro-ph0207408_arXiv.txt": {
        "abstract": "Estimates of the Galactic coalescence rate (${\\cal R}$) of close binaries with two neutron stars (NS--NS) are known to be uncertain by large factors (about two orders of magnitude) mainly due to the small number of systems detected as binary radio pulsars.  We present an analysis method that allows us to estimate the Galactic NS--NS coalescence rate using the current observed sample and, importantly, to assign a statistical significance to these estimates and to calculate the allowed ranges of values at various confidence levels. The method involves the simulation of selection effects inherent in all relevant radio pulsar surveys and a Bayesian statistical analysis for the probability distribution of the rate. The most likely values for the total Galactic coalescence rate (${\\cal R}_{\\rm peak}$) lie in the range $2-60$ Myr$^{-1}$ depending on different pulsar population models. For our reference model 1, where the most likely estimates of pulsar population properties are adopted, we obtain ${\\cal R}_{\\rm tot} = 8_{-5}^{+9}$ Myr$^{-1}$ at a 68\\% statistical confidence level. The corresponding range of expected detection rates of NS--NS inspiral are $3_{-2}^{+4}\\times 10^{-3}$ yr$^{-1}$ for the initial LIGO and $18_{-11}^{+21}$ yr$^{-1}$ for the advanced LIGO. ",
        "introduction": "\\label{sec:intro}  The detection of the double neutron star (NS--NS) prototype PSR~B1913+16 as a binary pulsar \\cite{ht75} and its orbital decay due to emission of gravitational waves \\cite{tfm79} has inspired a number of quantitative estimates of the coalescence rate of NS--NS binaries \\cite{cvs79,nps91,ph91,cl95}. In general, the coalescence rate of NS--NS binaries can be calculated based on: (a) our theoretical understanding of their formation (see Belczynski \\& Kalogera 2001 \\nocite{bk01a} for a review and application of this approach); (b) the observational properties of the pulsars in the binary systems and the modeling of pulsar survey selection effects (see \\nocite{nar87} e.g.~Narayan 1987). Interest in these coalescence derives from an intrinsic motivation of understanding their origin and evolution and their connections to other NS binaries. However, significant interest derives from their importance as gravitational-wave sources for the upcoming ground-based laser interferometers (such as LIGO) and their possible association with $\\gamma$-ray burst events (Popham et al.~1998 \\nocite{pop98} and references therein).  The traditional way of calculating the coalescence rate based on observations involves an estimate of the {\\it scale factor}, an indicator for the number of pulsars in our Galaxy with the same spin period and luminosity \\cite{nar87}. Corrections must then be applied to these scale factors to account for the faint end of the pulsar luminosity function, the beamed nature of pulsar emission, and uncertainties in the assumed spatial distribution.  The estimated total number in the Galaxy can then be combined with estimates of their lifetimes to obtain a coalescence rate, ${\\cal R}$. This method was first applied by Narayan et al.\\ (1991) and Phinney (1991) and other investigators who followed \\cite{cl95,vl96}. Various correction factors were (or were not) included at various levels of completeness. Summaries of these earlier studies can be found in Arzoumanian et al.\\ (1999) and Kalogera et al.\\ (2001; hereafter KNST). The latter authors examined all possible uncertainties in the estimates of the coalescence rate of NS--NS binaries in detail, and pointed a small-number bias that introduces a large uncertainty (more than two orders of magnitude) in the correction factor for the faint-pulsar population that must be applied to the rate estimate.  They obtained a total NS--NS rate estimate in the range ${\\cal R}=10^{-6} - 5\\times 10^{-4}$\\,yr$^{-1}$, with the uncertainty dominated by the small-number bias.  Earlier studies, which made different assumptions about the pulsar properties (e.g.~luminosity and spatial distributions and lifetimes), are roughly consistent with each other (given the large uncertainties). Estimated ranges of values until now were not associated with statistical significance statements and an ``all-inclusive'' estimated Galactic coalescence rate lies in the range $\\sim10^{-7} - 10^{-5}$ yr$^{-1}$.  The motivation for this paper is to update the scale factor calculations using the most recent pulsar surveys, and present a statistical analysis that allows the calculation of statistical confidence levels associated with rate estimates.  We consider the two binaries found in the Galactic disk: PSR~B1913+16 \\cite{ht75} and PSR~B1534+12 \\cite{wol91}.  Following the arguments made by Phinney (1991) and KNST, we do not include the globular cluster system PSR~B2127+11C \\cite{pr91}; this system will be the subject of a later paper. Radio-pulsar-survey selection effects are taken into account in the modeling of pulsar population. As described in what follows, the small-number bias and the effect of a luminosity function are {\\it implicitly included} in our analysis, and therefore a separate correction factor is not needed.  For each population model of pulsars, we derive the probability distribution function of the total Galactic coalescence rate weighted by the two observed binary systems. In our results we note a number of important correlations between ${\\cal R}_{\\rm peak}$ and model parameters that are useful in generalizing the method. We extrapolate the Galactic rate to cover the detection volume of LIGO and estimate the detection rates of NS--NS inspiral events for the initial and advanced LIGO.  The plan for the rest of this paper is as follows.  In \\S \\ref{sec:method}, we describe our analysis method in a qualitative way.  Full details of the various pulsar population models and survey selection effects are then given in \\S \\ref{sec:models} and \\S \\ref{sec:selfx} respectively.  In \\S \\ref{sec:analysis}, we derive the probability distribution function for the total Galactic coalescence rate and calculate the detection rate of LIGO. In \\S \\ref{sec:results}, we summarize our results and discuss a number of intriguing correlations between various physical quantities. Finally, in \\S \\ref{sec:discussion}, we discuss the results and compare them with previous studies. ",
        "conclusions": "\\label{sec:discussion}   In this paper we present a new method of estimating the total number of pulsars in our Galaxy and we apply it to the calculation of the coalescence rate of double neutron star systems in the Galactic field. The method implicitly takes into account the small number of pulsars in the observed double-neutron-star sample as well as their distribution in luminosity and space in the Galaxy. The modeling of pulsar survey selection effects is formulated in a ``forward'' way, by populating the Galaxy with model pulsar populations and calculating the likelihood of the real observed sample. This is in contrast to the ``inverse'' way of the calculation of scale factors used in previous studies. The formulation presented here allows us to: (a) calculate the probability distribution of coalescence rates; (b) assign statistical significance to these estimates; and (c) quantify the uncertainties associated with them.  As originally shown by KNST, the most important uncertainties originate from the combination of a small-number observed pulsar and a pulsar population dominated by faint objects. The probability distribution covers more than 2 orders of magnitude in agreement with the uncertainties in excess of two orders of magnitudes asserted by KNST. However, for the reference model, even at high statistical confidence level (99\\%) the uncertainty is reduced to a factor of $\\sim 60$. At confidence levels of 95\\% and 68\\%, the uncertainty is further reduced to just $\\sim 25$ and $\\sim 5$, respectively. We use our results to estimate the expected detection rates for ground-based interferometers, such as LIGO. The most likely values are found in the range $\\sim (1 - 30) \\times 10^{-3}$\\,yr$^{-1}$ and $\\sim 4 - 140$\\,yr$^{-1}$, for the initial and advanced LIGO.  The statistical method developed here can be further extended to account for distributions of pulsar populations in pulse periods, widths, and orbital periods. Most importantly the method can be applied to any type of pulsar population with appropriate modifications of the modeling of survey selection effects.  Currently we are working on assessing the contribution of double neutron stars formed in globular clusters as well as the formation rate of binary pulsars with white dwarf companions that are important for gravitational-wave detection by LISA, the space-based interferometer planned by NASA and ESA for the end of this decade."
    },
    "0207/astro-ph0207122_arXiv.txt": {
        "abstract": "Deep near-infrared images obtained with adaptive optics (AO) systems on the Gemini North and Canada-France-Hawaii telescopes are used to investigate the bright stellar content and central regions of the nearby elliptical galaxy Maffei 1. Stars evolving on the upper asymptotic giant branch (AGB) are resolved in a field 3 arcmin from the center of the galaxy. The locus of bright giants on the $(K, H-K)$ color-magnitude diagram is consistent with a population of stars like those in Baade's Window reddened by $E(H-K) = 0.28 \\pm 0.05$ mag. This corresponds to A$_V = 4.5 \\pm 0.8$ mag, and is consistent with previous estimates of the line of sight extinction computed from the integrated properties of Maffei 1. The AGB-tip occurs at $K = 20.0$, which correponds to M$_K = -8.7$; hence, the AGB-tip brightness in Maffei 1 is comparable to that in M32, NGC 5128, and the bulges of M31 and the Milky-Way. The near-infrared luminosity functions (LFs) of bright AGB stars in Maffei 1, M32, and NGC 5128 are also in excellent agreement, both in terms of overall shape and the relative density of infrared-bright stars with respect to the fainter stars that dominate the light at visible and red wavelengths. It is concluded that the brightest AGB stars in Maffei 1, NGC 5128, M32, and the bulge of M31 trace an old, metal-rich population, rather than an intermediate age population. It is also demonstrated that Maffei 1 contains a distinct red nucleus, and this is likely the optical signature of low-level nuclear activity and/or a distinct central stellar population. Finally, there is an absence of globular clusters brighter than the peak of the globular cluster LF in the central $700 \\times 700$ parsecs of Maffei 1. ",
        "introduction": "The brightnesses and spatial distributions of stars evolving on the asymptotic giant branch (AGB) provide clues about the past evolution of galaxies. Although stars near the AGB-tip are relatively bright, efforts to resolve these objects in the dense main bodies of nearby spheroids typically require angular resolutions approaching the diffraction limit of 2.5-metre or larger telescopes. Spheroids within the Local Group are obvious first targets for any studies of resolved stellar content, and Davidge (2000a), Davidge et al. (2000), and Davidge (2001a) found that the brightest AGB stars in the compact elliptical galaxy M32 and the bulge of M31 have similar brightnesses, and are well mixed with the fainter stars in these systems. This suggests that the most luminous AGB stars in M32 and the bulge of M31 belong to a population that formed when the structural characteristics of these galaxies were imprinted (Davidge 2001a). A similar AGB population may be present in the bulge of the Milky-Way (Davidge 2001a) -- a system that deep photometric studies indicate has an old age (Feltzing \\& Gilmore 2000; Ortolani et al. 1995).  Davidge (2001a) suggested that the bright AGB stars detected in Local Group spheroids trace an old, metal-rich population, and this suggestion can be tested by examining the brightness and spatial distribution of AGB stars in a larger sample of spheroids. With a distance between 4 and 4.5 Mpc (Davidge \\& van den Bergh 2001; Luppino \\& Tonry 1993), Maffei 1 is one of the closest large elliptical galaxies, and hence is an obvious target for efforts to study the bright stellar content. A complicating factor is that Maffei 1 is viewed through the Galactic disk, and so is subject to significant foreground extinction. Buta \\& McCall (1983) concluded that A$_V = 5.1 \\pm 0.2$ based on the integrated color of Maffei 1 and the column density of hydrogen along the line of sight. Hence, efforts to resolve stars in Maffei 1 will likely be most successful in the near-infrared, which is also the prime wavelength regime for investigating stars on the upper AGB. Davidge \\& van den Bergh (2001; hereafter DvdB) investigated the near-infrared photometric properties of the brightest AGB stars in a field with a projected distance of 6 arcmin from the center of Maffei 1. Despite long integration times (3 hours per filter) and image quality close to the theoretical diffraction limit of the CFHT, only stars within $\\sim 1.5$ mag of the AGB-tip were detected.  In the present study, two datasets are used to investigate different regions of Maffei 1. The first dataset consists of deep $H$ and $K'$ images obtained with the University of Hawaii Adaptive Optics (UHAO) system on the Gemini North telescope, which sample a field 3 arcmin from the center of Maffei 1. Not only is the stellar density higher than in the field studied by DvdB, but stars that are 1 mag in $K$ fainter than in the DvdB dataset are detected, permitting a detailed comparison with the bright stellar contents of other galaxies.  The second dataset consists of $J, H,$ and $Ks$ images of the center of Maffei 1 that were obtained with the CFHT AO system. While not diffraction-limited, these data have a resolution that permits the central $700 \\times 700$ parsecs of the galaxy to be investigated at sub-arcsec angular scales. Data of this nature can determine if Maffei 1 has a photometrically distinct nucleus, which in turn provides clues about the past evolution of this galaxy.  The paper is structured as follows. Details of the observations and the data reduction techniques are presented in \\S 2, while the photometric properties of stars in the Gemini dataset are discussed and compared with those of other galaxies in \\S 3 and 4. The photometric properties of the central regions of Maffei 1 are examined in \\S 5, while the results of a search for globular clusters in both datasets is presented in \\S 6. A summary and discussion of the results of this paper follows in \\S 7. ",
        "conclusions": "\\subsection{The Stellar Content of Maffei 1 and Other Spheroids}  Deep $H$ and $K'$ images obtained with the UHAO $+$ QUIRC on the Northern Gemini telescope have been used to investigate the bright AGB content of a field 3 arcmin ($\\sim 4$ kpc) from the center of Maffei 1. If it is assumed that the brightest giants in Maffei 1 have the same intrinsic $H-K$ color as late M giants in BW then a line-of-sight extinction is computed that is consistent with previous estimates, which have relied largely on the integrated properties of the galaxy.  The main result of this paper is that the infrared-bright stellar content of the Maffei 1 deep field, as gauged by (1) the brightness of the AGB-tip, (2) the shape of the AGB LF, and (3) the density of AGB stars measured with respect to surface brightnesses at visible wavelengths, does not differ significantly from that in other nearby spheroids. The AGB-tip in the Maffei 1 deep field occurs near M$_K \\sim -8.7$, and thus is comparable to the peak brightnesses in M32 (Davidge 2000a; Davidge et al. 2000), and the bulges of the Milky-Way and M31 (Davidge 2001a). Rejkuba et al. (2001) find that the peak M$_K$ in the outer regions of NGC 5128 is $\\sim -8.8$, which is also in remarkable agreement with the peak brightness in Maffei 1.  The near-infrared LFs of bright stars in the Maffei 1 deep field, and the outer regions of M32 and NGC 5128 are in excellent agreement. In some respects, the good agreement between the bright stellar contents of Maffei 1 and NGC 5128 is perhaps not surprising, given that these galaxies have comparable integrated brightnesses, distances, and environments. However, the chemical enrichment history of a galaxy is thought to depend on factors such as galaxy mass (e.g. Yoshii \\& Arimoto 1987), and it might be anticipated that the photometric properties of the brightest stars in a massive elliptical galaxy like Maffei 1 might differ from those in a smaller system like M32, due to differences in metallicity. Indeed, the integrated Mg$_2$ index of M32 is markedly lower than in more massive ellipticals (Burstein et al. 1984). However, the metallicity distribution function (MDF) of M32 measured by Grillmair et al. (1996) is similar to that of the outer regions of NGC 5128 (Harris \\& Harris 2000; Harris, Harris, \\& Poole 1999), suggesting that the stellar contents of M32 and larger ellipticals may not be so different.  Insight into the nature of the brightest stars in nearby spheroids can be obtained by examining their distribution within these systems. In M32 and the bulge of M31 the brightest stars are uniformly distributed throughout these systems, with a number density that scales with $r-$band surface brightness (Davidge 2000a, Davidge et al. 2000, and Davidge 2001). There are indications that the brightest stars also are uniformly distributed in NGC 5128, as Harris \\& Harris (2000) show that the relative number density of the brightest AGB and RGB stars does not change with radius, although it is evident from their Figures 7 and 8 that their data are not sensitive to luminous giants with solar or higher metallicities. Curiously, a comparison of the $K$ LFs of the DvdB field and the deep field suggests that the outer regions of Maffei 1 may be deficient in stars near $K = 20.5$, suggesting that the brightest stars in Maffei 1 may not be uniformly distributed throughout the entire galaxy. A survey of the outer regions of Maffei 1, covering a square arcmin or more and sampling stars as faint as $K = 21$ would provide the data that is needed to confirm if the deficiency of bright stars in the DvdB field is real, or a statistical fluke.  Soria et al. (1996) detected stars as bright as M$_{bol} \\sim -5$ in the inner halo of NGC 5128, and concluded that these objects belong to an intermediate-age population. Marleau et al. (2000) reached a similar conclusion after analyzing near-infrared observations of a portion of the Soria et al. field, and it is these data that have been compared with the Maffei 1 deep field observations. However, peak AGB luminosity is not an ironclad means of judging the age of a population, as the peak AGB brightness is a function of metallicity as well as age, and this introduces uncertainty in the age calibration of the AGB-tip. In fact, Guarnieri, Renzini, \\& Ortolani (1997) examined the brightest members of moderately metal-rich globular clusters, which have old ages (e.g. Ortolani et al. 1995), and found that the brightest AGB stars have M$_{bol}$ between --4.5 and --5.0 when [Fe/H] $> -1.0$; thus, the bright stars detected by Soria et al. (1996) may have ages comparable to Galactic metal-rich globular clusters.  The galaxy-to-galaxy similarity in peak M$_K$ and steller density are difficult to explain if the brightest stars are young or of intermediate-age, as these systems must then have experienced fortuitously similar star-forming histories during intermediate epochs: not only would age and metallicity have to be tuned to produce similar peak AGB luminosities, but the intermediate-age components would also have to be uniformly distributed throughout these systems with similar spatial densities. Both of these difficulties vanish if the bright stars are old; in this case the problem of tuning the AGB-tip brightness is less acute because the rate of change of this parameter with time decreases with increasing age. Likewise, the uniform distribution of infrared-bright stars occurs naturally if they formed during the initial collapse of the system, when the main structural characteristics of the galaxies were defined and there was likely a system-wide period of star formation. Finally, stars with a peak brightness like that in Maffei 1, M32, and the bulge of M31 occur in the Galactic Bulge (Davidge 2001a), which appears to have an old age (Feltzing \\& Gilmore 2000; Ortolani et al. 1995).  Clearly, NGC 5128 has experienced recent star formation, with the younger populations being centrally concentrated. However, based on the comparison with Maffei 1, the main body of this galaxy is old. It thus appears that NGC 5128 is a nearby example of the `frosting' model proposed to explain the integrated spectroscopic properties of many early-type galaxies in the field, in which a modest young or intermediate age population is superimposed on an old stellar substrate (Trager et al. 2000b).  Hierarchal models of galaxy formation, which assume that large galaxies are assembled by the accretion of smaller systems, are able to reproduce many observed properties of present-day galaxies (e.g. Somerville \\& Primack 1999; Cole et al. 2000). One prediction of these models is that 50\\% of all stars formed prior to z = 1.5 (Cole et al. 1994; 2000). It can be anticipated that most of the stars (or their remnants) that formed prior to z = 1.5 will be in spheroidal systems at the current epoch, since mergers and feedback from star formation likely prevented disks from forming until z $\\sim 1$ (e.g. Weil, Eke, \\& Efstathiou 1998). That spheroids are dominated by stars that formed early-on is consistent with the Mg$_2 - \\sigma_0$ relation of these systems, which indicates that their basic structural properties were imprinted at high redshift (Bernardi et al. 1998).  It is somewhat suprising that the brightest stars appear to be uniformly distributed throughout the main bodies of systems like M32, the bulge of M31, and NGC 5128, as the evolution of a region within a galaxy is influenced by the local mass density, which defines the escape velocity and (possibly) the star formation rate (e.g. Schmidt 1959). A radial variation in escape velocity may be the physical basis for metallicity gradients in elliptical galaxies (Franx \\& Illingworth 1991; Martinelli, Matteucci, \\& Colafrancesco 1998), as well as the tight relations between absorption line strengths and local surface brightness (Kobayashi \\& Arimoto 1999; Davidge \\& Grinder 1995). Local surface brightness is a relative measure of mass density, at least to the extent that spheroidal systems have similar M/L ratios. The surface brightnesses of the various fields that have been compared in this paper and in Davidge (2001a) are summarized in Table 1, and it is evident that these span a wide range of values. These data ostensibly suggest that the progenitors of the bright AGB stars studied in this paper can form in regions with surface brightnesses in M$_V$ at least as low as $\\sim 1$ mag pc$^{-2}$ ($\\sim 30$ M$_{\\odot}$ pc$^{-2}$).  \\subsection{The Central Regions of Maffei 1}  The data presented in this paper indicate that Maffei 1 contains a red nucleus, that extends out to $\\sim 1\\arcsec$. The nature of this nucleus is not clear, although there are hints that it is not a low level AGN. The absence of a discrete x-ray point source in Maffei 1 has been noted by Reynolds et al. (1997). In addition, Spinrad et al. (1971) discuss the only spectroscopic observations of Maffei 1 that are known to us. Their spectrum, obtained with a $2\\arcsec$ wide slit, shows strong line absorption, with no hint of central line or continuum emission.  While there is no evidence for a systematic age gradient in Maffei 1 (\\S 5), the presence of a young nucleus can not be completely discounted. However, if the nucleus is younger than the main body of the galaxy then it must be heavily extincted and/or viewed at an evolutionary stage when the AGB dominates the infrared light. Buta \\& McCall (1999) do find dust north of the nucleus of Maffei 1.  As a moderately large (M$_V \\sim -21.6$) elliptical galaxy, Maffei 1 should have a well-populated globular cluster system. However, the central $700 \\times 700$ parsecs of Maffei 1 is devoid of globular clusters brighter than the peak of the GCLF. While dynamical evolution is expected to disrupt clusters in the central regions of galaxies (e.g. Portegies Zwart et al. 2001; Murali \\& Weinberg 1997; Vesperini 1997), some nearby elliptical galaxies have bright globular clusters within a few hundred parsecs of their centers. Forbes et al. (1996) investigated the central globular cluster contents of a sample of elliptical galaxies. One of the nearest galaxies in their sample is NGC 4494, which has an integrated brightness similar to Maffei 1; 6 clusters brighter than the peak of the GCLF were found within 400 parsecs of the center of this galaxy. Interestingly, despite having what appears to be a well-populated central cluster system, NGC 4494 may have a lower than average global specific cluster frequency (Larsen et al. 2001).  The specific globular cluster frequency of Maffei 1 is not known. Because of the heavy extinction at visible wavelengths any survey for globular clusters in Maffei 1 should likely be conducted in the near-infrared. Based on the relation between the globular cluster system core radius and host galaxy brightness calibrated by Forbes et al. (1996), the core radius of the Maffei 1 cluster system should be $\\sim 2.5$ kpc, so a number of clusters should be present within $\\sim 2$ arcmin of the galaxy center. Foreground star contamination is an obvious concern, although this does not present an insurmountable hurdle, since field stars with brightnesses comparable to those of bright clusters in Maffei 1 have relatively blue colors (\\S 6). Hence, it should be possible to distinguish between clusters and stars using $J-K$ colors."
    },
    "0207/astro-ph0207252_arXiv.txt": {
        "abstract": "{ The nature and the location of the lenses discovered in the microlensing surveys done so far towards the LMC remain unclear. Motivated by these questions we compute the optical depth and particularly the number of expected events for self-lensing for both the MACHO and EROS2 observations. We calculate these quantities also for other possible lens populations such as thin and thick disk and galactic spheroid. Moreover, we estimate for each of these components the corresponding average event duration and mean mass using the mass moment method. By comparing the theoretical quantities with the values of the observed events it is possible to put some constraints on the location and the nature of the MACHOs. Clearly, given the large uncertainties and the few events at disposal it is not possible to draw sharp conclusions, nevertheless we find that certainly at least 3-4 MACHO events are due to lenses in LMC, which are most probably low mass stars, but that hardly all events can be due to self-lensing. This conclusions is even stronger when considering the EROS2 events, due to their spatial distribution. The most plausible solution is that the events observed so far are due to lenses belonging to different intervening populations: low mass stars in the LMC, in the thick disk, in the spheroid and possibly some true MACHOs in the halo. ",
        "introduction": "The location and the nature of the microlensing events found so far towards the Large Magellanic Cloud (LMC) is still a matter of controversy. The MACHO collaboration found 13 to 17 events in 5.7 years of observations, with a mass for the lenses estimated to be in the range $0.15 - 0.9 \\, \\mathrm{M_{\\sun}}$  assuming a standard spherical Galactic halo (Alcook et al. 2000a) and derived an optical depth of $\\tau= 1.2^{+0.4}_{-0.3} \\times 10^{-7}$. An analysis of the spatial distribution of the events as well as a detailed study of the source star location done on HST images (Alcook et al. 2001a) are both consistent with an extended lens distribution such as the Milky Way halo, however an LMC distribution is only slightly disfavoured. Thus this test is not conclusive given the few events at disposal. The EROS2 collaboration (Milsztajn \\& Lasserre 2001) announced the discovery of 4 events, with an estimated average mass of 0.2 $\\mathrm{M_{\\sun}}$, based on three years of observation but monitoring about twice as much stars as the MACHO collaboration. The MACHO collaboration monitored primarily 15 deg$^2$ in the central part of the LMC, whereas the EROS experiment covers a larger solid angle of 64 deg$^2$ but in less crowded fields. The EROS microlensing rate should thus be less affected by self-lensing. This might be the reason for the fewer events seen by EROS as compared to the MACHO experiment.  It has been argued that some if not all the events could be due to LMC self-lensing. Indeed, several authors have already studied in detail self-lensing by focusing, however, mainly on the value of the microlensing optical depth. Sahu (1994) and Wu (1994) suggested that self-lensing within the LMC could explain the observed optical depth. This claim has been questioned by further studies (Gould 1995; Alcock et al. 1997a). A major problem is the uncertainty related to a precise knowledge of the shape and the total mass of the LMC. Indeed, for instance a disaligned bar, as suggested by Weinberg, could increase the self-lensing optical depth (Weinberg 2000; Evans \\& Kerins 2000; Zhao \\& Evans 2000) as well as assuming the LMC being much more extended along the line of sight (Aubourg et al. 1999). These authors find then an optical depth for self-lensing of $0.5 - 1.5 \\times 10^{-7}$, thus comparable with the measured value by the MACHO collaboration. Other authors, instead, find lower values for the optical depth in the range $0.5 - 8.0\\times 10^{-8}$ (see for instance Gyuk et al. 2000). The discrepancy are of course due to the different adopted models for the shape of the LMC, which are based on various arguments and by weighing differently the observation on the LMC star distribution which are available today.  We notice that a direct comparison of the theoretical values for the optical depth to self-lensing and the measured one is not straightforward. Indeed, the theoretical values are computed for a given line of sight and then to compare with the measured value one takes the average over the lines of sight corresponding to the various monitored fields (Gyuk et al. 2000). On the other hand the measured optical depth is computed assuming that it stays constant over the monitored fields of the LMC. Such a procedure is certainly adequate if the lenses are in the halo and thus their optical depth does practically not vary on the size of the LMC, however, this is no longer true when dealing with self-lensing.  Some of the events found by the MACHO team are most probably due to self-lensing: the event MACHO-LMC-9 is a double lens with caustic crossing (Alcock et al. 2000b) and its proper motion is very low, thus favouring an interpretation as a double lens within the LMC; the source star for the event MACHO-LMC-14 is double (Alcock et al. 2001b) and this has allowed to conclude that the lens is most probably in the LMC. The expected LMC self-lensing optical depth due to these two events has been estimated to lie within the range $1.1-1.8\\times10^{-8}$ (Alcock et al. 2001b), which is still below the expected optical depth for self-lensing even when considering models giving low values for the optical depth.  The event LMC-5 is due to a disk lens (Alcock et al. 2001c) and indeed the lens has even been observed with the HST. The lens mass is either $\\sim 0.04 ~\\mathrm{M_{\\sun}}$ or in the range $0.095 - 0.13~\\mathrm{M_{\\sun}}$, so that it is a true brown dwarf or a M4-5V spectral type low mass star. The other stars which have been microlensed were also observed but no lens could be detected, thus implying that the lens cannot be a disk star but has to be either a true halo object or a faint star or brown dwarf in the LMC itself.  Some work has also been done on the Small Magellanic Clouds (SMC) where only two microlensing events have been found up to now (Alcock et al. 1997b, 1999; Palanque-Delabrouille et al. 1998). One was a resolved binary event, allowing the determination of the lens distance (Alcock et al. 1999; Afonso et al. 1998; Albrow et al. 1999), which most probably resides in the SMC itself and thus clearly being due to self-lensing. The other event is of long duration and a detailed analysis seems also to favour a self-lensing interpretation (Palanque-Delabrouille et al. 1998). It has been argued that the SMC self-lensing optical depth should be higher than the corresponding LMC value since the SMC is tidally disrupted, due to its interaction with the Milky Way and the LMC, and is thus quite elongated along the line of sight to the Milky Way (Caldwell \\& Coulson 1986; Welch et al. 1987).  Thus up to now the question of the location of the observed MACHO events is unsolved and still subject to discussion. Clearly, with much more events at disposal one might solve this problem by looking for instance at their spatial distribution. However, since the MACHO collaboration data taking stopped at the end of 1999 and the EROS experiment is still underway, but will hardly lead to a large amount of events in the next few years, there is not much hope to have substantial more data at disposal within the next few years and it is thus of importance to explore further ways which can give insight into the MACHO location using only the already existing data. This is the main aim of this paper.  As a first point we calculate the optical depth, the average duration and particularly the number of expected events for self-lensing. For the latter quantity we take for the number of monitored stars and for the exposure time the values corresponding to the MACHO and EROS experiments, respectively. Moreover, we compute the same quantities for MACHOs located in a thin or thick disk or in a spheroid around our Galaxy. As a main point we compute following the mass moment method (De Rujula et al. 1991, 1992) the average mass of the lenses for the observed events under the assumption that the lenses are located in the LMC, the halo or in one of the Galactic components. This allows to see whether the assumption of having the lenses in the LMC itself or in one of the Galactic components leads to consistent values for the masses. Our analysis indicates that most probably not all lenses originate from the same population but are due to different ones. Therefore, some of the conclusions reported so far in the literature, especially on the average mass and thus on the nature of the lenses, have to be taken with caution. Finally, we give an estimate of the fraction of the local dark mass density detected in form of MACHOs in the Galactic halo.  The paper is organized as follows: in Sect. 2 we give the density profiles of the stellar distributions for the different components of the LMC, which will then be used in the following. In Sect. 3 we present the different equations to compute the optical depth, the microlensing rate and the mass moments we shall use. In Sect. 4  we report the results for the LMC self-lensing, whereas in Sect. 5 we present the values we find for the various Galactic lens populations: thin and thick disks, spheroid and halo. The mass of the lenses is discussed in Sect. 6. In Sect. 7, by using different models, we determine the fraction of the local dark mass density detected in form of MACHOs in the Galactic halo. In Sect. 8 we analyse the asymmetry in the spatial distribution of the observed events. We conclude in Sect. 9 with a summary of our results. ",
        "conclusions": "As mentioned at the beginning the issue of the location and nature of the objects which act as lenses in the observed microlensing events is still an open problem, which can possibly be solved once more events will be available. The number of events found by the experimental collaborations are too few in comparison with that predicted for a halo composed entirely by MACHOs. In the last years several possible explanations of the experimental data have been proposed, but they are all not definitive and not completely satisfactory. One possibility might be that there are more MACHOs in the halo, but that, for instance, they are associated with gas cloud (Bozza et al. 2002), which would then produce non-achromatic events, that due to the present selection criteria have not been considered. On the other hand, one should also consider the possibility that some events might not be due to microlensing at all. This obviously underlines the fact that the present results have to be taken with care and that more observations are needed in order resolve this issue.  From the presently few available events and our above discussion it emerges clearly that, especially for the MACHO events, the lenses are due to different populations. Some are certainly due to LMC self-lensing, but this can hardly be the case for all the observed events especially for the few EROS events.  For the preferred LMC model with a dispersion velocity of about $30~ \\mathrm{km\\, sec^{-1}}$ we expect some 3 events among the MACHO ones (not including the possible self-lensing binary event LMC-9) to be due to self-lensing. Moreover, with the mass moment method we find an average mass in the range $(0.1 - 0.5)\\, \\mathrm{M_{\\sun}}$ for the self-lensing events, which is consistent with the expectation that the lenses are low mass stars. The contribution of a LMC halo is, even if it exists, also very minor of at most 1 event unless we are in the rather strange situation where the LMC halo is, contrary to our own, made almost entirely of MACHOs. From the galactic component, thin and thick disk, we expect roughly $1-3$ events on the MACHO data, but no contribution on the EROS data. This result seems to be in agreement with the fact that LMC-5 is a disk event. The inferred mass is small, though compatible with a low mass star, but clearly with only one event at disposal one has to take this value just as indicative.  As a result of our analysis we find that a plausible solution is that among the MACHO data some $3-4$ events are due to self-lensing, $1-2$ to the thick disk and for the remaining ones about half are due to the spheroid and the others to the halo. If so the optical depth gets contributions from each of these components, namely: about $2.3 \\times 10^{-8}$ from self-lensing, some $(2-4) \\times 10^{-8}$ from the halo, $4 \\times 10^{-8}$ from the spheroid and $4 \\times 10^{-8}$ from thin and thick disk. This way we get a total optical depth of about $(1.2 - 1.5) \\times 10^{-7}$ which is in good agreement with inferred value by the MACHO team of $\\tau=1.2_{-0.3}^{+0.4}\\times 10^{-7}$. Since the EROS events are less and given their spatial position, we do not expect that they get much contribution from self-lensing and the thick disk so that it is clearly not possible to draw more conclusions from them. On the other hand assuming that the EROS events are due to lenses in the halo, with or without some contribution from the spheroid, leads to a halo mass fraction which is in reasonable agreement with the corresponding MACHO value once the self-lensing and the disk events are subtracted. This way the MACHO and EROS results nicely fit together. Moreover, the above mentioned asymmetry, if not just due to a statistical fluctuation, is also not compatible with only self-lensing events. Given our results it is also clear that one has to take with care values on the lens mass based on the assumption that all lenses belong to just one population. Clearly, once more data will be available, by using also the methods outlined in this paper, it will be possible to draw more firm conclusions."
    },
    "0207/astro-ph0207544_arXiv.txt": {
        "abstract": "{ We measure the eddy viscosity in the outermost layers of the solar convection zone by comparing the rotation law computed with the Reynolds stress resulting from f-plane simulations of the angular momentum transport in rotating convection with the observed differential rotation pattern. The simulations lead to a {\\em negative } vertical  and a {\\em  positive } horizontal angular momentum transport. The consequence is a subrotation of the outermost layers, as it is indeed indicated both by helioseismology and the observed rotation rates of sunspots. In order to reproduce the observed gradient of the rotation rate a value of about $1.5 \\times 10^{13}$ cm$^2$/s for the eddy viscosity is necessary. Comparison with the magnetic eddy diffusivity derived from the sunspot decay yields a surprisingly large magnetic Prandtl number of 150 for the supergranulation layer. The negative gradient of the rotation rate also drives a surface meridional flow towards the poles, in agreement with the results from Doppler measurements. The successful  reproduction of the {\\em abnormally positive} horizontal cross correlation (on the northern hemisphere) observed for bipolar groups then provides an independent test for the resulting eddy viscosity. ",
        "introduction": "Over the last years, while observing the solar oscillations on longer timescales and with higher precision, it has become evident that these oscillations play an important role in understanding the solar interior, bearing more or less the only information from the deeper parts of the sun, which cannot be probed otherwise. Helioseismology reveals a {\\em maximum} of the angular velocity at {\\em all} latitudes rather close to the surface, as shown in Fig.~\\ref{f1} (Howe et al.~2000).  It  is known, on the other hand, that sunspots rotate faster than the solar surface plasma by about 4 \\% or 80 m/s at all latitudes. Such a clearly verified subrotation of the outermost layer of the convection zone is easiest understood as a result of angular momentum conservation of fluid elements with purely radial motions. But in this domain of the solar convection zone the velocity field is dominated by horizontal motions. Fluctuating fields with predominantly horizontal intensity should produce superrotation rather than subrotation. There is, however, another strong argument for considering the exceptional behavior of the {\\em horizontal motions} in more detail.  Ward (1965) was the first to consider the horizontal cross-correlation of the proper motions of sunspot groups, the faster of which tend to move toward the equator. He found \\begin{equation} Q_{\\theta\\phi} \\approx \\ 0.1\\  ({\\rm deg/day})^2 \\approx 2 \\times 10^7 {\\rm cm}^2/ {\\rm s}^2 \\label{0} \\end{equation} on the northern hemisphere. More recent observations found smaller, but always positive values (Gilman \\& Howard 1984; Nesme-Ribes \\ea 1993; Komm \\ea 1994, see an overview  by Meunier et al. 1997). This result has a strong implication for theory confirming the existence of the positive $H$ coefficient in the expression (\\ref{3}) for the horizontal Reynolds stress. \\begin{figure}[ht] \\center \\mbox{ \\psfig{figure=gong.ps,width=7cm,height=5cm} } \\label{f1} \\caption{The internal rotation of the Sun as found by helioseismology. Image: NSF's National Solar Observatory } \\end{figure} We shall demonstrate the close relation between the negative radial gradient of the rotation rate and the positive sign of the horizontal cross correlation by solving the Reynolds equation on the basis of new data from hydrodynamic simulations of rotating convection in the outermost layer of the convection zone. The computations provide a tool for measuring the eddy viscosity in the outer solar convection zone. ",
        "conclusions": "The theory of turbulent angular momentum transport in stellar convection zones based on the Second Order Correlation Approximation (SOCA) in KR93 predicts a positive vertical and vanishing horizontal \\L-effect in the solar supergranulation layer, and negative vertical as well as positive horizontal \\L-effect in the bulk of the solar convection zone. Solutions of the Reynolds equation with the stress tensor from KR93 therefore reproduce the observed variation of the rotation rate with latitude remarkably well, but lack the decrease of the rotation rate with increasing radius in the outermost part of the convection zone.  Simulations of rotating convection in the upper part of the solar convection zone show a strong and positive horizontal and a negative vertical \\L-effect. With the \\L-coefficients as derived from the simulations, the solutions show a negative shear, $\\partial \\Omega/\\partial r < 0$, of the observed amplitude when a value of $1.5 \\times 10^{13} {\\rm cm}^2/{\\rm s}$ is chosen for the eddy viscosity. As there is no way to directly measure the eddy viscosity, this is the only method to derive its value from observations.  As an independent test, we have computed theoretical Ward profiles. With the \\L-effect from KR93, the horizontal cross-correlation is always negative because the horizontal \\L-effect vanishes in the surface layer. With the large positive value of $H$ from the simulations, on the other hand, $Q_{\\theta \\phi}$ is always positive at low latitudes, as observed, and the amplitudes are in good agreement as well.  In models of the solar differential rotation with positive radial \\L-effect in the outermost layers of the convection zone, the radial shear is always positive in that layer, and the surface gas flow is directed towards the equator, both in contradiction to the observations. The negative radial \\L-effect derived from the Chan (2001) simulations removes both these contradictions by maintaining a negative gradient of the rotation rate, which in turn drives the surface flow towards the poles.  We conclude that the KR93 expressions for the Reynolds stress are invalid in the outermost part of the solar convection zone. Possible reasons are the proximinty of the outer boundary, the increasing importance of radiative energy transport with decreasing depth, and the neglect of the inherent anisotropy of turbulent convection, which is driven by a large-scale entropy gradient rather than a random force."
    },
    "0207/astro-ph0207258_arXiv.txt": {
        "abstract": "We present optical broad- and narrow-band imaging of a sample of a dozen barred galaxies. These images are analyzed in conjunction with our previously published near-infrared imaging of their central regions and with literature values for, e.g., bar strengths and the total star formation activity of the galaxies. We present $B$, $I$ and H$\\alpha$ images, and radial profiles derived from these, to infer geometric and dynamical parameters of the structural components of the galaxies, such as bar lengths, bar ellipticities, and location of star formation and dust. We find that the more centrally concentrated the \\ha\\ emission in a galaxy is, i.e., the higher the fraction of star formation originating in the circumnuclear region, the higher the overall star formation rate, as measured from far-infrared flux ratios.  Stronger bars host smaller nuclear rings, but the strength of the bar does not correlate with either the intrinsic ellipticity of the ring or the offset between the position angles of the bar and the ring. We interpret these results in comparison with modelling of gas inflow in the circumnuclear region, and show that they were theoretically expected.  We confirm observationally, and for the first time, the anti-correlation predicted from theory and modelling between the degree of curvature of the bar dust lanes and the strength of the bar, where stronger bars have straighter dust lanes. ",
        "introduction": "Most galaxies are barred (e.g., Sellwood \\& Wilkinson 1993; Knapen, Shlosman \\& Peletier 2000a) and a substantial fraction of barred galaxies show enhanced star formation (SF) activity in or near their centres, often in the form of complete or incomplete nuclear rings (e.g., Buta \\& Combes 1996; Knapen 1999). It is believed that there is a causal connection between the existence of a bar and circumnuclear SF activity. Through its non-axisymmetric potential, a bar can facilitate gas inflow by extracting angular momentum from the gas through gravitational torques. The inflowing gas may then accumulate in the vicinity of inner Lindblad resonances (ILRs), triggering massive SF (see, e.g., review by Shlosman 1999). It is therefore natural to infer that the properties of the circumnuclear starburst region are connected to those of the large-scale structure, most obviously the bar, and it is this inference that we wish to test observationally in the present paper.  Correlations have been established between morphological properties of the stellar bar, such as size, shape and surface brightness (Martin 1995; Elmegreen \\& Elmegreen 1985), the location of star forming regions (Sersic \\& Pastoriza 1967; Phillips 1993, 1996), the presence of nuclear activity (Chapelon, Contini \\& Davoust 1999; Knapen et al. 2000a; Shlosman, Peletier \\& Knapen 2000; Laine et al. 2002) and the Hubble type of their host galaxy. Bars are found to be longer in early-type barred galaxies (Martin 1995), longer bars are more elliptical (Martinet \\& Friedli 1997), bar radial profiles appear to be flatter in early-type galaxies and exponential in late-type galaxies (Elmegreen \\& Elmegreen 1985), and there is some evidence that the more elliptical the bar is, the higher is its overall SF rate (Aguerri 1999). There also appears to be some correlation between the location of SF in galaxies and their morphological classification. Phillips (1993) noted that early-type barred galaxies exhibit SF in rings but neither in the bar nor the centre, whereas late-type barred spirals display many star forming regions along their bars.  Sersic \\& Pastoriza (1967) realised that enhanced SF is found in the inner parts of some spiral galaxies, and that this SF is frequently arranged into a ring or pseudo-ring pattern. Combes \\& Gerin (1985) and Athanassoula (1992a) found that such spiral galaxies are usually of early type and host a bar. Rings in galaxies, including nuclear rings, were extensively reviewed by Buta \\& Combes (1996). Buta \\& Crocker (1993) compiled a catalog of rings in galaxies, including ten nuclear rings studied in the present paper, and derived metric characteristics for the rings. Case studies combining optical and near-infrared (NIR) imaging with velocity fields and dynamical modelling show in detail how the bar and the circumnuclear regions (CNRs) are part of the same dynamical system. Gas is driven inward under the influence of a large stellar bar, where it piles up near or between the ILRs, and possibly results in massive SF. Examples of such case studies are those by Knapen et al.  (1995a,b, 2000b) on M100 (=NGC~4321), Regan et al. (1996, 1997) on NGC~1530, Buta, Crocker \\& Byrd (1999) on ESO~565-11 (see also Rautiainen \\& Salo 2000), and Laine et al. (1999, 2001) and Jogee et al. (2002a,b) on NGC~5248. In thin bars (high axial ratio) the gas inflow may be most effective, but according to the simulations by Piner, Stone \\& Teuben (1995), nuclear rings will not form in such bars and gas will flow directly to the nucleus.  Nuclear bars, spirals, or rings, are frequently observed on scales of a kiloparsec or less, and can be modelled as directly resulting from the dynamics of the overall system (see Knapen et al. 2001 for many examples).  We have performed an imaging study of a dozen barred spirals with evidence for the presence of star-forming CNRs, to investigate the coupling between the CNRs and their host galaxies in a more statistical manner.  Our main goal is to test the hypothesis (Knapen et al. 1995b; P\\'erez-Ram\\'\\i rez \\& Knapen 1998) that the structure and dynamics of the large-scale bar and disk of the host galaxies is intimately connected with that of the inner, active region (either AGN or SF), and thus determines whether such activity occurs in a given galaxy, and in what form.  In Paper~I (P\\'erez-Ram\\'\\i rez et al. 2000), we presented NIR images of the central regions of a sample of 12 barred galaxies to study the properties of circumnuclear structures in stars and dust that exist at small scales (a few hundred pc to a few kpc). In this paper, we present optical broad-band and H$\\alpha$ imaging of the whole extent of the disks of the same 12 barred galaxies, and combine parameters derived from the NIR and optical parts of our imaging survey with parameters taken from the literature to identify several correlations which illuminate the important connections that exist between the central regions and the host galaxy structures at large.  We summarise the observations and data reduction procedures in $\\S$~2, and present our observational results in $\\S$~3. Relative \\ha\\ fluxes from different parts of the sample galaxies are explored in $\\S$~4. In subsequent sections, we explore correlations between the strength of the bar and the geometric parameters of the nuclear rings ($\\S$~5), and the shapes of the dust lanes within the bar ($\\S$~6). We summarise our main findings in $\\S$~7. Notes on individual objects can be found in Appendix~1. ",
        "conclusions": "This paper is the second of a pair presenting results of our study of how the general structure of barred galaxies is related to, and possibly determines, the presence and properties of dynamical features such as rings and spirals in their CNRs.  We present optical broad and narrow-band imaging of our sample galaxies in $B, I$ and H$\\alpha$ in order to map the structural components across the disks of these barred galaxies. The 12 galaxies in our sample are all nearby spiral galaxies known to host star-forming circumnuclear ring-like structure. We present graphically our optical images and the radial profiles we derived from these by ellipse fitting. From these data, we infer geometric and dynamical parameters such as bar lengths and ellipticities, and location of SF and general dust structures, and combine these with parameters obtained from the literature.  Our main results can be summarised as follows:  \\begin{itemize}  \\item For the first time using bar strengths ($Q_{\\rm b}$ values from the literature) rather than ellipticity, and using NIR imaging of the bars, we checked the previously found correlation between bar strength and length, i.e., longer bars are also stronger. Our results are confusing because bar length as such correlates with bar strength, but bar length relative to, e.g., the size of the galaxy does not. Careful study of larger samples is needed.  \\item The FIR index log($S_{25}$/$S_{100}$), a SF indicator, is used to verify whether the strength of the bar is connected to the global SF activity. We find no relation with the bar strength $Q_{\\rm b}$ nor with the length of the bar.  \\item We measure the relative H$\\alpha$ flux from the disk, bar and CNR components in the galaxies. We find that the more centrally concentrated the \\ha\\ emission is (i.e., the higher the fraction of total SF from the CNR), the more SF takes place in absolute terms in the galaxy, as estimated from the FIR flux ratios. This is possibly due to the IRAS index being more sensitive to CNR SF, which is concentrated in a dust-rich environment. Alternatively, the CNR SF activity as measured from \\ha\\ is related to the overall SF activity, as measured in the FIR, of the disk. We find no correlation between the relative H$\\alpha$ flux from the disk, the bar, and the CNR with bar strength, nor with bar length.  \\item We explore relations between the strength of the bar and the geometric parameters of the nuclear rings, namely their size relative to the galaxy disk, intrinsic inclination, and PA offset between the major axes of the nuclear ring and the bar which hosts it. We find that stronger bars tend to host smaller rings, but that the bar strength is unrelated to both the PA offset and the intrinsic ring ellipticity. The fact that smaller rings occur in stronger bars can be understood from bar orbit theory, which predicts that stronger bars limit the region where $x2$ orbits can exist, and thus in turn the size of the nuclear ring. As known from numerical modelling, both PA offset and intrinsic ring ellipticity depend primarily on the mass inflow rate into the nuclear ring, which in turn is determined only in part by the strength of the bar, but more directly on the availability of fuel, and on temporal variations therein.  \\item We present a new, reproducible, measure of the curvature of dust lanes in bars, a parameter which was theoretically linked about a decade ago to the strength of the bar. This link is through gas shocking under the influence of the bar potential, at the location where pairs of symmetric dust lanes are seen along the bars in optical images. We indeed confirm that the stronger the bar, the straighter its dust lanes. This is in fact the first observational confirmation of this theoretical result, based on the use of both a sample (admittedly still rather small) of galaxies, and of quantitative measures of both dust lane curvature and bar strength.  \\end{itemize}  From these results, we can confirm that indeed the CNRs of barred galaxies are intimately linked to their host galaxies, and that their SF and morphological properties are determined to a significant degree by the properties of, mainly, the large bar. The small size of our sample precludes a detailed exploration of the correlations that we have found. Studies using larger samples of nuclear ring galaxies are needed to more fully test the predictions and results from theory and modelling, and to explore further relations between the parameters governing the CNR, the bar, and the disk.  {\\it Acknowledgements}  We thank Milagros Ru\\'\\i z and Sharon Stedman for help during the observing runs and for discussions, and Reynier Peletier for his help in the preparation of Figure~1. We thank Isaac Shlosman for illuminating discussion on bars and rings, and David Axon for discussions about subtracting the H$\\alpha$ continuum. We thank Ron Buta for calculating $Q_{\\rm b}$ values for NGC~1530 and NGC~6951. An anonymous referee is acknowledged for comments which greatly helped to improve the presentation of our results. 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. The Jacobus Kapteyn and Isaac Newton Telescopes are operated on the island of La Palma by the ING in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\\'\\i sica de Canarias. Data were partly retrieved from the ING archive."
    },
    "0207/astro-ph0207128_arXiv.txt": {
        "abstract": "We report the detection of four new extra-solar planets from the Anglo-Australian Planet Search orbiting the somewhat metal-enriched stars HD\\,73526, HD\\,76700, HD\\,30177 and HD\\,2039. The planetary companion of HD\\,76700 has a circular orbit with a period of 3.98\\,d. With M\\,$\\sin i$=0.197$\\pm$0.017\\Mjup, or 0.69 times the mass of Saturn, is one of the lowest minimum mass extra-solar planets yet detected. The remaining planets all have elliptical orbits with periods ranging from 190.5\\,d to 4.4\\,yr. All four planets have been found orbiting stars from a sub-sample of twenty metal-enriched and faint (V$<$9) stars, which was added to the Anglo-Australian Planet Search's magnitude-limited V$<$7.5 main sample in October 1998. These stars were selected to be metal-enriched on the basis of their Str\\\"omgren photometry, and their enrichment has been subsequently confirmed by detailed spectroscopic analysis. ",
        "introduction": "The Anglo-Australian Planet Search (AAPS) is a long-term planet detection program which aims to perform extra-solar planet detection and measurement at the highest possible precision. Together with programmes using similar techniques on the Lick 3\\,m and Keck I 10\\,m telescopes \\citep{fischer01,vmba00}, it provides all-sky planet search coverage for inactive F,G,K and M dwarfs down to a magnitude limit of V=7.5. Initial results from this programme \\citep{AAPSI,AAPSII,AAPSIII,AAPSIV,AAPSV,AAPSVI} demonstrate  that it achieves long-term, systematic velocity precisions of 3\\,\\ms\\ or better, for suitably stable stars, down to our main sample magnitude limit of V$<$7.5.  AAPS is being carried out on the 3.9\\,m Anglo-Australian Telescope (AAT), using the University College London Echelle Spectrograph (UCLES) and an I$_2$ absorption cell. UCLES is operated in its 31\\,lines\\,mm$^{-1}$ mode. Prior to 2001 September, it was used with  a MIT/LL 2048$\\times$4096 15$\\mu$m pixel CCD, and since then has been used with an EEV 2048$\\times$4096 13.5$\\mu$m pixel CCD. Our target sample includes 178 FGK stars with $\\delta < -20$\\arcdeg\\ and V$<$7.5. Where age/activity information is available from \\rhk\\ indices -- see for example \\citet{hsdb96,CaHKI} -- we require target stars to have $\\log$\\rhk $>$ -4.5 (corresponding to ages greater than 3\\,Gyr). The observing and data processing procedures follow that described in \\citet{bmwmd96} and \\citet{AAPSII}.  In addition to our primary sample of V$<$7.5 dwarfs, a small sub-sample of twenty fainter dwarfs (down to V$<$9)  was added in October 1999, following suggestions that metal-enriched stars seemed to be preferentially revealing planets (see eg. \\citet{laughlin00} and references therein). These stars all had $uvby$ photometry suggesting metal-enrichment. All stars in this fainter sub-sample were observed with a maximum exposure time of 300s regardless of observing conditions. As a result, the velocity precisions achieved for these targets are not as high as those demonstrated for the AAPS main sample \\citep{AAPSIV,AAPSII}. ",
        "conclusions": "The resultant minimum companion mass for HD\\,76700 is \\Msini\\,=\\,0.197$\\pm$0.017\\,\\Mjup, with an orbital semi-major axis $a$\\,=\\,0.049$\\pm$0.004\\,au and eccentricity $e$=0.00$\\pm$0.04. This zero eccentricity is consistent with the expectation that a planet with a period of just 3.971$\\pm$0.001\\,d will almost certainly lie in an orbit which has been tidally circularised \\citep{mb98}. The resulting orbital parameters for HD\\,76700 place it amongst the planetary companions with the lowest known minimum masses. HD\\,76700, joins HD\\,49674, HD\\,16141, HD\\,168746 \\& HD\\,46375 \\citep{butler02b, mbv00, pepe02} in the group of extra-solar planetary companions with measured minimum masses less than a Saturn mass (0.299\\,\\Mjup).  The remaining three extra-solar planets all have elliptical orbits, though the ellipticity of the orbit for the companion to HD\\,30177 ($e$=0.22$\\pm$0.17) is not different from zero with great significance. As further data are acquired over the coming years this parameter will become far better constrained, and it is possible that this extra-solar planet could turn out to be in a substantially circular orbit. If so this system would join with the other known nearly circular systems with gas-giant planets lying in orbits between where the Earth and Jupiter lie in our own Solar System ($\\epsilon$\\,Ret, HD\\,4208, the outer components of 47\\,UMa, HD\\,28185 \\citep{AAPSIV}). These extra-solar planets would seem to indicate that gas-giants exist in nearly circular orbits with semi-major axes all the way out to, and beyond that of Jupiter, as confirmed recently by the detection of the outer planet in the 55\\,Cnc system \\citep{55Cnc}.  Of the twenty metal-enriched stars included as a sub-sample along with our main sample in late 1999, five have revealed the presence of planetary companions (the four planets discussed here -- HD\\,30177, 73526, 76700 and 2039 -- along with the previously known companion to HD\\,83443 \\citep{AAPSV,mayor_83443}). This gives us a lower limit (there may be longer period or lower mass planets present which we cannot yet detect) to the discovery rate of 25$\\pm$11\\% for this ``metallicity-biassed'' sub-sample. This compares with the overall discovery rates estimated for the Keck, Lick and AAPS of $\\sim$8\\% (ie 8\\% of stars surveyed have planets in orbits within 3.5\\,\\au\\ of their host stars \\citep{properties}) -- a difference which, while not of great statistical significance, is not unexpected given that extra-solar planets seem to be being found preferentially around metal-enriched stars (eg. \\citet{reid02,laughlin00} and references therein). It is also interesting to note (even if perhaps not statistically significant) that all four of the stars in this paper would seem to be beginning their evolution off the main sequence (see Fig \\ref{models}).  These results would suggest that the biassing of planet surveys toward metal-enriched host stars {\\em may} offer a benefit in the planet detection rate. However, such an increased discovery rate must be balanced against the fact that it will produce an inherently biassed sample of extra-solar planets. With the total number of extra-solar planets still numbering less than a hundred, and the parameters of this ensemble of planets still poorly placed in a scheme of extra-solar planetary formation and evolution, now is not the time for planet searches to begin biassing their large surveys in the chase for better ``hit rates'' at the expense of scientific utility."
    },
    "0207/astro-ph0207402_arXiv.txt": {
        "abstract": "We have analyzed changes in the acoustic oscillation eigenfrequencies measured over the past 7 years by the GONG, MDI and LOWL instruments. The observations span the period from 1994 to 2001 that corresponds to half a solar cycle, from minimum to maximum solar activity.  These data were inverted to look for a signature of the activity cycle on the solar stratification. A one-dimensional structure inversion was carried out to map the temporal variation of the radial distribution of the sound speed at the boundary between the radiative and convective zones. Such variation could indicate the presence of a toroidal magnetic field anchored in this region.  We found no systematic variation with time of the stratification at the base of the convection zone. However we can set an upper limit to any fractional change of the sound speed at the level of $3 \\times 10^{-5}$. ",
        "introduction": "Changes in the frequency of the solar $p$-mode oscillations have now been observed for more than a decade. Such changes affect both the central frequencies, $\\nu_{n \\ell}$, and the frequency splittings, $\\Delta \\nu_{n \\ell m}$ of low degree \\citep{bi40}, intermediate degree \\citep{bi41,bi43} and very high degree modes \\citep{bi42}. A number of mechanisms have been proposed to explain these variations on frequency. \\citet{bi1} have argued, on the basis of observations of intermediate degree modes, that the source of the perturbations must lie near the solar surface. \\citet{bi2} and \\citet{bi3} concluded that magnetic fields located near the base of the convection zone, with strengths significantly lower than $10^6$ G have no observable effect on $p$-mode frequencies. The stability analysis for magnetic fields by \\citet{bi5} has shown that fields with strengths significantly larger than $10^5$ G cannot be stored in this region.  The $p$-mode frequency variations track rather well the changes of the activity strength of the solar cycle with time. It is thus plausible that these frequency shifts are due to variations of the mean magnetic field near the photosphere \\citep{bi2,bi8}. The influence of thin magnetic fibrils on the frequency shifts has been investigated by \\citet{bi15}. Another possible cause for these frequency shifts is the presence of sunspots during solar activity \\citep{bi14}. The dominant effect of sunspots on the propagation of acoustic waves is believed to be the dissipation of the acoustic energy and therefore, should decrease their amplitudes and lifetimes \\citep{bi44,bi42}, while it has been observed that frequencies increase with increased magnetic activity.  In the work presented here, we explore a mechanism first suggested by \\citet{bi6}, in which the sound speed perturbation associated to the observed changes of the photospheric latitudinal temperature distribution might be responsible for the frequency shifts seen during the solar cycle. Changes in the temperature distribution are themselves due to the heat transport through the convection zone induced by the solar dynamo.  Let us suppose that there is a magnetic field anchored below the base of the convection zone. To maintain pressure equilibrium, the gas pressure and thus the density inside the magnetized region must be lower than in its surroundings. The magnetized fluid will thus experience a larger radiative heating and therefore the temperature at the top of the magnetized region will increase. This will also induce a change in the temperature gradient that could be large enough to make the region above the magnetic field convectively unstable. In such a scenario, the base of the convection zone would locally drop, allowing the magnetized fluid to ascend by convective upflows, transporting excess entropy to the photosphere \\citep{kus}.  Numerical experiments by \\citet{kus} have shown that entropy perturbations in the deep convection zone can produce strongly peaked temperature changes in regions below $\\tau = 1$ that have a substantial acoustic signature (where $\\tau$ is the optical depth). They have also shown that the thermal perturbations that account for the solar acoustic variability are consistent with the observed solar irradiance and luminosity changes that occur during the 11 year solar cycle.  Luminosity changes, even if no larger than $0.1 \\%$, must come from the release of energy stored somewhere in the solar interior and must be accompanied by a change in the solar radius. The ratio between relative luminosity and radius changes, hereafter $W$, can help estimate the location of the region where this energy is stored \\citep[and references therein]{go2000}. Theoretical calculations indicate that $W$ increases when increasing the depth of the source of the variations in luminosity. For instance $W \\approx 2 \\cdot 10^{-4}$ if the source is located in the outer layers of the convection zone, while $W \\approx 0.5$ if the source is located in the solar core. Unfortunately there is a large scatter in the observed values of $W$. Indeed, recent measurements of $W$ range from 0.021 as estimated by \\citet{bi45} and \\citet{bi46} to an upper limit of 0.08 derived by \\citet{emi}.  \\citet{ku} estimated that a $0.1 \\%$ luminosity perturbation integrated over a solar cycle corresponds to about $10^{39}$ erg. If this energy originates in the tachocline and if the tachocline thickness is $0.05\\,\\Rsun$, the associated relative variation in sound speed at that depth would be on the order of $\\delta c / c \\approx 10^{-5}$ or $10^{-6}$. Fractional changes in sound speed as small as $10^{-4}$ are easily accessible by helioseismic inversion techniques. Some attempts to find solar-cycle variations of the sound speed asphericity and the latitude-averaged sound speed have been carried out using MDI and/or GONG data \\citep{bi45,bi41,bi47,bi48,bi49}. However, none of them found any systematic variation of the solar structure at the base of the convection zone that could be associated with the presence of a local toroidal magnetic field.  We have extended the previous analysis to the latest data available, including LOWL data. Only common modes to all data sets were used, in an attempt to obtain comparable and significant results to all instruments. This way, we can give a robust upper limit on the temporal variations of the solar internal stratification during the period 1994-2001. ",
        "conclusions": ""
    },
    "0207/astro-ph0207634_arXiv.txt": {
        "abstract": "We show the initial results of our 3D MHD simulations of the flow of the Galactic atmosphere as it responds to a spiral perturbation in the potential. In our standard case, as the gas approaches the arm, there is a downward converging flow that terminates in a complex of shocks just ahead of the midplane density peak. The density maximum slants forward at high $z$, preceeded by a similarly leaning shock. The latter diverts the flow upward and over the arm, as in a hydraulic jump. Behind the gaseous arm, the flow falls again, generating further secondary shocks as it approaches the lower $z$ material. In cases with two arms in the perturbing potential, the gaseous arms tended to lie somewhat downstream of the potential minimum. In the four arm case, this is true at large $r$ or early evolution times. At smaller $r$, the gaseous arms follow a tighter spiral, crossing the potential maximum, and fragmenting into sections arranged on average to follow the potential spiral. Structures similar to the high $z$ part of the gaseous arms are found in the interarm region of our two-armed case, while broken arms and low column density bridges are present in the four-armed case. Greater structure is expected when we include cooling of denser regions.  We present three examples of what can be learned from these models. We compared the velocity field with that of purely circular rotation, and found that an observer inside the galaxy should see radial velocity deviations typically greater than $20 \\kms$. Synthetic spectra, vertical from the midplane, show features at velocities $\\approx -20 \\kms$, which do not correspond to actual density concentrations. Placing the simulated observer outside the galaxy, we found velocity structure and arm corrugation similar to those observed in H$\\alpha$ in NGC 5427. ",
        "introduction": "Even though spiral structure is one of the most prominent features of disk galaxies, details of the spiral arms in our own Galaxy remain uncertain. \\citet{geo76} traced the spiral structure of the Milky Way using \\ion{H}{2} regions, and developed a model with four arms. More recent attempts concluded that the Milky Way might actually have a superposition of two and four arm structures, each one with different pitch angles, which might arise from different components of the galactic disk \\citep{dri00, lep01}, suggesting that the stellar and gaseous disks might not be tightly coupled. Similar behavior has been frequently observed in external galaxies \\citep[for example]{pue92, gro98}.  \\citet{rob69} showed that the gas must generate a large scale shock in the presence of a spiral perturbation. It was proposed that the density enhancement induced by this shock might generate a sequence of molecular clouds and star formation downstream from the shock, which itself was associated with the strong dust lane observed in the inner region of the spiral arms in external galaxies. Two dimensional numerical models by \\citet{tub80} and \\citet{sou81} showed that the gas forms a vertical shock perpendicular to the plane of the galactic disk. The post-shock gas remained close to hydrostatic equilibrium, even with an adiabatic equation of state. Their results did not show vertical motions larger than $3 \\kms$. In fact, the largest downflow they found was due to the pre-shock gas readjusting its vertical structure as it flows into the arm potential. Therefore, when \\ion{H}{1} observations on face-on galaxies showed extended velocity components with dispersions of the order of $20 \\kms$, they were attributed to other phenomena, such as galactic fountains, a warping of the \\ion{H}{1} disk, or intermediate velocity clouds \\citep{dic90, kam92, kam93}.  Since then, we have realized that the ISM is thicker and with a higher pressure than previously thought. The pressure scale height has been found to be larger than the density scale height, and the non-thermal pressures (turbulent, magnetic and cosmic ray) are at least as large as the thermal component \\citep{bad77, rey89, bou90}. Therefore, less compressible gas needs to be considered in order to generate more realistic models of the ISM. Such a medium, with a larger effective $\\gamma$ (the ratio of the specific heats) would be more likely to display the vertical motions characteristic of a hydraulic jump. With this in mind, \\citet[MC]{mar98} performed 2D MHD simulations of the flow of the gaseous disk and found diverse structures that differed from the vertical near-hydrostatics found in previous studies. In many cases, the gas moved up ahead of the stellar arm, sped up over it, and fell behind with large bulk velocity. Frequently, there was a downstream shock at higher $z$ as this downflow was arrested, sometimes resulting in secondary midplane density maxima.  The goal of our investigation is to extend calculations like those of MC to three dimensions, to a large fraction of the Galaxy, and to look for its possible observational signatures. In this paper, we present the early results of these simulations. In Section 2 we describe the numerical setup and the procedure to achieve the initial hydrostatic equilibrium, in Section 3 we describe the results of the simulations, in Section 4 we present three examples of synthetic observations that can be done with this type of simulation, and in Section 5 we present our conclusions. ",
        "conclusions": "In this work, we present our early results in the 3D MHD modeling of the large scale interaction of the ISM with a spiral potential. The presence of a thicker, more pressurized gaseous disk, together with the extra freedom the gas has in 3D simulations, allows the generation of density and velocity structures that previous work failed to reveal.  We confirmed and extended the work by MC, in which large scale vertical motions of the gas are an intrinsic feature of the response to the spiral perturbation. The downflow occurs along a much broader region and at higher velocities than the upflow. The falling gas can have large regions with a very similar vertical component, which translates into velocity crowding. In the present models, this gas appears as peaks at about $20 \\kms$ in spectra taken directly up from the midplane. These motions are accompanied with rapid flow above the arms and similar ``up and over'' motions in the radial direction. So far, there are some hints that such motions might be occuring in NGC 5427 \\citep{alf01}.  We also found significant differences in the midplane line-of-sight velocity distribution as compared with a purely circular rotation model. We think that the presence of streaming motions generated by the spiral arms must be considered when estimating the distance to elements of the ISM using their velocity as reference. In the future, when we obtain a more realistic model for the Milky Way, we may be able to provide a reasonable recipe for translating radial velocities and galactic longitude data to distance in a more reliable way.  Our models have a number of numerical simplifications (low resolution, closed boundaries, short run times) and omission of physical processes (heating and cooling of the gas, cosmic rays, self gravity, ionization, star formation or associated energy injection). Improvement on the run times and resolution will allow us to follow the structures to maturity, better examine cyclic features, explore substructure formation such as feathers, bridging and gaseous interarms, follow the magnetic field energy density and geometry to saturation, and to explore radial migrations of material and angular momentum. Addition of a more realistic equation of state to represent heating and cooling of the gas will allow the formation of truly dense regions. The interaction between magnetized flow and these regions may qualitatively alter the general arm structure, the velocity field and the complexity of the magnetic field configuration.  Our results show that failure to consider high $z$ and non-circular motions of the ISM associated just with the response to the spiral potential can easily lead to confusion when interpreting observational results. The study of the gaseous structure of the Milky Way and other galaxies require the consideration of three dimensional effects and a more realistic model of the nature of the ISM and its interaction with other dynamical elements of the system."
    },
    "0207/astro-ph0207152_arXiv.txt": {
        "abstract": "A new astronomical window into the southern skies has been opened with the high-frequency upgrade to the Australia Telescope Compact Array (ATCA), which allows radio-interferometric mapping of sources at wavelengths as short as 3mm.  In anticipation of the upgrade's completion, a two-day workshop was held at the University of Melbourne in November 2001.  The workshop covered a diverse range of fields, tied together by a common theme of identifying key areas where ATCA observations can have an impact.  More than half of the talks were concerned with molecular clouds and star formation, with the remainder covering topics such as molecular gas in the Galactic Centre, Seyfert nuclei, and high-redshift objects.  Some early results from the 3mm and 12mm prototype systems were also presented.  In consultation with the speakers, we are presenting in this article a summary of the talks.  The original slides are available from the ATNF website. ",
        "introduction": "Millimetre (mm) astronomy has grown rapidly in the past decade, spurred by advances in receiver sensitivity and growing scientific interest in what might be termed the ``cold universe.''  While molecular lines are still the ``bread and butter'' of millimetre astronomy, a wide range of continuum sources can also be studied at these wavelengths, and today mm astronomy spans nearly the full range of astrophysical research, from solar-system studies to cosmology.  The success of centimetre-wave interferometry, demonstrated by facilities such as the Westerbork Synthesis Radio Telescope (WSRT), led to the development of millimetre interferometers in the 1970's and 1980's.  There are currently four dedicated mm arrays, all located in the Northern Hemisphere at latitudes of 35\\arcdeg\\ to 45\\arcdeg.  A much larger array, to be built in Chile by an international consortium, will provide a manyfold increase in collecting area over all existing mm arrays put together.  This Atacama Large Millimeter Array (ALMA) will be completed around 2011.  In the meantime, however, another millimetre array is making its debut in the Southern Hemisphere: the Australia Telescope Compact Array (ATCA), being upgraded to operate at wavelengths as short as 3mm, is poised to open a new astronomical window into the southern skies.  In anticipation of the completion of the ATCA upgrade, a workshop on millimetre-wave science was convened at the University of Melbourne on 29-30 November 2001.  Participants from the Australian astronomical community, as well as guests from overseas, came to discuss and learn about areas of science that could be addressed with the new ATCA.  As is inevitable with a small meeting, many areas could not be properly represented---including, but not limited to, cosmology and VLBI---so this summary should be considered only a sampling of what might be possible.  We also apologise in advance if, in trying to summarise the meeting in a cohesive way, we have been rather selective in our recollections.  Slides from the talks are available for download from the ATNF website\\footnote{http://www.atnf.csiro.au/whats\\_on/workshops/mm\\_science2001}. ",
        "conclusions": ""
    },
    "0207/astro-ph0207478_arXiv.txt": {
        "abstract": "{We have obtained ground-based {\\em I}, {\\em J} and {\\em K} band images of the spiral galaxy, Messier 74 (NGC 628). This galaxy has been shown to possess a circumnuclear ring of star formation from both near-infrared spectroscopy of CO absorption and sub-millimetre imaging of CO emission. Circumnuclear rings of star formation are believed to exist only as a result of a bar potential. In this paper we show evidence for a weak oval distortion in the centre of M74. We use the results of Combes \\& Gerin (1985) to suggest that this weak oval potential is responsible for the circumnuclear ring of star formation observed in M74. ",
        "introduction": "Circumnuclear rings of star formation have been shown to occur in barred spiral galaxies since the early 1980s (e.g. Benedict 1980) and they have been studied in great detail in many galaxies since this pioneering work (e.g. Knapen 1996; Knapen et al. 1999). They are thought to be a result of a funneling of material to the central regions of the galaxies by a bar potential. Indeed, hydrodynamical simulations of galaxies have shown that gaseous material is shocked at the leading edge of a bar and diverted towards the centre of the galaxy (Roberts, Huntley \\& van Albada 1979). The material can then accumulate at the inner Lindblad resonance (ILR) until it reaches a critical density at which star formation can be induced.  The spiral galaxy, Messier 74 (NGC 628), is classed as a non-barred spiral galaxy (its Hubble classification is SAc - de Vaucouleurs et al. 1991). However, a circumnuclear ring of star formation does exist in the central regions of M74. This has been observed in $^{12}$CO $J=1-0$ sub-mm imaging (Wakker \\& Adler 1995) and 2.3 $\\mu$m CO absorption spectroscopy (James \\& Seigar 1999). Is it therefore possible that a bar--like structure exists in the centre of M74, but is shrouded in dust? One way to answer this question is to observe M74 at near-infrared wavelengths. The first demonstration that bars are more common in the near-infrared was performed by Hackwell \\& Schweizer(1983). Since then, this method has proved successful for uncovering bars in spiral galaxies in many cases (e.g. Seigar \\& James 1998; Eskridge et al. 2000). In this paper we present {\\em I}, {\\em J} and {\\em K} band images of M74, the longest wavelength images available, in order to uncover a bar in its centre.  This paper is arranged as follows: section 2 describes the observations; section 3 is a discussion of the results presented in this paper; in section 4 we summarise our main results. ",
        "conclusions": "We have presented in this paper, $I$, $J$ and $K$ band images of the spiral galaxy, M74. These images highlight the presence of a weak oval distortion in the central regions of the galaxy. We believe that this oval distortion is responsible for the circumnuclear ring of star formation observed in molecular CO emission (Wakker \\& Adler) and 2.3 $\\mu$m CO absorption (James \\& Seigar 1999). We have also argued that even weak oval distortions can be responsible for such circumnuclear star formation, using the simulations presented by Combes \\& Gerin (1985)."
    },
    "0207/astro-ph0207222_arXiv.txt": {
        "abstract": "A decade after the publication of the Hipparcos Catalogue, the Space Interferometry Mission (SIM) will be capable of making selected high-precision astrometric measurements about three orders of magnitude more accurate than the Hipparcos survey.  We present results from a detailed set of end-to-end numerical simulations of SIM narrow-angle astrometric measurements and data analysis to illustrate the enormous potential that SIM has for the discovery and characterization of planets outside the Solar System. Utilizing a template observing scenario, we quantify SIM sensitivity to single planets orbiting single normal nearby stars as function of measurement errors and properties of the planet: SIM will detect over 95\\% of the planets with periods between a few days and the 5-year nominal mission lifetime that produce astrometric signatures $\\sim 2.2$ times larger than the single-measurement accuracy. We provide accuracy estimates of full-orbit reconstruction and planet mass determination: at twice the discovery limit, orbital elements will be determined with a typical accuracy of 20-30\\%; the astrometric signature must be $\\sim 10$ and $\\sim 15$ times the minimum signal required for detection to derive mass and inclination angle estimates accurate to 10\\%. We quantify the impact of different observing strategies on the boundaries for secure detection and accurate orbit estimation: the results scale with the square root of both the number of observations and the number of reference stars. We investigate SIM discovery space, to gauge the instrument ability in detecting very low-mass planets: around the nearest stars, SIM will find planets as small as Earth, if they are present. Some of these might be orbiting inside the parent star's Habitable Zone.  Extra-solar planets figure prominently among SIM scientific goals: our results reaffirm the importance of high-precision astrometric measurements as a unique complement to spectroscopic surveys based on radial velocity. For example, establishing the existence of rocky, perhaps habitable planets would constitute both a fundamental test of theoretical models, and progress towards the understanding of formation and evolution processes of planetary systems. Such discoveries would also provide the Terrestrial Planet Finder (TPF) with prime targets to investigate with direct spectroscopy in terms of the potential for life. ",
        "introduction": "Six years ago, research in planetary science was essentially synonymous with studies of our Solar System alone. Only four years later, thanks to extensive precision radial velocity surveys of the solar neighborhood, Butler et al.~\\cite{butler00} could list 50 nearby Main-Sequence stars orbited by at least one planet candidate with projected masses\\footnote{Radial velocity techniques cannot determine the viewing geometry of the orbit, and consequently only lower limits to the companion mass can be inferred} below the so-called deuterium burning threshold ($M\\sin i < 13 M_\\mathrm{J}$, where $M_\\mathrm{J}$ is the mass of Jupiter), as discussed by Oppenheimer et al.~\\cite{oppen00}. The number of planets continues to grow; as of July 2002, 26 more candidate planets have been identified, bringing the number of stars harboring planetary-mass companions to 88. Recently, one of these candidate extra-solar planets was confirmed to be a Jupiter-mass object in a few-days period orbit (a {\\it Hot Jupiter}) via transit observations of the star HD 209458~\\citep{henry00,charbon00}, and the presence of sodium in its atmosphere detected~\\citep{charbon02}. Furthermore, radial velocity measurements have also proven the existence of candidate planetary {\\it systems}~\\citep{butler99,marcy,udry,fischer02}, and of systems composed of a planet and a brown dwarf candidate~\\citep{udry,marcy2,els}. Except for the case of HD 209458, all low-mass companions to solar-type stars having $M\\sin i < 13 M_\\mathrm{J}$ have been classified by some as extra-solar planets solely on the basis of their small projected masses, and thus, under the reasonable assumption of orbital planes randomly oriented in space, small true masses. In fact, the true nature of such objects is still matter of ongoing debates among the scientific community. For example, the unexpected orbital configurations of the majority of the planet candidates, such as companions having $M\\sin i\\geq M_\\mathrm{J}$ and orbital periods of a few days~\\citep{mayor95,butler97} or large eccentric orbits~\\citep{mazeh96,coch97} have raised crucial questions about their origin. In response to these challenging discoveries, new theoretical models have been proposed, which invoke diverse mechanisms like orbital migration~\\citep{murray98,trilling98,delpopolo02} or {\\it in situ} formation~\\citep{ward97,wuch97,bode00}. Statistical analyses have also been carried out~\\citep{heac99,step00,mayor00,step01,mazeh01}, which, highlighting the striking similarity between the distributions of eccentricities and periods of the two populations, suggest alternative scenarios implying common formation processes for planet and brown dwarf candidates, and for stellar binaries. On the other hand, recent attempts to determine the actual mass distribution of low-mass companions to nearby stars~\\citep{mazeh01,jorissen01,zucker01b,halbwachs01} have confirmed the indications obtained by early studies~\\citep{basri97,mayor98a,mayor98b,marcy98} which pointed out remarkable differences in the mass distribution of planet candidates and low-mass stellar secondaries. In particular, the two populations appear to be separated by a gap in mass of roughly an order of magnitude in the range 10-100 $M_\\mathrm{J}$. This is the so-called ``brown-dwarf desert'' (see for example the early works of Campbell et al.~\\cite{campbell88}, Marcy \\& Benitz~\\cite{Marcy89}, Marcy \\& Butler~\\cite{marcy94}, and more recently Halbwachs et al.~\\cite{halbwachs00}), commonly thought of as supporting the idea that the two populations are actually distinct, and therefore suggesting that planet candidates are indeed planets. On the basis of joint analyses of Hipparcos Intermediate Astrometric Data and ground-based astrometric observations, Gatewood et al.~\\cite{gatewood01} and Han et al.~\\cite{han01} cast doubts on the actual planetary nature of the low-mass objects detected by radial velocity surveys, disputing the hypothesis of randomness of the orbital planes. Recently, Pourbaix~\\cite{pourbaix01a}, Pourbaix \\& Arenou~\\cite{pourbaix01b}, and Zucker \\& Mazeh~\\cite{zucker01a} have questioned the statistical significance and robustness of their method and results, arguing that the present milli-arcsecond precision of the most accurate astrometric measurements today available is insufficient to derive sensible conclusions on the exact nature of these objects. Furthermore, the Main-Sequence stars harboring the planet candidates have been shown to have higher metallicity than the average of field stars with the same mass in the solar neighborhood~\\citep{laughlin00,gonzalez01,santos01}, and these findings may support the evidence for significant correlation between high stellar metallicity and the presence of orbiting giant planets, under the assumption that core accretion is the primary planet formation mechanism. However, if disk instability is the preferred mechanism for forming extrasolar giant planets, then the metallicity dependence may turn out to be an artifact due to observational selection effects or stellar pollution by ingestion of planetary material, and even low-metallicity stars should harbor giant planets~\\citep{boss02}. As it can be easily understood, such a plethora of diverse interpretations clearly indicates how our present understanding of the origin of planetary systems is {\\it de facto} still limited, and significant contributions in terms of data obtained via means other than Doppler shift measurements are essential in order to be able to discriminate between biased theoretical and observational models.  Radial velocity surveys, accurate to 3-5 m/s~\\citep{butler96}, have been so far a unique tool for planet discovery. However, we anticipate that high-precision astrometry, both from ground \\citep{mariotti98,booth99,colavita99} and in space \\citep{danner99,roser99,gilm00}, will be among the preferred means for helping fill regions of the parameter space Doppler techniques cannot reach. Astrometry has a significant advantage over the radial velocity technique because it measures two rather than one projection of an orbit and thus describes the full three-dimensional geometry. Astrometry, removing the degeneracy on the inclination angle, provides unambiguous mass estimates and directly determines coplanarity for systems of planets. Astrometric techniques can be used to search for planets around young and bright stars (earlier than F), and late M dwarfs. These objects cannot be searched by radial velocity, either because of their spectral properties (absence of relevant spectral lines) or because of intrinsic instability of the stellar atmospheres (active chromospheres, spots, significant rotation). Furthermore, astrometric sensitivity increases for planets with longer periods, thus complementing radial velocity searches, which favor short-period planets. Finally, radial velocity detection limits are currently of about a Saturn mass within 1 AU. As we will see, astrometry with SIM's exquisite sensitivity pushes detection two orders of magnitude lower, down to Earth masses.  SIM (Space Interferometry Mission) is under development as NASA's first space-based optical interferometer devoted to micro-arcsecond ($\\mu$as) astrometry~\\citep{danner99}. It represents a First Generation Mission within NASA's {\\it Origins} Program (\\url{http://origins.jpl.nasa.gov/}), which has the long-term goal of direct imaging of Earth-like planets around nearby solar-type stars. The instrument is scheduled for launch by mid-2009, with a nominal mission lifetime of 5 years. SIM will perform pointed observations, unlike astrometric missions such as Hipparcos (\\url{http://astro.estec.esa.nl/Hipparcos/}), DIVA~\\citep{roser99}, or the recently approved ESA Cornerstone Mission GAIA~\\citep{perryman01}, which are designed to survey the sky using a well-defined scanning law, in order to build global astrometric catalogues. On one hand, this will limit the total mission throughput. A few tens of thousands objects will be observed, compared to the 120\\,000 stars surveyed by Hipparcos or to the $10^7-10^9$ objects which are expected to be charted by DIVA and GAIA, respectively. On the other hand, SIM's pointed observations can achieve unprecedented astrometric accuracy that will bring new light in the exploration of our galactic neighborhood.  Detection and measurement of planets will be carried out primarily with SIM operated in narrow-angle astrometric mode. The instrument is expected to achieve a narrow-angle {\\it single measurement} accuracy of $\\sim 1$ $\\mu$as in 1 hr integration time on bright targets ($V\\leq 11$), which corresponds to the amplitude of the gravitational perturbation induced on a solar-mass star by an Earth-mass planet on a 1 AU orbit, as seen from 3 pc. SIM local astrometry is therefore uniquely suited for detection and measurement of planets with masses as small as a few Earth masses in the vicinity of the Solar System.  This is the first of two papers which will connect and relate the basic SIM capabilities to the properties of extra-solar planetary system. We have built a detailed software suite to $a)$ simulate sample narrow-angle SIM observing campaigns of stars with planets, and $b)$ analyze the simulated datasets resulting from such observations. The purpose of this first paper is to show how these tools can be used to evaluate the detectability of single planets around single stars and their measurability in terms of mass and orbital characteristics, as a function of both SIM mission parameters and properties of the planet. In the second paper we will address the issues of the detectability and measurability of systems of planets with SIM, as well as extensive analyses of the instrument capability to determine the coplanarity (or non coplanarity) for a variety of orbital arrangements in multiple-planet systems.  The first paper is organized as follows. In the second Section we briefly describe SIM narrow-angle astrometric mode. In the third Section we present a description of the software for the simulation of SIM narrow-angle observations. Details on detection and orbit determination methods are given in the fourth Section. The most significant results obtained so far are presented in the fifth Section, followed by summary and conclusions. ",
        "conclusions": "Since the establishment of the existence of the first extra-solar planet orbiting a solar-type star~\\citep{mayor95}, the approach to sciences of stars and planets has dramatically changed. Now, answers are sought to more advanced questions about the formation and evolution of planetary systems and the existence of rocky, perhaps habitable planets.  Precision astrometry constitutes a fundamental complement to other search techniques. Today, monolithic telescopes and optical interferometers are being built or designed, which will provide accurate astrometric measurements, both from ground~\\citep{mariotti98,booth99,colavita99} and in space~\\citep{danner99,roser99,perryman01}.  In this paper we have used extensive end-to-end numerical simulations of narrow-angle astrometric measurements with the Space Interferometry Mission and the subsequent statistical analysis of the simulated dataset in order to quantify the potential of SIM for the discovery and characterization of single planets around single stars in the vicinity of the solar system. Utilizing a simplified, but realistic, error model for SIM operated in narrow-angle mode, and adopting a reasonable, flexible template observing scenario (Sections~\\ref{scenario}), we have: $a)$ defined the boundaries for secure planet detection and accurate determination of orbital elements and masses, as function of the basic SIM capabilities and properties of the observed systems (Sections~\\ref{detect},~\\ref{errors}, and~\\ref{known}), $b)$ evaluated the impact of different observing strategies on the boundaries for detection and orbit reconstruction (Sections~\\ref{refer} and~\\ref{timing}); $c)$ adopting template observing strategies for both bright ($V\\leq 11$) and faint targets (Section~\\ref{faint}), illustrated SIM discovery potential in terms of its ability to detect terrestrial planets around a sample of the closest stars (Section~\\ref{space}). Our main results can be summarized as follows.  \\begin{itemize} \\item[(1)] secure detection (at the 95\\% confidence level) will be possible for planets producing an astrometric signature $\\alpha_{\\mathrm{min}}\\sim 2.2$ times larger than the Standard Visit accuracy $\\sigma_\\mathrm{d}$, for periods shorter than 5 years, the nominal mission lifetime; \\item [(2)] in the same period range, the mass of the planet and the full set of orbital elements will be determined with a typical accuracy of 20-30\\% for objects producing a signal $\\sim 2\\,\\alpha_{\\mathrm{min}}$; for mass and inclination measurements accurate to 10\\%, the required signal is $\\sim 10\\,\\alpha_{\\mathrm{min}}$ and $\\sim 15\\,\\alpha_{\\mathrm{min}}$, respectively; analyzing how the set of presently known extra-solar planets would fall within the boundaries for reliable detection and accurate mass and orbit determination, we find that about 75\\% will be detected and 50\\% will have orbital elements and masses measured to 10\\%, or better; \\item[(3)] the detection threshold scales similarly with the number of observations ($\\sqrt{N_o}$) and reference stars ($\\sqrt{N_r}$); random uniform and geometric distributions of the observations are preferred for achieving the best detection sensitivity and more accurate estimates of orbital parameters and masses in the period range between $\\sim 1$ month and 5 yr, and $\\leq 1$ month, respectively; \\item [(4)] due to the very small astrometric signature induced on the parent star, reliable detection of Earth-class planets in the Habitable Zone of the closest solar-type stars will be possible, but demanding in terms of number of full observations per target and measurement precision; instead, around the nearest M dwarfs, more relaxed constraints on the number of observations and single-measurement errors would still ensure detection of planets as small as Earth. \\end{itemize}  Our findings indicate how SIM, with its unprecedented astrometric precision, will be a valuable tool for discovering planets around stars other than the Sun. Among the new generation of instruments designed to study extra-solar planets, SIM will be able to provide unique insights towards the understanding of planetary systems in their generality and investigating the habitability of other worlds than Earth. Today two factors hamper the transition from the present cataloguing phase to the more fundamental classification phase, where, for example, mass might be operationally used as one of the genesis indicators which would help discriminate between planets and brown dwarfs: $a)$ mass uncertainty for the radial-velocity discoveries (due to inclination angle ambiguity), and $b)$ incompleteness in the mass range corresponding to solar system planets (due to inadequate sensitivity). By determining the true rather than the projected orbit of the planet (as with radial-velocity techniques), SIM measurements will remove the inclination angle degeneracy and associated companion-mass uncertainty for the presently existing planets, as well as for those the instrument will discover directly. Furthermore, by ruling out the presence of Earth-mass planets around the nearest stars, SIM will be capable of addressing for the first time the role of rocky cores in the complex scenarios of planetary formation and evolution, and start to investigate their potential habitability. In fact, SIM astrometry will be important in investigating the Habitable Zones of stars with known planets in wide orbits: those systems in which the Habitable Zone and the zone in which planet formation has not been disrupted by the presence of the known giant planet overlap~\\citep{wetherill96} would immediately become high-priority targets for SIM narrow-angle observations, to search for terrestrial planets and find evidence of the existence of planetary systems resembling our own.  Another crucial area in which SIM measurements might have a significant impact is the study of multiple-planet systems: the remarkable pattern of low-eccentricity orbits and coplanar structure of the solar system are commonly thought to be fossil evidence of the planets having accumulated in a dissipative protoplanetary disk~\\citep{lissa93,pollack96}. The wide variety of planetary masses and orbits found by radial velocity techniques have called into question the generality of such ideas, suggesting that significant orbital evolution may be needed to explain the high-eccentricity orbits~\\citep{arty92,weiden96,lin97,mazeh97} and the {\\it Hot Jupiters} at very small orbital radii~\\citep{lin96,murray98,lin00}. By answering the seemingly simple question of whether multiple-planet orbits are coplanar, SIM might confirm that some other planetary systems are similar to our own and similarly indicative of origin in a quiescent, flattened disk. Or, SIM measurements might provide evidence that other systems are truly different, with large relative orbital inclinations, which could point to either an early, chaotic phase of orbital evolution or formation by another mechanism such as disk instability~\\citep{kuiper51,cameron78,boss97,boss00,bossetal02}. The simulation of SIM observations of extra-solar multiple-planet systems, the quantification of the instrument capability in discovering and measuring systems of planets, as well as its ability in determining coplanarity of multiple-planet orbits, will constitute the core of the results presented in paper II."
    },
    "0207/astro-ph0207236_arXiv.txt": {
        "abstract": "From the peak of a gravitational microlensing high-magnification event in the A component of QSO 2237+0305, which was accurately monitored by the GLITP collaboration, we derived new information on the nature and size of the optical $V$-band and $R$-band sources in the far quasar. If the microlensing peak is caused by a microcaustic crossing, we firstly obtained that the standard accretion disk is a scenario more reliable/feasible than other usual axially symmetric models. Moreover, the standard scenario fits both the $V$-band and $R$-band observations with reduced chi-square values very close to one. Taking into account all these results, a standard accretion disk around a supermassive black hole is a good candidate to be the optical continuum main source in QSO 2237+0305. Secondly, using the standard source model and a robust upper limit on the transverse galactic velocity, we inferred that 90 per cent of the $V$-band and $R$-band luminosities are emitted from a region with radial size less than 1.2 10$^{-2}$ pc (= 3.7 10$^{16}$ cm, at 2$\\sigma$ confidence level). ",
        "introduction": "In the optical continuum, QSO 2237+0305 is a gravitational mirage that consists of four compact components (A-D) round the nucleus of the deflector (lens galaxy). The light bundles corresponding to the components are passing through the bulge of the lens galaxy, and thus, if the galactic mass at the QSO image positions is mainly in stars, the optical depths to microlensing are as high as $\\sim 0.5$ (e.g., Schmidt, Webster \\& Lewis 1998). In this scenario with large normalized surface mass densities, given a component, microlensing violent events will result from the source either crossing a microcaustic, passing close to a microcusp, or traveling through a network of microcaustics (e.g., Schneider, Ehlers \\& Falco 1992). On the other hand, a microlensing violent episode in the light curve of a component can be easily distinguished from a intrinsic variation, since the intrinsic variability must be observed in all four components of the system with extremely short time delays (e.g., Wambsganss \\& Paczy{\\' n}ski 1994; Chae, Turnshek \\& Khersonsky 1998; Schmidt, Webster \\& Lewis 1998). However, when the four brightness records of the QSO images have different non-flat shapes, a direct separation between the true microlensing fluctuations and the possible intrinsic variation cannot be achieved. In the case of four incoherent and non-flat observational trends, to find the true microlensing behaviours we must do some hypothesis on the intrinsic variability.  Irwin et al. (1989) discovered microlensing variability in the quadruple system QSO 2237+0305, and that first evidence was confirmed by other observers (Corrigan et al. 1991; {\\O}stensen et al. 1996; Wo{\\' z}niak et al. 2000a,b; Schmidt et al. 2001; Alcalde et al. 2002). In very recent years, two gravitational microlensing high-magnification events (HMEs) were clearly detected by the OGLE collaboration (Wo{\\' z}niak et al. 2000b) and corroborated by the GLITP monitoring (Alcalde et al. 2002). As each individual HME is directly related to the intrinsic surface brightness of the source, the two HMEs in the $V$ band reported by the OGLE team were used to obtain two measurements of both the optical continuum ($V$-band) source size and the rate of brightness decline (Shalyapin 2001). Shalyapin (2001) compared the observational data included in each HME with the time evolution expected from an axially symmetric source crossing a single straight fold caustic. To describe the brightness distribution of the $V$-band source, he used a power-law model: $I_V(r) = 2^{p_V} I_V (1 + r^2/R_V^2)^{-p_V}$, which is determined by a typical radius $R_V$, a typical intensity $I_V = I_V(R_V)$, and a power-law index $p_V$. The expected microlensing light curves depend on five free parameters, and the most relevant ones are $\\Delta t = R_V/V_{\\perp}$ and $p_V$. We note that a direct measurement of the typical radius $R_V$ is not possible, however, using some upper limit on the quasar velocity perpendicular to the caustic line ($V_{\\perp}$), it can be obtained a very interesting constraint on the source size as measured by means of the typical radius of the 2D brightness distribution. A crossing time of $\\Delta t \\approx$ 90 days is inferred from the HME observed in the light curve of the component A, whereas a shorter crossing time of about 30 days is consistent with the HME corresponding to the image C. However, the whole HME of image A (from day 1200 to day 1800) seems to be caused by a complex magnification, which is different to the single straight fold caustic magnification law. In other words, when the main portion of the source is far away from the fold caustic of interest (before day 1400 and after day 1600), there is evidence for a {\\it rare} behaviour, and so, the fit to the whole microlensing event in the brightness record of the component A could give biased estimates of the parameters. Taking into account this perspective, we only take the results based on the HME corresponding to the image C as non-biased parameter estimates. The value of $\\Delta t \\approx$ 30 days together with the velocity constraint by Wyithe, Webster \\& Turner (1999) give an upper limit on the $V$-band typical radius of $R_V \\leq$ 3 10$^{-4}$ pc. A similar conclusion was obtained by Yonehara (2001), who analyzed the same microlensing event but using a different picture and a typical microlens mass of $\\approx$ 0.1 $M_{\\odot}$ (Wyithe, Webster \\& Turner 2000). On the other hand, Shalyapin (2001) found that the value of the power-law index is close to the validity limit of the model ($p_V \\sim$  1) and the best-fit reduced $\\chi^{2}$ is significantly greater than 1.  Apart from the very recent papers by Shalyapin (2001) and Yonehara (2001), other previous works also discussed the size of the optical continuum source (e.g., Wambsganss, Paczynski \\& Schneider 1990; Webster et al. 1991; Wyithe et al. 2000). These first studies are based on a poorly sampled HME which was observed in Q2237+0305A during late 1988. While the continuum mostly arises from a compact source, the line emission comes from a much larger region. From two-dimensional spectroscopy of QSO 2237+0305, Mediavilla et al. (1998) found an arc of extended C {\\sc III}]$\\lambda$1909 emission that connects the components A, D, and B. The observed arc is consistent with a source radius larger than 100 pc.  The GLITP (Gravitational Lenses International Time Project) collaboration has monitored QSO 2237+0305 in the period ranging from 1999 October to 2000 February (Alcalde et al. 2002). The GLITP/PSFphotII photometry in the $V$ and $R$ bands (see Fig. 1 in Alcalde et al. 2002), showed a peak in the flux of the component A and a relatively important gradient in the flux of the component C. These features are related to the two HMEs that were discovered by the OGLE team. The GLITP light curve for image A traced the peak of the corresponding HME, i.e., the maximum and its surroundings, with an unprecedented quality. For example, the OGLE collaboration sampled the $V$-band peak at 19 dates, whereas the GLITP record included measurements of the $V$-flux at 52 dates. Moreover, the rest of global behaviours (components B-D) were accurately drawn from the GLITP photometry (in this paper, we will use the PSFphotII variant). The $VR$ light curves of the four components A-D have been also analyzed from a phenomenological point of view, and the global flat shape for the light curve of Q2237+0305D suggested that the microlensing signal in D and the intrinsic signal are both globally stationary. In consequence of this result, it seems that the global variabilities in A-C are unambiguously caused by microlensing.  As mentioned here above, the whole HME of Q2237+0305A seems to be originated by a complex magnification law. However, in principle, the peak of the HME (just when the main portion of the source crossed a microcaustic) could be a structure mainly caused by a single straight fold caustic, i.e., the curvature and other possible close microcaustics do not significantly perturb the simple magnification law. We adopt this last point of view, and take the GLITP light curve for image A, including only data points very close to the maximum of the HME, to be fitted to the microlensing curves resulting from sources crossing a single straight fold caustic. We remark that the high asymmetry of the peak as well as probabilistic arguments are two strong reasons against an interpretation of the microlensing peak based on a source passing close to a single cusp caustic.  In Section 2 we present the expected microlensing light curves when axially symmetric sources cross a single straight fold caustic, and the fitting procedure. We use a set of axisymmetric sources: brightness distributions enhanced at the centre of the source (standard physical profile, Gaussian profile, and $p$ = 3/2, 5/2 power-law profiles) and the uniform brightness distribution (e.g., Shakura \\& Sunyaev 1973; Schneider \\& Weiss 1987; Shalyapin 2001). Section 3 is devoted to the parameter estimation from the comparison between the GLITP microlensing peak in the component A and the expected time evolutions for the different source models. A discussion on the V-band and R-band source sizes, the source size ratio ($R_V/R_R$), and the reliability of the source models is also included in Section 3. In this paper, to obtain information on the dimension of the $V$-band and $R$-band sources, we will use the measurements of the transverse galactic velocity reported by Wyithe, Webster \\& Turner (1999). Finally, in Section 4 we summarize our results and conclusions. ",
        "conclusions": "We have analyzed the peak of a microlensing high-magnification event which was accurately monitored by the GLITP collaboration (Alcalde et al. 2002). The prominent event has occurred in image A of the quadruple system QSO 2237+0305 (Wo{\\' z}niak et al. 2000b), and the GLITP team observed its peak in two optical bands ($V$ and $R$). As both $V$-band and $R$-band light curves are characterized by high flux and time resolutions, we attempted to interpret the observational trends in terms of some theoretical microlensing light curves. In principle, the observed microlensing peak could mainly result from the source crossing a microcaustic or passing close to a microcusp. However, taking into account probabilistic arguments and the important asymmetry around day 1500 (see Fig. 3), we have discarded the hypothesis of a source in the vicinity of a single cusp caustic. Thus we only considered the most simple alternative picture: a source crossing a single straight fold caustic. Different axially symmetric source models were used, and consequently, several theoretical microlensing curves were compared with the observed ones. To discuss the reliability/feasibility of different intrinsic intensity profiles, all source models were chosen to cause theoretical microlensing curves with the same number of free parameters. These parameters are two characteristic fluxes, the time of caustic crossing by the source centre, and the typical crossing time.  Our main results and conclusions are: \\begin{enumerate} \\item From the fits and the upper limits on the transverse galactic velocity claimed by Wyithe, Webster \\& Turner (1999), we inferred constraints on the $V$-band and $R$-band source sizes. For a uniform source model, the $V$-band and $R$-band radii should be less than 6.3 10$^{-4}$ pc = 1.9 10$^{15}$ cm, at 1$\\sigma$ confidence level. For a Gaussian disk, the typical $V$-band radius is $<$ 7.8 10$^{-4}$ pc (= 2.4 10$^{15}$ cm, at 1$\\sigma$ CL), while the typical $R$-band radius can be a little larger ($<$ 1.6 10$^{-3}$ pc = 4.9 10$^{15}$ cm, at 1$\\sigma$ CL). Taking the total size of the $R$-band source as the source diameter for a uniform disk, or the full-width at one-tenth maximum for a Gaussian profile, we found that the 1$\\sigma$ upper limits on the $R$-band source size are 3.9 10$^{15}$ cm (top-hat) and 1.5 10$^{16}$ cm (Gaussian). These bounds agree well with the results by Wyithe et al. (2000). On the other hand, for the standard accretion disk, we obtained that 90\\% of the $V$-band and $R$-band luminosities are radiated within a radius of 8 10$^{-3}$ pc (= 2.5 10$^{16}$ cm, at 1$\\sigma$ CL), 1.2 10$^{-2}$ pc (= 3.7 10$^{16}$ cm, at 2$\\sigma$ CL). As the total luminosity of a standard accretion disk around a 10$^8$ $M_{\\odot}$ black hole will be probably enclosed in a circle with radial size of $\\approx$ 10$^{-2}$ pc, the {\\it standard} results are highly consistent with the current paradigm on the central engine in QSOs. However, we remark that other central masses also agree with the constraints. Once upper limits of $\\sim$ 10$^{15}$ -- 10$^{16}$ cm are known, we may attempt to test the hypothesis of a negligible caustic curvature (at $V_{\\parallel} <$ 7000 km s$^{-1}$, the parallel path length during $\\sim$ 100 days will be also less than 10$^{15}$ -- 10$^{16}$ cm). The typical caustic curvature radius is usually assumed to be the Einstein-ring radius on the source plane, i.e., $R_{CC} \\sim R_E \\sim 10^{17}(m/M_{\\odot})^{1/2}$ cm, where $m$ is the microlens mass ($\\Omega$ = 1, $H$ = 60 -- 70 km s$^{-1}$ Mpc$^{-1}$). Therefore, for $m \\sim 1 M_{\\odot}$, the ratios between the optical radii and the typical caustic curvature radius will be smaller than 0.01 -- 0.1, and this result supports the straight fold caustic approximation. For a small microlens mass ($m \\sim 0.1 M_{\\odot}$), we however have a smaller $R_E$ and cannot confirm the weakness of the curvature effects. \\item We also studied the source size ratio, i.e., the ratio between the $V$-band radius and the $R$-band radius. The source size ratio values are independent of the transverse quasar velocity. At 1$\\sigma$ CL, only the Gaussian source model led to a $V$-band source being inside the $R$-band one. For the standard source, one hopes for a ratio of about 0.8, but this expected value could not be confirmed from our microlensing experiment. The 1$\\sigma$ confidence interval (0.53 -- 1.26) is in agreement with the two possible situations: the $V$-band source being inside the $R$-band source, and both sources having a similar size. \\item An important issue is the reliability/feasibility of several source models tested by us. With respect to this point, we deduced very interesting conclusions. The usual top-hat and Gaussian profiles are not favored from the data in the $V$ and $R$ bands. The results are better when power-law profiles are assumed, particularly for the model with a power index of 1.5, which more closely resembles the one-dimensional intensity profile of the exact standard accretion disk. For the standard source model, the reduced chi-square values are very close to one. So, an accretion disk around a supermassive black hole seems a good candidate to be the optical continuum main source in QSO 2237+0305 (see Fig. 3), and a measurement of the black hole mass and the mass accretion rate could be made in that system (e.g., Yonehara et al. 1998). This last topic will be discussed in a separate paper. We also note that a hybrid scenario in which the light in the $V$ band is emitted from an accretion disk and the $R$-band light comes from the accretion disk and another extended region, is not in contradiction with the observed microlensing peaks. This result totally agrees with the main conclusion of Jaroszy\\'nski, Wambsganss \\& Paczy\\'nski (1992), who previously tested the hybrid model from old observational data. The assumption of a complementary $R$-band extended source only modifies the expression of the $R$-band background flux $F_{0R}$ in Eqs. (11) and (13), and thus, the interpretation of its best-value and uncertainties. Exclusively using GLITP data, one can try to analyze the possible existence of a light excess in the $R$ band. To do this task, it is needful an accurate calibration of the flux in both optical bands and small uncertainties in the measurements of the background fluxes. A detailed study of this topic is however out of the scope of the paper, which is focused on the optical continuum main (compact) source. On the other hand, in order to fit the observed spectrum of the quasar, Rauch \\& Blandford (1991) did not consider additional $R$ light coming from a large region, but adopted a non-classical accretion disk model. Unfortunately, the disk model by Rauch \\& Blandford (1991) cannot account for the old microlensing variations in QSO 2237+0305. \\end{enumerate}  We remark that the OGLE collaboration also monitored the $V$-band event in Q2237+0305A (http://www.astro.princeton.edu/$\\sim$ogle/ogle2/huchra.html). The GLITP $V$-band photometry traced the peak of the microlensing event, whereas the OGLE $V$-band dataset described the behaviour of the whole fluctuation. In comparison with the OGLE observational procedure, the GLITP observations are of higher quality because they were obtained with a larger telescope, using a detector with better resolution, and on nights with better seeing. Therefore, as due to the quality of the observations and the excellent sampling rate, the GLITP $V$-band peak is probably the best tracer of the underlying signal around the maximum of the whole flux variation. On the other hand, the whole event seems to be caused by a complex magnification, which may include ingredients such as the curvature of the fold caustic (e.g., Fluke \\& Webster 1999), the presence of another caustic, and a non-constant background magnification (e.g., Gaudi \\& Peters 2002). The simple fit to the whole event may thus give a wrong estimation of the parameters. In any case, we chose two OGLE observation periods to infer {\\it standard} solutions and compare them with the {\\it standard} fit presented in Table 1. The first period included only the end of 1999, from day 1450.6 to day 1529.5. The dataset is called OGLE99, and it corresponds to the GLITP monitoring period. The second dataset (OGLE99-00) covers the 1999-2000 seasons, more exactly from day 1289.9 to day 1766.7. In Table 4 we can see all at once the fits from the GLITP, OGLE99, and OGLE99-00 light curves. As $\\hat{\\chi}^2 (min) >$ 1 for the OGLE best solutions, in the time parameter estimation from the OGLE datasets, we considered the 1$\\sigma$ bounds associated with $\\Delta \\chi^2 = \\hat{\\chi}^2 (min)$ rather than $\\Delta \\chi^2$ = 1 (e.g., Grogin \\& Narayan 1996). The GLITP and OGLE99 fits are not consistent each other, but the amplitudes and the crossing time from the OGLE99-00 brightness record are close to the values of $F_0$, $F_C$, and $\\Delta t$ from the GLITP dataset.  Finally, we must emphasize that an important progress can be made from accurate and detailed data of a microlensing peak. One can discuss on the reliability/feasibility of different source models leading to the same number of free parameters, and thus, to discard some of them and to find the models that agree with the observations. Given a good model, which is consistent with the data, it is possible to obtain a robust upper limit on the size of the optical source. If, for example, the good model is an accretion disk around a massive black hole, then one can try to work out a technique to measure the central mass and the accretion rate. Moreover, as the accurate and well-sampled microlensing peaks seem to be inconsistent with some intensity profiles, it could be very interesting to apply the deconvolution method to these peaks and to derive the best brightness distribution in a more direct way (Grieger, Kayser \\& Schramm 1991; Agol \\& Krolik 1999; Mineshige \\& Yonehara 1999). Although we had not success in the accurate and robust indirect estimation of the ratio between the $V$-band radius and the $R$-band radius (the parameter $q$), new multiband monitoring campaigns could also lead to relevant measurements of source size ratios. With regard to this last issue, we note that a direct estimate of $q$ from the cross-correlation of our $V$-band and $R$-band light curves is now in progress, but the expectations are not very promising."
    },
    "0207/gr-qc0207047_arXiv.txt": {
        "abstract": "#1{\\vskip 7mm \\begin{center}{\\large Abstract}\\par \\smallskip \\begin{minipage}[c]{12cm} \\small #1 \\end{minipage} \\end{center} } \\def\\title#1{\\begin{center}{\\Large\\bf #1}\\end{center}} \\def\\author#1{\\vskip 5mm \\begin{center}{#1}\\end{center}} \\def\\address#1{\\begin{center}{\\it #1}\\end{center}} \\def\\beq{\\begin{equation}} \\def\\eeq{\\end{equation}} \\def\\C{{\\cal C}} \\def\\B{{\\cal B}} \\def\\R{{\\cal R}} \\def\\d{{\\delta}} \\def\\E{{\\cal E}} \\def\\H{{\\cal E}} \\def\\A{{\\cal A}}  \\makeatletter \\@ifundefined{lesssim}{\\def\\lesssim{\\mathrel{\\mathpalette\\vereq<}}}{} \\@ifundefined{gtrsim}{\\def\\gtrsim{\\mathrel{\\mathpalette\\vereq>}}}{} \\def\\vereq#1#2{\\lower3pt\\vbox{\\baselineskip1.5pt \\lineskip1.5pt \\ialign{$\\m@th#1\\hfill##\\hfil$\\crcr#2\\crcr\\sim\\crcr}}} \\makeatother  \\begin{document}  \\title{% Gravitation and cosmology in a brane-universe \\smallskip \\\\ } \\author{% David Langlois\\footnote{E-mail:langlois@iap.fr} } \\address{% Institut d'Astrophysique de Paris, \\\\ 98bis Boulevard Arago, 75014 Paris, France }  \\abstract{Recent theoretical developments have generated a strong interest in the ``brane-world''  picture, which assumes that ordinary matter is trapped in a three-dimensional submanifold, usually called brane, embedded in a higher dimensional space. The purpose of this  review is to introduce some  basic results concerning gravity in these models and then to present various aspects of  the cosmology in a brane-universe. } ",
        "introduction": "The idea that our world may contain hidden extra dimensions is rather old since one can trace  this idea, in the modern context of general relativity, back to the beginning of the twentieth century, with the pioneering works of Kaluza and Klein, trying to reinterpret electromagnetism as a geometrical effect from a fifth dimension. The idea of extra-dimensions was revived more recently with the advent of string theory as the most promising avenue for conciling gravity and quantum field theory. In order to get a consistent theory at the quantum level, ten spacetime dimensions are needed in superstring theories (eleven in M-theory), which means that six dimensions must be somehow hidden to four-dimensional observers such as us.  The simplest way to hide  extra dimensions is to assume that they are flat, compact with a  radius  sufficiently small to be unobservable. Consider for example the case of one extra-dimension described by the  coordinate $y$ and compactified via the identification $$ y\\rightarrow y+ 2\\pi R, $$ $R$ being the ``radius'' of the extra-dimension. Any matter field, for example a scalar field, depends on both the ordinary spacetime coordinates $x^\\mu$ and the extra-coordinate $y$. It can be Fourier expanded  along the extra-dimension so that \\beq \\phi(x^\\mu,y)=\\sum_{p=-\\infty}^\\infty e^{ipy/R} \\phi_p(x^\\mu). \\eeq The corresponding Fourier modes $\\phi_p$ are  designated as Kaluza-Klein modes, and each of them can be seen as a four-dimensional scalar field satisfying the four-dimensional Klein-Gordon equation with the effective squared mass \\beq M^2_p=m^2+{p^2\\over R^2}. \\eeq A simple way to identify an extra-dimension would thus be to detect the characteristic spectrum of the Kaluza-Klein modes. To do this one needs enough  energy to excite a least the first Kaluza-Klein modes and the non observation of Kaluza-Klein modes can be interpreted as meaning that the size of the extra-dimension is smaller than the inverse of the energy scale probed by the experiment. Present constraints from colliders thus imply $$ R\\lesssim 1 \\ ({\\rm TeV})^{-1}. $$  Let us now turn to gravity. The natural way to extend Einstein gravity to higher dimensions is to start from  the  Einstein-Hilbert action defined in a generalized  spacetime with, say, $n$ extra-dimensions: \\beq S_{grav}=\\int d^4x\\, d^ny \\,  {R\\over 2\\kappa^2}, \\eeq where $R$ is the scalar curvature in the $(4+n)$-dimensional spacetime. Variation of the action including the matter part leads to the generalized Einstein equations  which read \\beq G_{AB}\\equiv R_{AB}-{1\\over 2}R\\  g_{AB}=\\kappa^2 T_{AB}. \\label{einstein} \\eeq This equation has exactly the same form as the familiar one, with the difference that all tensors are now  $(4+n)$-dimensional tensors.  In the static weak field regime (in a flat $(4+n)$-dimensional spacetime) Einstein equations imply  as usual the Poisson equation. However, the solution of Poisson's equation depends on the  number of space dimensions, and the general form of the Newtonian potential is \\beq \\phi_N(r)\\propto {G_{(4+n)}\\over  r^{n+1}}, \\eeq where the generalized Newton's constant $G_{(4+n)}$ is proportional to the gravitational coupling $\\kappa^2$ introduced above in the Einstein equations (\\ref{einstein}).  This means that, a priori, the presence of extra-dimensions implies  that gravity is  modified. The simplest way to recover the familiar gravity law is once more to compactify the extra-dimensions. The resulting gravity is \\begin{itemize} \\item the $(4+n)$ dimensional gravity on scales smaller than the  radius $R$ of compactification (assumed here to be the same in all extra-dimensions for simplicity) \\item the usual $4$-dimensional gravity on scales much larger than $R$, with the Newton's constant given by \\beq G_{(4)}\\sim G_{(4+n)}/R^n, \\eeq or, equivalently, expressed in terms of the Planck mass, \\beq M_{(4)}^2\\sim M_{(4+n)}^{2+n} R^n. \\label{vol_extra} \\eeq \\end{itemize} This result can be understood is the following way. In a compactified space,  the gravitational field  induced by a point mass $m$  can be computed by unwrapping the extra-space and by summing  the contributions of  all the images of the true mass $m$. At small distances with respect to  $R$, the influence of the image masses can be ignored and one gets $(4+n)$ dimensional gravity. By contrast, on scales much larger than $R$, all image masses contribute to the gravitational field and they can be assimilated to a continuous massive ``line'' with constant ``linear'' mass density. Applying then Gauss' law to a cylinder surrounding the massive line yields the usual gravitational force with the above gravitational coupling.  As a consequence, like in particle physics, an upper constraint on the compactification radius can be deduced from the absence of any observed  deviation from ordinary Newton's law. The present experimental  constraints yield (see e.g. \\cite{grav_exp}) \\beq R \\lesssim  0.2 \\, {\\rm mm}. \\eeq  The latest developments in the models with extra-dimensions  come from the realization that the observational constraint on the size of extra-dimensions from gravity experiments is much weaker than that from accelerator experiments. This suggests  the idea to decouple the extra-dimensions ``felt'' by ordinary particles from the extra-dimensions ``felt'' by gravity. Concretely, this can be realized by invoking a mechanism that confines fields of the particle physics Standard Model to a subspace with three spatial dimensions, called three-brane, within a higher dimensional space where gravity  lives.  The purpose of this review, far from being exhaustive, is to present  the basic results concerning gravity  and cosmology in the brane scenarios where the self-gravity of the brane is taken into account, and to illustrate some more advanced aspects like cosmological perturbations or brane collisions. Complementary information can be found in two recent reviews, one by R. Maartens \\cite{maartens} and the other by V. Rubakov \\cite{rubakov}, the latter more  focused on the particle physics aspects. ",
        "conclusions": ""
    },
    "0207/astro-ph0207370_arXiv.txt": {
        "abstract": "Despite extensive study, the mechanisms by which Be star disks acquire high densities and angular momentum while displaying variability on many time scales are still far from clear. In this paper, we discuss how magnetic torquing may help explain disk formation with the observed quasi-Keplerian (as opposed to expanding) velocity structure and their variability.  We focus on the effects of the rapid rotation of Be stars, considering the regime where centrifugal forces provide the dominant radial support of the disk material.  Using a kinematic description of the angular velocity, \\vphi$(r)$, in the disk and a parametric model of an aligned field with a strength $B(r)$ we develop analytic expressions for the disk properties that allow us to estimate the stellar surface field strength necessary to create such a disk for a range of stars on the main-sequence.  The fields required to form a disk are compared with the bounds previously derived from photospheric limiting conditions. The model explains why disks are most common for main-sequence stars at about spectral class B2 V. The earlier type stars with very fast and high density winds would require unacceptably strong surface fields ($> 10^3$ Gauss) to form torqued disks, while the late B stars (with their low mass loss rates) tend to form disks that produce only small fluxes in the dominant Be diagnostics. For stars at B2 V the average surface field required is about 300 Gauss. The predicted disks provide an intrinsic polarization and a flux at \\Halpha\\ comparable to observations. The radial extent of our dense quasi-Keplerian disks is compatible with typical estimates. We also discuss whether the effect on field containment of the time dependent accumulation of matter in the flux tubes/disk can help explain some of the observed variability of Be star disks. ",
        "introduction": "The presence of disks around rotating stars has long been realized from observations, but has never been explained in a satisfactory way. Particularly pronounced are the disks around the emission line Be stars. These can be detected because of the strong double-peaked emission lines of \\Halpha, large IR continuum excesses, and intrinsic polarizations. Because of these interesting and easily detectable properties, Be stars and their disks have been intensively studied for many years (see reviews in Smith, Henrichs, \\& Fabregat (2000) and Yudin (2001)). Nevertheless, the origins of their high densities, angular momentum, spatial structure, and variability remain a major puzzle. In addition to the dense equatorial disks, it is well known that Be stars have substantial winds, roughly isotropic, with typical $\\Mdot \\simeq 10^{-(10 \\to 8)} \\Msunyr$ and terminal velocities near $\\vinf\\simeq 1000~\\kms$ ( Marlborough \\& Peters 1986; Grady, Bjorkman, \\& Snow, 1987; Prinja 1989)  Hot stars exhibit rotational line broadening indicative of surface equatorial rotation speeds which are high, but less than about 80 percent the centrifugal (Keplerian) limit. Thus rotation alone is not sufficient to centrifuge material off the equator to form a disk. Owing to the high luminosity of the stars, wind material leaves the general stellar surface due to radiative forces on the line opacity. Line-driving, however, only works well where the wind material is not too optically thick, and has a sufficient velocity gradient for the absorbing atomic lines to sweep through the stellar continuum spectrum. Observational evidence suggests that Be star equatorial disks are too dense and too slowly expanding for line driving alone to support them against gravity.  One possible mechanism proposed to create disks is the Wind Compressed Disk (WCD) model of Bjorkman \\& Cassinelli (1993). In this, stellar wind matter accelerating out from an intermediate latitude on a rotating star has a trajectory orbital plane  which crosses the equatorial plane. The continuous outflows of such matter from the upper and lower hemispheres `collide' in the equatorial plane and form a compressed disk with density enhancement by a factor of $\\sim 100 $, which is on the low side for what is needed observationally. Irrespective of other theoretical issues raised concerning the WCD model (Owocki, Cranmer, \\& Gayley 1996), the fact is that all the observational indications seem to point toward a disk which is mainly moving azimuthally, supported against gravity by centrifugal forces, i.e.,~a near Keplerian disk $(v_r \\ll \\vphi \\approx \\sqrt{GM/r})$ rather than one which is mainly outflowing $(v_r \\gg \\vphi)$.  The near Keplerian and high density disk inferred from observations poses questions which we address in this paper: How does the material acquire more angular momentum than it possessed when it left the stellar surface? How does such an equatorial disk accumulate material, given that its density and small radial velocity gradient prohibit radiative driving from overcoming gravity in the equatorial plane? Proposed answers to these two questions have included schemes in which matter is ejected from the star selectively in the pro-grade direction to boost its rotation up to orbital values. In particular, Hanuschik (1999) considered how an outflow of particles with an isotropic Maxwellian distribution could create an orbital regime comprising the faster pro-grade particles with the rest falling back.  Owocki \\& Cranmer (2001) proposed non-radial pulsations with an anisotropic radiation force favoring pro-grade gas flow. While these ideas are interesting, the direct centrifuging action that a strong enough stellar magnetic field could provide seems to be a more direct way of producing the disk.  The possible role of magnetic fields in Be disk production has been mentioned quite often (Friend \\& MacGregor 1984; Poe \\& Friend 1986; Ignace, Cassinelli, \\& Bjorkman 1996). These dealt with applications of magnetic rotator wind theory, which explored the various degrees of control of the outflowing gas by the magnetic fields (Belcher \\& MacGregor 1976; Hartmann \\& MacGregor 1982; Poe, Friend \\& Cassinelli, 1989; Mestel 1990). The magnetic rotator models considered the effects of the field on the azimuthal motion of the gas, the terminal speed of the wind, and the equatorial mass loss rate.  Several papers have discussed aspects of the time-dependent accumulation of stellar wind matter channeled by strong magnetic fields. This has been mainly in the context of early type Chemically Peculiar (CP) stars with very strong magnetic fields (Havnes \\& Goertz 1984; Babel \\& Montmerle 1997a), but with ideas being extended recently toward applications to other hot stars and Be star phenomena (Babel \\& Montmerle 1997b; Donati \\etal\\ 2001). Specifically, Havnes and Goertz (1984) discussed the effect of the growing `magneto-spheric' density and  mechanisms which might disrupt the field, releasing energy and producing X-rays. Babel and Montmerle (1997a) applied these ideas to develop a magnetically confined wind-shocked model for X-rays from the Ap star IQ Aur and extended this (Babel and Montmerle, 1997b) to the young star $\\theta^1$ Ori C. Donati \\etal\\ (2001) applied the model to $\\beta$ Cep which has a magnetic field much lower than those on CP stars, and which, they concluded, is highly oblique to the axis of slow rotation of this star. All of these papers, however, neglected the contribution of rotation to the dynamics of the magnetically channeled wind.  This approximation cannot be valid for stars of near critical rotation since the rotational energy density of the channeled wind, and even more so of the shock compressed disk, may far exceed the outflow energy density. Nor can it say anything about how disk material acquires the quasi-Keplerian angular momentum values observed. In this paper we therefore review the problems and address the other limiting regime where the rotational term is dominant (centrifugal force balancing gravity). We do so to evaluate what field is necessary to provide sufficient torquing to produce a dense disk with quasi-Keplerian speeds.  We make use of the fact that near main-sequence B stars have a mass outflow and the channeling of that flow can lead to shock compressed disk material. The relatively high densities in that disk lead to the strong \\Halpha, excess free-free fluxes and electron scattering polarization for which Be stars are known.  We develop equations to find the field needed both to transfer angular momentum to material that is driven up from the star by line driven wind forces, and to redirect that flow toward the equatorial plane. This regime of a magnetically directed flow is non-spherical and, because there is a build up of matter, is also non-steady. The directed flow results in the formation of a disk-like structure the azimuthal motion of which is `quasi-Keplerian'. By this we mean that the velocity distribution is not strictly Keplerian but the motion is primarily azimuthal and with speeds of order the Keplerian value.  We have chosen to use an aligned field of dipole-like shape for the basic magnetic topology, $\\Bvec(r,\\theta)$, with the $r$-dependence parameterized.  As for the gas, it will for the most part be assumed to be confined by the field. Our main interest is not in the flow trajectories $\\Vvec(r,\\theta,\\phi)$, but rather in the azimuthal velocity, \\vphi(\\dist), in the equatorial plane at radial distance, \\dist. The closed field will channel the gas to a specific radial distance, \\dist$(\\lami)$, that depends on the latitude \\lami\\ the field tube had at the stellar surface, the larger \\lami\\ corresponding to larger \\dist.  Field lines from high latitudes, $\\lami \\gap \\pi/4$, near the pole, would intercept the equator at very large distances, but before reaching those distances the wind energy density will exceed that of the field and no longer be magnetically dominated, since even a dipole field declines as $d^{-3}$. Thus, there will be a broad, open, polar wind that fills the envelope volume far beyond the disk and above the closed lines of force that are channeling lower latitude material toward the disk.  For our study, we consider both the strictly closed field regions and the transition to the open ones as illustrated in Figure~\\ref{Fig:structure}.  The simplest picture of the situation would be an abrupt ``switch'' model, in which the field is assumed to be fully capable of causing solid body rotation, $\\vphi(d) \\propto d$, out to an \\Alfven\\ distance, at which the field energy density has decreased to that in the gas flow. Beyond that transition radius the gas would be dominant and in particular its angular velocity would satisfy conservation of angular momentum, $\\vphi \\propto 1/d$. In this representation, there is a discontinuous shift in \\vphi(\\dist) (from $\\vphi \\propto d$ to $ \\vphi \\propto d^{-1}$). Although this picture may be useful for understanding some basic concepts involved in the torquing of disks, we know from magnetic rotator stellar wind theory that no such abrupt shift occurs at the \\Alfven\\ point in either \\vphi\\ or the angular momentum.  To allow for a more realistic behavior of the angular speed and angular momentum transfer, we introduce a ``smooth transition'' model, in which a continuous parametric function, $\\vphi(d)$ is specified for the angular velocity as a function of distance in the equatorial plane. This function reaches a maximum, but the angular momentum continues to increase with increasing radius by a factor of about two, transmitted by the bent field lines. We are interested in cases in which $\\vphi(d)$ reaches at least to the Keplerian speed $\\vK = \\sqrt{G M/d}$. The field lines will not be considered drawn open until at least a distance, \\dmax, where the Keplerian ratio, i.e. $\\vphi(d) /\\vK(d)$ is at its maximal value. The maximal Keplerian ratio radius is identified with the \\Alfven\\ radius and, from that radius outward, we assume that the gas motion then tends to the field-free flow limit. Thus, at large $d$, it obeys conservation of angular momentum, and the particles remain in their orbital planes defined at the point where they became free from the magnetic field forces.  The subsequent flow trajectories in the far field region become similar to those in wind compression theory, either in the form of the WCD model or the Wind Compressed Zone (WCZ) model of Ignace, Cassinelli, \\& Bjorkman (1996). The WCZ model was developed to describe the structure of a wind which has a rotation that is too slow to form a shock compressed disk, but nonetheless can produce density enhancements in the low latitude regions. The enhancements in density (by a factor of 10 or so) have an effect on the observational properties of a wide range of stars. The field lines are drawn out in the far field regime as described by Ignace, Cassinelli, \\& Bjorkman (1998).  Neither the WCD nor the WCZ models can explain the quasi-Keplerian disks of Be stars.  The relation to WCD theory of the current model for a shock compressed disk is as follows. The WCD model is ``kinematic'' in that the {\\it radial velocity structure} of the outflow is chosen to obey the well-known beta law distribution, instead of being computed from the force of radiation on line opacities. In WCD theory, the velocity component \\vphi\\ is assumed to be determined by conservation of angular momentum, and thus \\vphi\\ decreases rapidly with radius and does not attain the quasi-Keplerian angular speeds inferred from observations. In the current paper, we again use a kinematic approach, but now specify an {\\it azimuthal velocity structure}, $\\vphi(d)$, that assures quasi-Keplerian velocities are reached for a sufficiently strong magnetic field. This kinematic approach allows us to derive constraints on the surface field, \\BZ, and the surface spin rate parameter, \\SZ, that are needed for a disk to form. Mass and magnetic flux conservation requirements provide information regarding the disk density distribution, the radial extent, and the timescale for variability. Where needed we also use the wind beta velocity law to provide an estimate of the flow velocities driven by the star's radiation field. The derived structure for stars with ``torqued disks'' allows us to estimate observational properties such as \\Halpha\\ line strengths, intrinsic polarizations, and continuum fluxes at a variety of wavelengths useful for comparisons of the model with observations.  A major goal is to derive lower limits on the field strengths required for the existence of centrifugally-supported disks produced by an aligned dipole-like field on a Be star. As a first rough estimate, showing that the fields required to channel a flow are not unreasonable, we consider magnetic rotator wind theory. In that theory, as explained in Ch 9 in Lamers \\& Cassinelli (1999) (hereafter L\\&C), there is a ``primary mass loss mechanism'' (in our case the line driven wind theory), which sets the stellar mass loss rate, and the terminal velocity.  An amplification of both \\Mdot\\ and \\vinf\\ can be produced by the co-rotating magnetic field. A transition occurs from the case in which the field acceleration effect is unimportant (a ``slow magnetic rotator'') to the case in which the field produces a strong effect (a ``fast magnetic rotator'', or FMR). This transition occurs when the Michel velocity, \\VM, (which is very close to the terminal velocity in FMR theory) is as large as the terminal velocity owing to the primary mass loss mechanism. The Michel velocity is related to the Poynting flux which makes the enhanced outflow energetically possible. As a starting point, we consider it reasonable to assume that the field necessary to dominate the flow and channel it to the equatorial disk should be at least comparable to the transition field strength from slow to fast magnetic rotator theory for a particular star. The Michel Velocity, \\VM, is given by quantities defined at the surface of the star; the surface field, (\\BZ), rotation rate (\\OmegZ), and mass loss rate, \\Mdot. Solving for \\BZ\\ in the equation for \\VM\\ gives a minimal field based on FMR considerations.  \\begin{eqnarray} B_{\\rm FMR} &=& \\frac{(\\Mdot \\VM^3)^{1/2}}{R^2 \\OmegZ}\\\\ &\\approx& \\frac{(\\Mdot \\vinf^3)^{1/2}}{\\SZ \\sqrt{RGM}}\\\\ & =     & 6.9\\units{\\rm Gauss} \\frac{\\Mdot_{-9}^{1/2} v_8^{3/2}}{\\SZ (R_{12} M_{10})^{1/2}} \\label{Bfmr} \\end{eqnarray} where $R$ is the stellar radius.  We set $R \\OmegZ$ equal to the surface equatorial speed $\\vZ$, which is expressed as the equatorial Keplerian speed, \\vK(R) times a fraction \\SZ. We also set $\\VM = \\vinf \\approx \\vesc$ (where \\vesc\\ is the escape speed at the stellar surface). For numerical values, we express the velocity by the ratio $v_{8} = \\vinf/(10^8~ {\\rm cm~ s^{-1}})$; the stellar mass and radius ratios by $M_{10}= M/(10 \\Msun)$, $R_{12} = R/(10^{12}\\units{cm})$; and the mass loss rate by $\\Mdot_{-9} = \\Mdot/(10^{-9} \\Msunyr)$.  A roughly similar value for the magnetic field, as in Equation~(\\ref{Bfmr}), follows from requiring that the energy density of the surface field $\\BZ^2/8 \\pi$ should exceed that of the gas $\\rho_{\\circ} v_{\\circ}^2/2$. The derived field of order 10 Gauss (for \\SZ$\\simeq 1$) is a very modest one. However, as we shall see, it is a lower bound on \\BZ\\ because it effectively uses the density in the wind, which is much below (\\ie\\ $\\simeq 10^{-3}$--$10^{-4} \\times$) the densities that are produced in the centrifugally supported disk, so that the field ($\\propto \\rho^{\\half}$) torquing the disk is 30 to 100 times larger than that in Equation (3). ",
        "conclusions": "The central point of our model is that magnetic torquing and channeling of wind flow from the intermediate latitudes on a B star can, for plausible field strengths, create a dense disk a few stellar radii in extent in which the velocity is azimuthal and of order the local Keplerian speed. The model was motivated by the fact that line profile data (e.g. Hanuschik 1996) indicate that disk emission lines of Fe II, extend to $\\Delta v > v_{\\circ} \\sin i$, and reach disk Keplerian values. Also the sharp absorption features seen in the center of those lines do not show evidence for inflow or outflow. The latter of course would be expected in the case of the WCD model. The present paper aimed to show that fields could confine and torque the flow to create a quasi-Keplerian disk of roughly the spatial extent suggested by these observations. Future work to test the real viability of the model will address two important questions:  \\begin{enumerate}  \\item[a)] Is the azimuthal velocity distribution $\\vphi(d)$ across the radial range of the dense disk predicted by our magnetic model compatible with observed line profiles? This model distribution starts as roughly corotational, \\ie\\ $\\vphi(d) \\propto d$, and transitions to a \\vphi\\ decreasing with $d$. This distribution is clearly not truly Keplerian disk $v_\\phi(d) \\propto 1/ \\sqrt{d}$. However, the typical $v_\\phi$ is comparable with $v_K$ and exists over only a rather narrow range of $d$, unless the field is very strong, which is unlikely on theoretical grounds (Maheswaran \\& Cassinelli 1988, 1992; c.f.~Figure~\\ref{Fig:mahesTHSZ}). Outside the model disk domain, \\ie\\ at $x> \\xalfven$, we predict quantitatively a somewhat flattened WCD region changing with increasing $d$ to a WCZ.  At even farther distances from the star, the wind will be nearly spherical.  All of these regions involve strong outflow velocities but with densities much below those in the quasi-Keplerian disk. It is important to estimate whether the high $v_r$ but low density distributions there are compatible with line profile data \\ie\\ does there exist an outer expanding flattened disk which affects line profiles very little but is detectable by radio observations etc. These and other diagnostic predictions of the Magnetically Torqued Disk model such as X-ray production will be treated in subsequent papers.  \\item[d)] Another issue, that is perhaps related to the $V/R$ variations is the winding up of the field - something we have ignored so far. In a model with strictly corotational $\\vphi(r)=\\vZ r/R$ out to some point and angular momentum conservation beyond there, the field would corotate in the disk region. However, in practice and for our parametric model (Equation \\ref{vphi/v0}), \\vphi$(r)$ falls behind corotation and the field will be wound up. The  resulting \\Bphi\\  component will have opposite sign on opposite sides of the disk plane and will grow until field reconnection can occur allowing a quasi-steady state with the disk matter slipping across the field lines. In such a situation gravitational control of the disk may dominate sufficiently for the disk to behave as quasi-Keplerian and allow density waves. Indeed, in the light of this, one might argue that our estimate of $\\BZTH$ is too conservative since the crucial feature of our model is that the field should deliver wind matter to the disk with sufficient angular momentum (from magnetic torquing) and with essentially no $v_r$ but only $v_\\theta$ so that shock compression can occur without any outflow (in contrast to the WCD model). This needs working out in detail but if only the wind density needs torquing action our field limits would be much reduced. Furthermore, in that case the disk $\\rho_D v_\\phi^2$ would be so high compared to $B^2/8\\pi$ that the matter would completely dominate the field once delivered into the dense disk, which would consequently be essentially Keplerian and so support density waves under the action of stellar gravity.  \\end{enumerate}   In this paper we have followed a used a kinematic description of the azimuthal velocity of a disk zone, along with constraints based on magnetic rotator theory to develop an analytic model which we call the Magnetically Torqued Disk (MTD) model, We derived the independent parameters for disk formation problem: the spin rate parameter. \\SZ, the surface field \\BZTH\\ and the stellar parameters contained in the quantity, \\Ystar. We assumed a rather simple rigid magnetic field picture, but one that allowed us to derive an estimate of the field required to produce a torqued disk for main sequence stars ranging from O3 V to B9 V. Although an unreasonably large field was found to be needed for the early O stars, we concluded that at the early B spectral class at which the Be phenomena are most prominent, the model provided about the right observational properties, and for reasonable surface field strength. Also, we considered several time scales associated with Be stars, and found start-up, growth and fill-up time scales that are in the month to decade long range that seem appropriate. For the late B stars the start-up time scale and the \\Halpha\\ equivalent widths appear to be small and we suggest that perhaps the mass loss rates for these stars need to be enhanced by centrifugal magnetic rotator forces. Modifications that are needed to address V/R variations and observational consequences at a variety of wavelengths will be considered in later papers."
    },
    "0207/astro-ph0207620_arXiv.txt": {
        "abstract": "\\sgra, the compact radio source, believed to be the counterpart of the massive black hole at the galactic nucleus, was observed to undergo rapid and intense flaring activity in X-rays with Chandra in October 2000. We report here the detection with XMM-Newton EPIC cameras of the early phase of a similar X-ray flare from this source, which occurred on September 4, 2001. The source 2-10~keV luminosity increased by a factor $\\approx$~20 to reach a level of 4~10$^{34}$~erg~s$^{-1}$ in a time interval of about 900~s, just before the end of the observation. The data indicate that the source spectrum was hard during the flare. This XMM-Newton observation confirms the results obtained by Chandra and suggests that, in \\sgra, rapid and intense X-ray flaring is not a rare event. This can constrain the emission mechanism models proposed for this source, and also implies that the crucial multiwavelength observation programs planned to explore the behaviour of the radio/sub-mm and hard X-ray/gamma-ray emissions during the X-ray flares, have a good chance of success. ",
        "introduction": "The bright, compact and variable radio source \\sgra\\ is believed to be the radiative counterpart of the 2.6~10$^{6}$~M$_{\\odot}$ black hole which governs the dynamics of the central pc of our Galaxy \\cite{mefa01}. The compelling evidence for the presence of a dark mass concentration at the Galactic Center \\cite{gen97,ghe00}, which implies the presence of a massive black hole, contrasts remarkably with the weak high-energy activity of such an extreme object. In spite of the fact that some amount of material, either provided by stellar winds from a close stellar cluster or by the hot surrounding medium, is probably feeding a moderate/low level of accretion, the total bolometric luminosity of the source amounts to less than 10$^{-6}$ of the estimated accretion power \\cite{mefa01,gol01}.  This has motivated the development of several black hole accretion flow models with low radiative efficiency, some of which have also been applied to binary systems, low luminosity nuclei of external galaxies and low luminosity active galactic nuclei. These models include spherical Bondi accretion in conditions of magnetic field sub-equipartition with a very small Keplerian disk located within the inner 50 Schwarzschild radii (R$_S$), large hot two-temperature accretion disks dominated by advection (ADAF) or non-thermal emission from the base of a jet of relativistic electrons and pairs, and some other variants or combination of the above (see the review by Melia $\\&$ Falcke (2001)). However any such model still predicts some level of X-ray emission from \\sgra\\ and determining the properties of such emission would constrain the theories of accretion and outflows in the massive black holes and in general in compact objects.  The search for high energy emission from \\sgra\\ with focussing X-ray telescopes dates back to the end of the '70s \\cite{pre94,bag01a,gol01}, but has recently come to a turning point with the remarkable observations made with the Chandra X-ray Observatory in 1999 and in 2000. Baganoff et al. (2001a) first reported the detailed 0.5$''$ resolution images obtained with Chandra in the range 0.5-7~keV, which allowed, finally, the detection of weak X-ray emission from the radio source. The derived luminosity in the 2-10~keV band was 2~10$^{33}$~erg~s$^{-1}$, for a distance of 8~kpc, and the measured spectrum was steep, with power law photon index of 2.7. Marginal evidence that the source is extended on a 1$''$ scale was also reported, but at low significance level. Then, in October 2000, the same source was seen to flare up by a factor of $\\approx$~45 in a few hours \\cite{bag01b}. The luminosity increased from a quiescent level similar to the one measured in 1999 to a value of 10$^{35}$~erg~s$^{-1}$. The flare lasted a total of 10~ks but the shortest variation took place in about 600~s, implying activity on length scales of $\\approx$~20~R$_S$, for the above quoted mass of the galactic center black hole. Evidence of spectral hardening during the flare was also reported by the authors who determined a source power law photon index during the event of 1.3 ($\\pm$~0.55), significantly flatter than observed during the quiescent state. These results constrain models of the accretion flow and radiation mechanism for \\sgra. A confirmation of the Chandra results and in particular a better determination of the flaring properties of the source are therefore crucial for the modeling of the physics of the Galactic Center and in general for the theories of accretion in black hole systems.  XMM-Newton, the other large X-ray observatory presently in operation, features three large area X-ray telescopes coupled to three CCD photon imaging cameras (EPIC) operating in the 0.1-15 keV range and to two reflection grating spectrometers (RGS) working in the 0.1-2.5 keV band \\cite{jan01}. Although its angular resolution (6$''$ FWHM) is insufficient for properly resolving \\sgra\\ in quiescence, the high sensitivity and wide spectral range of XMM-Newton allow deep studies of the X-ray emission of such a complex and crowded region like the Galactic Center. Indeed an intense flare such as the one seen by Chandra can be easily detected with XMM-Newton thanks to its large effective area, and its timing and spectral properties can be studied.  The Galactic Center region is one of the priority targets of the XMM-Newton mission and was included in the guaranteed time program. Visibility constraints and solar flare events have however delayed the monitoring of the very center of our Galaxy. A complete pointed observation was finally performed in fall 2001 and in this letter we report the detection with XMM-Newton of another X-ray flare from \\sgra\\ which occurred at this time. ",
        "conclusions": "The XMM-Newton discovery of a new X-ray flare of \\sgra\\ in September 2001 confirms the results obtained in the earlier Chandra observations. XMM-Newton observed only the first part of the flare, but the recorded event is fully compatible in intensity and time scale with the early phase of the flare seen by Chandra.  The count rate within 10$''$ from \\sgra\\ increased in 900~s by a factor 3, but if attributed to \\sgra\\, it implies that this source brightened by a factor about 20-30, which is compatible with the increase in the first 1000~s of the flare observed by Chandra. We have not detected the maximum in the flare rise. Therefore we cannot strictly apply the travel light argument to estimate the size of the emitting region. If we assume that the flare duration (900~s) we observed is the shortest time scale of variation of the present event, it corresponds to a size of about 30~R$_S$. This limit is a factor 1.5 larger than the shortest scale estimated with the Chandra data and does not constrain further the geometry of the region.  On the other hand the detection of another such a flare indicates that the event is not rare. The total reported observation time with Chandra amounts to $\\approx$~75~ks. Considering the XMM-Newton 26~ks exposure, the duty cycle of such event is 0.11 (= 11~ks / 101~ks), but it would increase to 0.18 (= 20~ks / 110~ks) if we assume that the flare we detected for only 1000~s would last for 10~ks. Though not much different than the value determined with Chandra, this estimate of the active time fraction of the source is now based on 2 events and it is therefore more significant.  The radio source on the other hand has been observed many times and the detected flux variability has never exceeded a factor 2 \\cite{zha01}. This implies that it is unlikely that radio or sub-mm emission present a comparable increase in flux. If this is confirmed the flare may not be due to a change in the accretion rate, since this variation would lead, at least in models which attribute the bulk of X-ray emission to self synchrotron Compton emission, to a comparable increase of radio and sub-mm radiation \\cite{mar01}.  The X-ray flare from \\sgra\\ cannot be explained by pure ADAF models (Narayan et al. 1998) as in these models the emission is due to thermal bremsstrahlung from the whole accretion flow and arises from an extended region (between 10$^3$ - 10$^5$ R$_S$) which cannot account for such rapid variability. Models which predict emission from the innermost regions near the black hole involve a mechanism acting either at the base of a jet of relativistic particles \\cite{mar01} or in the hot Keplerian flow present within the circularization radius of a spherical flow \\cite{mel01,liu02}. In both cases a magnetic field is present in the flow and the sub-mm radiation is attributed to optically thin synchrotron emission from the inner region, while the X-rays are produced by the synchrotron self-Compton (SSC) mechanism whereby radio to mm photons are boosted to X-ray energies by the same relativistic or subrelativistic electrons that are producing the synchrotron radiation. Large flux variations can be produced by a change in accretion rate or, in the jet model, by additional heating of the electrons caused for example by magnetic reconnection. The second mechanism would increase (and harden) the X-ray flux without significantly increasing the radio and sub-mm part of the spectrum and therefore it could be more compatible with the lower amplitude of radio changes compared to X-rays \\cite{mar01}.  However even emission from a circularized flow can provide low or anti correlation of the radio emission with the X-rays if the radiation mechanism for the X-rays is bremsstrahlung rather than SSC \\cite{liu02}. The sub-mm and far IR domain on the other hand would in this case show a large correlated increase, but at these frequencies the measurements have not been frequent enough to settle the issue. Though the exact modelling of radiation process depends on viscosity behavior and other uncertain details, the observed hardening of the spectrum during the flare indeed favours the bremsstrahlung emission mechanism in this model rather than the SSC one \\cite{liu02}. More compelling constraints on the models will be set when simultaneous observations in radio/sub-mm and X-ray wavelengths of such a flare are obtained.  Correlated radio and X-ray observations are indeed crucial because, althought the amplitude of the radio variability is low compared to the event recorded in X-rays, an intriguing correlation seems to be present between the X-ray flares and the rise of the radio emission. Indeed Zhao et al. (2001), using Very Large Array (VLA) data collected over two decades, detected a periodicity in the \\sgra\\ radio variability, with a 106 days cycle and a characteristic timescale of 25 days. Baganoff et al. (2001b) already remarked that the October 2000 X-ray flare occured at a radio-cycle phase corresponding to the beginning of the radio peak. We have computed the 106 days radio cycle phase of the X-ray flare that we detected with XMM-Newton and found that it differs by only 6 days from the phase of the flare detected with Chandra. The flare occurred at the day 64 in the light curve of Fig.~3 of Zhao et al. (2001), while the Chandra flare took place at phase 70 day and the 1.3~cm radio peak rise extends roughly from day 55 to day 75. Even though the light curve radio peak is wide and several other structures are present, both X-ray flares detected till now are very close in phase and take place during the rising part of the main radio flare. We have also compared the time of the flare to a recent radio light curve of \\sgra\\ obtained at 1.3 cm and 2 cm with the VLA between March and November 2001 \\cite{yua02}. The X-Ray flare occurred 1-2 days after a local maximum of the curve, but no radio data points are reported for the day when our XMM-Newton observation took place.  It will be also important to study the shape of the flare spectrum at energies higher than 10 keV to fully understand the radiation mechanism producing the high energy tail. In particular by measuring the high energy cut-off of the spectrum one could determine the electron temperature for a thermal emission or the Lorentz factor for non-thermal processes. We estimated that such a flare should be marginally visible in the range 10-60 keV with the low energy instruments onboard the new gamma-ray mission INTEGRAL, to be launched in October 2002, if the spectrum extends to these energies with the slope observed with Chandra and XMM-Newton.  Our more secure estimation of the duty cycle of the flares shows that multiwavelength observations of \\sgra\\ which involve XMM-Newton or Chandra will have good chance of observing an X-ray flare provided the simultaneous coverage is of the order of 100~ks."
    },
    "0207/astro-ph0207416_arXiv.txt": {
        "abstract": "Most of the Milky Way's evolved massive stellar population is hidden from view. We can attempt to remedy this situation with near-infrared observations, and in this paper we present our method for detecting Wolf-Rayet stars in highly extincted regions and apply it to the inner Galaxy. Using narrow band filters at K-band wavelengths, we demonstrate how WR stars can be detected in regions where they are optically obscured. Candidates are selected for spectroscopic follow-up from our relative line and continuum photometry. The final results of applying this method with a NIR survey in the Galactic plane will provide a more complete knowledge of the structure of the galactic disk, the role of metallicity in massive stellar evolution, and environments of massive star formation. In this paper we briefly describe the survey set-up and report on recent progress. We have discovered four emission-line objects in the inner Galaxy: two with nebular emission lines, and two new WR stars, both of late WC subtype. ",
        "introduction": "Optical surveys within our Galaxy are severely hampered by dust obscuration; therefore complete samples must be obtained with longer wavelength observations. Here we describe a survey in the Galactic plane at K-band wavelengths, where the extinction is much lower than traditional V-band surveys. Our scientific driver is the discovery of the young stellar population in our Galaxy through the detection of evolved massive stars. These stars have strong emission lines, which makes them relatively easy to detect using narrow-band filters.  As a massive O star evolves, its spectrum becomes dominated by emission lines, arising either in a dense stellar wind, or in circumstellar material produced by mass loss. Stellar emission lines are most pronounced in Wolf-Rayet (WR) stars which have lifetimes $<$~10~Myr, and thus are excellent tracers of recent star formation, and so also Galactic structure. A complete sample of Galactic WR stars will also enable us to understand the distribution of WN and WC sub-types in a high metallicity environment. Additionally, WR stars are critical components in our quest to understand how star formation proceeds. For example, most of the previously known WRs are relatively isolated, but recent searches in the IR have found a plethora of these objects in stellar clusters near the Galactic center (e.g. Figer et al. 1999). ",
        "conclusions": ""
    },
    "0207/astro-ph0207285_arXiv.txt": {
        "abstract": "I review the status of large-scale structure studies based on redshift surveys of galaxies and clusters of galaxies.  In particular, I compare recent results on the power spectrum and two-point correlation correlation function from the 2dF and REFLEX surveys, highlighting the advantage of X-ray clusters in the comparison to cosmological models, given their easy-to-understand mass selection function.  Unlike for galaxies, this allows the overall normalization of the power spectrum to be measured directly from the data, providing an extra constraint on the models.  In the context of CDM models, both the shape and amplitude of the REFLEX P(k) require, consistently, a low value for the mean matter density $\\Omega_M$.  This shape is virtually indistinguishable from that of the galaxy power spectrum measured by the 2dF survey, simply multiplied by a constant cluster-galaxy bias factor.  This consistency is remarkable for data sets which use different tracers and are very different in terms of selection function and observational biases.  Similarly, the knowledge of the power spectrum normalization yields naturally a value $b\\simeq 1$ for the bias parameter of $b_J$-selected (as in 2dF) galaxies, also in agreement with independent estimates using higher-order clustering and CMB data.  In the final part, I briefly describe the measurements of the matter density parameter from redshift space distortions in galaxy surveys, and show evidence for similar streaming motions of clusters in the REFLEX redshift-space correlation function $\\xi(r_p,\\pi)$.  With no exception, this wealth of independent clustering measurements point in a remarkably consistent way towards a low-density CDM Universe with $\\Omega_M\\simeq 0.3$. ",
        "introduction": "The last couple of years have witnessed an impressive series of achievements in the field of large-scale structure, thanks to \\footnote{Review to appear in {\\it DARK2002, 4th Heidelberg Int. Conf. on Dark Matter in Astro- and Particle Physics}, (Cape Town, February 2002), H.-V. Klapdor-Kleingrothaus \\& R. Viollier eds., Springer} new large surveys of galaxies and clusters of galaxies. The enthusiasm for new results on the clustering of galaxies and clusters has been strenghtened by the unprecedented possibility to couple these to the anisotropies in the cosmic microwave background over an overlapping range of scales (see contributions by Melchiorri and Cooray, this volume).  In this brief review I have tried and provide a general guide for the non-specialist through some of the large-scale structure results. Clearly, such a review is far from being complete, although the references indicated should allow the reader to find further links to the available literature on the subject (before February 2002).  I therefore apologize to those colleagues whose work has not been adequately covered. ",
        "conclusions": "We are definitely in a golden age for observational cosmology and in particular for the study of large-scale structure. We never had such a wealth of diverse data at our disposal, through which we are pinning down the values of cosmological parameters to a high accuracy (e.g. \\cite{CMB}). The observational facts we have reviewed here contribute to further reinforce the remarkable convergence among different observables (CMB, large-scale structure, distant Supernovae, cluster evolution, to mention a few) towards a model with flat geometry ($\\Omega_{total}=1$) provided by the combination of a dominating Cold Dark Matter component ($\\Omega_{M} = \\Omega_{CDM}+\\Omega_{baryon}\\simeq 0.3$, with $\\Omega_{baryon}\\simeq 0.04$) and a Dark Energy of unknown nature (the cosmological constant, $\\Omega_{\\Lambda}\\simeq 0.7$ ).  Still, we cannot avoid to note that such wonderful ``standard'' cosmological model is full of ``unseen'' ingredients, as a dark matter nobody has detected so far (but see contribution by P. Belli in this volume) and a dark energy we have little idea where it could come from.  Seen from outside, this might look as an almost epicyclic model and I believe understanding its foundations provides one of the major challenges for particle physics and cosmology in the next decade.  {\\bf Acknowledgments.} I thank the organizers of the DARK2002 meeting for inviting me to give this review.  I am grateful to all my collaborators for all our common results discussed here, in particular P. Schuecker, C. Collins and H. B\\\"ohringer.  I thank I. Zehavi and E. Hawkins for providing their clustering results in electronic form."
    },
    "0207/astro-ph0207599_arXiv.txt": {
        "abstract": "We present results from observations of persistent black hole candidates with the High Energy Transmission Gratings aboard the {\\em Chandra X-ray Observatory}. The sources include LMC X-1, LMC X-3, GRS 1758-258, and Cyg X-1. Along with the published results on 1E1740.7-2942, we have completed a high-resolution spectroscopic survey of such systems. The observed X-ray spectra of LMC X-1 and LMC X-3 show no prominent discrete features, while absorption edges (Mg K and Si K) are detected in the spectrum of GRS 1758-258. The edges are likely to be of interstellar origin. In most cases, the X-ray continuum can be described well by models that are often adopted in low-resolution studies of black hole candidates: a multi-temperature disk spectrum plus a Comptonization component. However, the relative contribution of the two components varies greatly among different sources. For instance, only the disk component is present for LMC X-1 and GRS 1758-258, while the Comptonized component is required for other sources. We discuss general issues related to obtaining disk parameters from modeling X-ray continuum. ",
        "introduction": "Prior to the launch of {\\em Chandra X-ray Observatory} (Chandra), emission lines and absorption lines or edges had already been detected in the X-ray spectrum of a few black hole candidates (BHCs), with X-ray spectrometers of low to moderate resolution (e.g., Barr et al. 1985; Kitamoto et al. 1990; Done et al. 1992; Ebisawa et al. 1996; Cui et al. 1998; Ueda et al. 1998; Cui et al. 2000; Feng et al. 2001). The features often appear in the energy range 6--8 keV and are, therefore, attributed to emission or absorption processes involving iron K-shell electrons. Although the exact location of the emitting or absorbing material is often debatable, there is evidence that, at least for some black hole candidates, the observed iron $K\\alpha$ line appears to originate in the innermost region of the accretion disk, very close to the central black hole. If this proves to be the case, the profile of the line would be distorted by the strong gravitational field of the hole (Fabian et al. 1989; Tanaka et al. 1995) and could be carefully modeled to constrain the intrinsic properties of the hole, such as its mass and spin (Laor 1991; Bromley et al. 1997).  Two major advances were brought about by the CCD spectrometer aboard {\\em ASCA}. First, several pairs of Doppler-shifted emission lines were observed and attributed to the relativistic jets in SS 433 (Kotani et al. 1996). A recent observation of the source with {\\em Chandra} confirmed the presence of these lines from the jets and provided many more details (Marshall et al. 2002), thanks to the improved spectral resolution of the {\\em High-Energy Transmission Gratings} (HETG) aboard. The studies of such lines have proven fruitful in gaining insights into the dynamics and physical conditions of the jets in SS 433 (Marshall et al. 2002; Kotani et al. 1996). Similar Doppler shifted lines also seem to exist in the spectrum of another well-known jet source, 1E1740.7-2942, based on a recent {\\em Chandra} HETG observation (Cui et al. 2001). Second, the {\\em ASCA} observations of micro-quasars GRS 1915+105 and GRO J1655-40 revealed the presence of absorption lines (Ueda et al. 1998; Kotani et al. 2000). The results are confirmed by a recent {\\em Chandra} observation of GRS 1915+105 (Lee et al. 2001). The absorption lines were interpreted as resonant absorption lines due to highly ionized (helium or hydrogen like) ions in a non-spherical configuration. There was speculation as whether such lines were characteristic of microquasars. This is clearly not the case, as shown by recent {\\em Chandra} observations of Cyg X-1 (Schulz et al. 2002; Marshall et al. 2001). The obtained HETG spectra of Cyg X-1 show numerous narrow absorption lines due to H-like or He-like ions of various elements. Therefore, the phenomenon may be common for BHCs.  The presence of energetic electrons in the vicinity of accretion disks has always been a key ingradient in nearly all attempts to model the X-ray continuum of a BHC. The observed power-law spectrum at high energies ($\\gtrsim 10$ keV) is invariably attributed to inverse Comptonization of soft photons by such electrons in an optically thin but geometrically thick configuration (review by Tanaka \\& Lewin 1995). At low energies, contribution from the accretion disk becomes important and often dominant. The $\\alpha$-disk solution (Shakura \\& Sunyaev 1973) is often adopted to describe the observed spectrum of the disk emission; this has led to the formulation of the popular multi-color disk (MCD) model (Mitsuda et al. 1984). The model has worked satisfactorily in describing low-resolution data, but so have other models (such as a simple blackbody). The high-resolution data from {\\em Chandra} (and {\\em XMM-Newton}) may now allow us to see subtle differences in the shape of the continuum as predicted by various models.  The MCD model contains fundamental physical information about the accretion disk, such as its temperature (distribution) and the distance from the inner edge of the disk to the black hole. Existing observational evidence shows that, under certain conditions, the accretion disk extends down to the last (marginally) stable orbit (Tanaka \\& Lewin 1995 and references therein). Therefore, modeling emission from the disk may lead to a determination of the radius of the last stable orbit and thus the properties of the black hole (e.g., its mass and angular momentum; Zhang, Cui, \\& Chen 1997; Tanaka \\& Lewin 1995). However, the local X-ray spectrum of an accretion disk is not a perfect blackbody, because in the hot inner portion of the disk the opacity is primarily due to electron scattering as apposed to free-free absorption. The spectrum can instead be described by a ``diluted'' blackbody (Ebisuzaki, Hanawa, \\& Sugimoto 1984), with its temperature higher than the effective temperature. Attempts have been made to quantify such effects theoretically in terms of a spectral hardening factor (Shimura \\& Takahara 1995), but the results are still quite uncertain. In particular, it has been shown that the hardening factor seems to vary with mass accretion rate (Merloni, Fabian, \\& Ross 2000). On the other hand, the derivation of disk properties depends critically on this factor. The question becomes whether we can determine it observationally.  In this paper, we present results from a survey of persistent BHCs with the HETG. The motivations for conducting the survey included: (1) using emission or absorption lines, if present, to study the physical properties of the emitting region and its environment; (2) obtaining profiles of those emission lines that might originate in the accretion disk, with a goal of inferring the properties of the black hole; (3) directly detecting X-ray emission from the jets in microquasars; and (4) better characterizing the disk component of the continuum, making use of good low-energy sensitivity as well as excellent energy resolution of the HETG. The survey might also have provided data for systematic studies of line production in persistent BHCs. The results are likely to be relevant to transient BHCs, which constitutes the majority. ",
        "conclusions": "Unlike Cyg X-1 (Schulz et al. 2002; Marshall et al. 2001; Miller et al. 2002) or 1E1740.7-2942 (Cui et al. 2001), we found no evidence for the presence of any emission or absorption lines in the spectra of LMC X-1, LMC X-3, and GRS 1758-258. This is at least partly due to the lack of statistics, as evidenced by the absence of interstellar absorption edges. In the case of GRS 1758-258, we did positively detect Mg K and Si K edges, which are most likely of interstellar origin.  For BHCs, X-ray illumination of the accretion disk is thought to be very important and is thus necessarily taken into account in the models. The re-processing of incident X-rays by the disk is often used to explain some of the continuum features. This process is expected to be accompanied by the production of fluorescent lines. On the other hand, depending on the ionization state of the disk, the lines can be destroyed by the Auger process (Ross \\& Fabian 1993). The lack of line emission can, therefore, shed light on these physical processes. Unfortunately, in our case, the quality of data is not good enough to constrain the models.  We have shown that for most sources in our sample the observed X-ray continuum can be well described by a two-component model including a MCD component and a Compton component. Such a model has been successfully applied to BHCs in general previously. The low-energy sensitivity of the HETG has allowed us to reliably model the emission from the hot inner region of the accretion disk (i.e., the MCD component). As in previous studies, however, the coupling between the two components still makes it challenging to completely separate them. Luckily, in two cases, LMC X-1 and GRS 1758-258, the data requires only the MCD component, indicating the X-ray photons from these sources originate entirely from the disk. Here, the derived disk parameters should, in principle, be quite reliable.  There are three important issues that make it uncertain to extract the physical properties of the accretion disk from the measured quantities (e.g., Zhang 1999; Cui 2001). First of all, it is still not clear what is the origin of seed photons for the inverse Comptonization process. If the accretion disk provide all the seed photons, as often assumed, these photons will not appear in the observed spectrum of the disk (i.e., the MCD component). Instead, they appear in the Compton component. In this scenario, therefore, the photons detected are originally all from the accretion disk. To compute intrinsic luminosity of the disk we must then add photons in the Compton component to the MCD component. It is the intrinsic disk luminosity that should be used to derive such important properties as the radius of the inner edge of the disk. In practice, however, the disk radius is often derived directly from the best-fit normalization of the MCD component, which is clearly erroneous, especially when the Compton component is seen to be strong. This erroneous procedure is likely to be responsible for the claimed correlation between the inner disk radius and the spectral hardness of the source: the stronger the Compton component the smaller the radius, since this is exactly what would be expected if the disk is the primary source of seed photons for Comptonization. Unfortunately, the derivation of the intrinsic luminosity of the disk (by taking into account the loss of disk photons to the Comptonization process) depends on many assumption, such as the exact origin of seed photons and the spatial distribution of Comptonizing electrons with respect to the disk, and is thus very model dependent. Robust results can be obtained only for cases in which the Compton component is either absent or weak enough to produce negligible effects.  Secondly, as discussed in \\S~1, the local spectrum of the inner portion of the disk is not a blackbody, because the opacity is predominantly determined by electron (Compton) scattering as opposed to free-free emission (e.g., Ebisuzaki, Hanawa, \\& Sugimoto 1984). Strictly speaking, therefore, the spectral shape of the disk emission is that of a (saturated) Comptonized spectrum, with the spectrum of seed photons being of MCD shape. The ratio of the temperature of electrons in the disk to that of seed photons gives the so-called ``color correction factor'' (or more precisely, spectral hardening factor). This factor is critical to deriving the radius of the inner edge of the disk (see discussion in \\S~1). A possible way to determine the hardening factor observationally is, therefore, to model the disk spectrum with a Compton component. We investigated this possibility by replacing the MCD component (see Table 2) in the model with ``comptt''. The new model fit the data equally well, if not better, for all cases. The results are summarized in Table 3. As a sanity check, we expect that the temperature ($T_e$) of Comptonizing electrons in the disk should be close to the effective temperature of the disk, when the optical depth is large. This is indeed the case (comparing results in Table 3 and Table 2), although the measured values of the optical depth are much smaller than that expected from the standard $\\alpha$ disk model (Shakura \\& Sunyaev 1973). It should be noted that ``comptt'' assumes a Wien (as opposed to MCD) spectrum for seed photons (Titarchuk 1994). Empirically, we found that in order to fit the peak of a MCD spectrum with a Wien function the temperature of the inner disk ($T_{in}$ in MCD) must be equal to 2.7 times that of the Wien distribution ($T_0$). Therefore, the spectral hardening factor is simply given by $f = T_e/(2.7T_0)$. With large uncertainties (mostly due to poor constraints on the seed photon temperature), the values of $f$ were found to be $2.6$, $\\gtrsim 0.3$, $1.6$, and $4.7$, respectively, for the sources in Table~3. While some of these values seem reasonable (compared to the results of Merloni et al. 2000), others are too high (but keep in mind the large uncertainties). Data of much improved statistics, especially at low energies, is required to constrain the temperature of seed photons.  Finally, the situation can be further complicated by the presence of a ``warm layer'' just above the disk (Zhang et al. 2000). Such a layer can be due to the heating of the disk by an illuminating hard X-ray source (e.g., Nayakshin \\& Melia 1997; Mistra et al. 1998), although observation evidence suggests that the layer exists even in the absence of any non-disk emission (Zhang et al. 2000). The presence of a low-density warm layer is in fact supported by our results (i.e., relatively small optical depth of the soft Compton component for all sources; see Table 3). To account for the warm layer, yet another Compton component needs to be added to the model, which is not warranted by the quality of our data. This can perhaps explain somewhat erratic values of the spectral hardening factor.  To conclude, we emphasize that although modeling the disk continuum is a viable approach to deriving physical parameters of the disk the derivation is complicated by the issues discussed. Progress can be made by improving the quality of the data, especially at the lowest energies. Equally important is simultaneous broadband coverages that make it possible to reliably separate out the disk component from other components. The value of broad spectral coverage was demonstrated by the modeling of the Cyg X-1 spectrum: the traditional MCD model was shown to fail in this case."
    },
    "0207/astro-ph0207550_arXiv.txt": {
        "abstract": "Two recent works have analyzed a solar large and steady coronal loop observed with Yohkoh/SXT in two filter passbands to infer the distribution of the heating along it. Priest et al. (2000) modelled the distribution of the temperature obtained from filter ratio method with an analytical approach, and concluded that the heating was uniform along the loop. Aschwanden (2001) found that a uniform heating led to an unreasonably large plasma column depth along the line of sight, and, using a two component loop model, that a footpoint-heated model loop (with a minor cool component) yields more acceptable physical solutions.  We revisit the analysis of the same loop system, considering conventional hydrostatic single loop models with uniformly distributed heating, and with heating localized at the footpoints and at the apex, and an unstructured background contribution extrapolated from the region below the analyzed loop.  The flux profiles synthesized from the loop models have been compared in detail with those observed in both filter passbands with and without background subtraction; we find that background-subtracted data are fitted with acceptable statistical significance by a model of relatively hot loop ($\\sim 3.7$ MK) heated at the apex, with a column depth $\\sim 1/10$ of the loop length. In discussing our results, we put warnings on the importance of aspects of data analysis and modeling, such as considering diffuse background emission in complex loop regions. ",
        "introduction": "\\label{sec:intro}  The most recent solar X-ray and UV missions (Yohkoh, SoHO, TRACE) have produced a significant amount of spectacular data and images of the corona. The improved spatial resolution and sensitivity of the detectors have revealed the highly structured and variable nature of the corona and have stimulated the interest of solar physics community.  Whereas the morphological interpretation of the images is often straightforward and has highlighted the presence of a variety of new structures and phenomena, the physical interpretation of the data and the inference of indications about items such as the heating of the corona is a much more delicate matter. A possible problem is that many physical and instrumental effects act concurrently and non-linearly. The coronal plasma is highly thermally conductive {\\it along} the magnetic field lines but much less {\\it across} them.  As a result, non-uniform and localized heating sources which ignite coronal loops will at the same time be elusive, because the temperature increases rapidly along the whole loop, and determine a significant thermal stratification across the loop and along the line of sight. Since the coronal plasma is mostly optically thin, a telescope observing a loop will then collect the emission from multiple thermal components along the line of sight, filter them with different weights, depending on its passband, and sum them up to yield a single data number for each image pixel. The reconstruction of the original thermal structure of the plasma column in a pixel is a difficult task and some attempts have been made possible by multi-line observations with SoHO (Schmelz et al.  2001). When analyzing data from wide-band imaging instruments such as Yohkoh/SXT, one approach is to obtain maps of the weighted average temperature of the various thermal components along the line of sight. The interpretation of temperature distributions is somewhat easier along bright loop structures, because there one has the reasonable expectation that one thermal plasma component has a much larger emission measure than the others, and therefore dominates the emission and can be determined reliably.  This vision has driven recent studies of the thermal structure along coronal loops, and of its usage as indicator of the distribution of the heating deposition within the loop (Priest et al. 2000, Aschwanden 2001). The concept of these works was to compare the temperature distributions of bright coronal loops derived from the data with those obtained from models of thermally structured single coronal loops (or combinations of them).  Priest et al. (2000) used an analytical approach and devised a method to infer the heating distribution along the loops from the analysis of temperature profiles, showing that along several bright steady-state large-scale loops observed at the solar limb the temperature distribution indicates uniformly distributed heating. Aschwanden (2001) revisited the analysis of one of the loops selected by Priest et al. (2000), and found that the previous approach led to inconsistent results and, in particular, to an unreasonably large plasma column depth required to match the emission measure inferred from the data.  Using a two component loop model to explain the observations, he also found that the data are consistently fitted instead by a model coronal loop with a heating highly localized at the footpoints and with a cool background component.  In this work we present a further analysis of the same loop system, motivated as follows.  The selected loop system is relatively bright, well outstanding at the solar limb. We might then suppose relatively large emission measure, high plasma pressure, and assuming that it is a steady structure, we can also infer, according to loop scaling laws (Rosner et al. 1978), high temperature at their apex, maybe higher than $\\sim 2$ MK, which is the temperature obtained from the previous data analyses.  On the other hand, neither previous analyses include, or include only partially, a component which instead may be important: ``the loop system is embedded in a hazy background'' (Aschwanden 2001).  This background is significant, especially in the region of the loop apex, where the loop signal is fainter than elsewhere, due to the plasma gravitational stratification. One possible hypothesis, that we pursue here, is that the haze may be the effect of the presence of many other disordered faint loops, or of a streamer, which extend over the whole loop system and intersect the analysed structure along the line of sight. Such background may not be described as a simple hydrostatic equilibrium, and deserves proper attention, since it may affect systematically the filter ratio and temperature values. Another item stimulating our analysis is the finding by Aschwanden (2001) of a solution with a very small scaling height of the heating and temperature maxima at the loop footpoints, quite a puzzling result in the light of previous well-established loop modeling (Rosner et al. 1978, Serio et al. 1981).  Here we revisit the loop system, describing the main loop component with conventional hydrostatic single loop models and deriving the background component directly from the data.  We compare the brightness profiles observed along the loop in the two SXT filter passbands with and without background subtraction to those obtained from the loop models. The modeling explores three different scenarios of heating deposition:  uniform distribution, localized at the footpoints and localized at the apex, i.e. a more complete sampling than in Aschwanden (2001).  The model results are compared to data with detailed fitting procedures, and attention is paid to the values of column depth obtained to match the observed fluxes.  Section~\\ref{sec:model} describes the analysis of the data and their comparison with loop models; Sect.~\\ref{sec:disc} discusses the results and their implications. ",
        "conclusions": "\\label{sec:disc}  This work has been stimulated by the contrasting results obtained from the analysis of the same loop system observed with Yohkoh/SXT by two subsequent works. In the earlier one (Priest et al. 2000) the data, and in particular the temperature distribution along the loop, were best matched by a loop model with a uniform heating distribution. The later one (Aschwanden 2001) found instead that a model dominated by a heating deposited at the loop footpoints, with a minor cooler component, provides more self-consistent results, because it predicts also reasonable values of the plasma column depth.  The present work extends the analysis of the same loop system, by including models with heating at the loop apex, and by considering a significant background emission possibly due to the intersection with other structures, such as disorganized fainter loops and/or a streamer, along the line of sight, which adds to the loop emission. The analysis shows that, if background emission is not subtracted, none of single loop models explored acceptably fits the data, indicating that the system cannot be described only as in a simple hydrostatic equilibrium.  A base-heated loop model with maximum temperature 2.2 MK yields the best results on unsubtracted data, but, besides the high $\\chi^2$, it involves a column depth larger than the solar radius and trends not matching the data in some regions along the loop.  We did not explore loop models with heating confined in a narrow region near the loop footpoints, as in Aschwanden (2001): such heating leads to a temperature maximum localized at the loop footpoints, and such configuration is well-known to be unstable (Rosner et al. 1978, Serio et al. 1981).  On the other hand, the coronal plasma density of $10^{10.3}$ cm$^{-3}$ predicted by the base-heated ``best'' model loop in Aschwanden (2001) is unusually large for large scale coronal structures (e.g. Vaiana \\& Rosner 1978).  The best-fit uniformly-heated loop (with maximum temperature 3.3 MK) matches the data slightly worse than the best-fit base-heated model, but implies much more reasonable values of the column depth.  We do find a model fitting acceptably the background-subtracted data: a loop heated at the loop apex, with a maximum temperature of 3.7 MK (model 17 in Table~\\ref{tab:data_models}).  The corresponding column depth ($\\sim 1/10$ of the loop length) represents a reasonable aspect for a loop system.  Although the lower statistics of background-subtracted data should allow in general to obtain more easily an acceptable fitting, the model above is {\\it the only one} to fit the data with high statistical significance. The combination of an acceptable fitting and of a realistic column depth provides a scenario globally self-consistent and consistent with the data.  Several considerations are now in order:  \\begin{itemize}  \\item Our best model is a loop heated at the apex, with a relatively high (still reasonable) temperature, a pressure of 0.3 dyne cm$^{-2}$ at the footpoints and a density decreasing to $\\sim 3 \\times 10^7$ cm$^{-3}$ close to the apex. Such pressure and density values are not unusual for large scale structures (e.g. Vaiana \\& Rosner 1978). A loop heating localized in corona is not a new result (e.g. Reale et al. 2000).  \\item The best fitting results have been obtained after subtraction of an unstructured background emission, extrapolated directly from the data, in the hypothesis that it is due mostly to chance alignment of disordered faint loops and/or of a streamer along the line of sight, and therefore not to be modelled in hydrostatic equilibrium.  This is one possible description of the scenario, not necessarily the best one, but, in our opinion, it is at least as realistic as other more local or model-dependent estimations of the background.  \\item High data statistics, obtained in this case by averaging over tens of exposures, is important to discriminate among loop models with different temperature and heating location.  Fig.~\\ref{fig:data_models} shows that larger error bars would made such discrimination much harder. This shows, once again, that the task of fitting data with loop models is not at all trivial, and requires both detailed modeling and high quality data.  \\end{itemize}  Of course, this work suffers from limitations. The background evaluation may not be unique, and a single loop model may be a simplified approximation to describe part of an arcade.  Also the exploration of the loop parameter space is forcedly limited, and we cannot exclude that a more refined tuning of the free parameters, and including other effects, such as strong geometrical variations at the base of the loop, could bring to equally good or better fitting results.  However, our analysis and modeling lead to a statistically acceptable description of the data, and to physically sound results (loop heating and plasma conditions, column depth), and appears therefore to be adequate to the data quality.  This work shows that the analyzed observational data can be described with a conventional hydrostatic and non-isothermal loop models and with an unstructured vertically stratified spurious emission.  The deposition of heat at the apex of large scale structures is a different result from those of both previous works, and may not be in agreement with recent results based on the analysis of data from the TRACE mission (Aschwanden et al. 2001).  Since there are serious independent arguments which make the analysis and temperature diagnostics with narrow band EUV instruments a very delicate issue (Schmelz et al. 2001), we believe that, at variance with Aschwanden (2001), the problems of the diagnostics of the heating function in corona and of the coherent interpretation of observations with different telescopes (e.g. Yohkoh and TRACE) are both still open.  We do not pretend to put a conclusive word on this topic, nor claim that our results are necessarily the best ones, rather we want to stimulate the community toward a more careful attention to aspects of the data analysis, such as the evaluation of background emission, important for a correct interpretation and modeling of the data. Other authors (McKay et al. 2000) come to similar conclusions, although with a totally different analysis. The present work also emphasizes the importance of the selection of the loops to be analyzed: loops embedded in crowded and complex coronal regions may not be the best ones, because emission from structures of different kind and characteristics may intersect along the line of sight, which may be difficult to evaluate. Further extensive and systematic analysis and modeling of bright, and as far as possible isolated, loop structures may provide more constraints on the issue of coronal heating deposition.  \\bigskip \\bigskip"
    },
    "0207/astro-ph0207021_arXiv.txt": {
        "abstract": "{ We are investigating the kinematics of jets in active galactic nuclei on parsec scales by studying a representative population of sources.  This study is being carried out using the Very Long Baseline Array at 15\\,GHz, with more than 800 images taken since 1994. In this contribution we present an overview of the diversity of kinematics for a complete sample of sources. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207217_arXiv.txt": {
        "abstract": "{Based on a sample of  1114 flares observed simultaneously in hard X-rays (HXR) by the BATSE instrument and in soft X-rays (SXR) by GOES, we studied several aspects of the Neupert effect and its interpretation in the frame of the electron-beam-driven evaporation model. In particular, we investigated the time differences ($\\Delta t$) between the maximum of the SXR emission and the end of the HXR emission, which are expected to occur at almost the same time. Furthermore, we performed a detailed analysis of the SXR peak flux -- HXR fluence relationship for the complete set of events, as well as separately for subsets of events which are likely compatible/incompatible with the timing expectations of the Neupert effect. The distribution of the time differences reveals a pronounced peak at $\\Delta t = 0$. About half of the events show a timing behavior which can be considered to be consistent with the expectations from the Neupert effect. For these events, a high correlation between the SXR peak flux and the HXR fluence is obtained, indicative of electron-beam-driven evaporation. However, there is also a significant fraction of flares (about one fourth), which show strong deviations from $\\Delta t = 0$, with a prolonged increase of the SXR emission distinctly beyond the end of the  HXR emission. These results suggest that electron-beam-driven evaporation plays an important role in solar flares. Yet, in a significant fraction of events, there is also clear evidence for the presence of an additional energy transport mechanism other than nonthermal electron beams, where the relative contribution is found to vary with the flare importance. ",
        "introduction": "Observations of solar flares in X-rays and microwaves frequently show that the shape of the rising part of the soft X-ray light curve closely resembles the time integral of the microwave or hard X-ray light curve. This led to the idea that there is a causal relationship between the nonthermal (microwave and hard X-ray) and thermal (soft X-ray) emission of a flare (Neupert 1968; Dennis \\& Zarro 1993), which has become known as the Neupert effect.  It has been shown that this effect can be reproduced by a model, in which the flare energy is released primarily in the form of nonthermal electrons (e.g., Brown 1971; Li et al. 1993). According to the so-called thick-target model, the hard X-ray (HXR) emission is electron-ion bremsstrahlung produced by electron beams encountering the dense layers of the lower corona, transition region, and chromosphere. The model assumes that only a small fraction of the energy of the nonthermal electrons is lost through radiation (for a discussion see McDonald et al. 1999). Most of the energy is transferred to heating of the ambient thick-target plasma via Coulomb collisions between the beam and the ambient electrons. Due to the rapid deposition of energy by the accelerated electrons, the energy cannot be radiated away at a sufficiently high rate and strong pressure gradients develop. The pre-flare hydrostatic equilibrium is lost and the heated plasma explosively expands up into the corona in a process known as chromospheric evaporation (e.g., Antonucci et al. 1984; Fisher et al. 1985; see also the review by Antonucci et al. 1999, and references therein). The hot dense plasma that has been convected into the corona gives rise to enhanced soft X-ray (SXR) emission via thermal bremsstrahlung.  Under such circumstances, the hard X-ray emission is directly related to the electron beam flux. On the other hand, the soft X-ray emission should be directly linked to the energy deposited by the same electrons up to a given time, i.e. to the time integral of the electron beam flux, and we can expect to see the Neupert effect. Thus, if the Neupert effect is observed, this can be considered as \\mbox{evidence} of electron-beam-driven chromospheric evaporation (see McTiernan et al. 1999). In recent years, the Neupert effect has also been observed on stars, which suggests the existence of the chromospheric evaporation process also in stellar flares (Hawley et al. 1995; G\\\"udel et al. 1996).  In the present study we utilize statistical properties of solar flares observed simultaneously in SXR and HXR emission to test several expectations from the Neupert effect. The main predictions are: (1) The maximum of the SXR emission and the end of the HXR emission should occur at the same time. (2) There should be a high correlation between the HXR fluence, i.e. the HXR flux integrated over the event duration, and the SXR peak flux. The correlation of the HXR fluence and the SXR peak flux and its relation to the involved nonthermal and thermal energies provide the fundamental link between the Neupert effect and the electron-beam-driven chromospheric evaporation model, which will be discussed in Sect.~\\ref{SectAnalysis}. The present analysis can be considered as a complementary approach to studies of the Neupert effect which use the actual SXR and HXR light curves (see, e.g., Dennis et al. 1992; Dennis \\& Zarro 1993; McTiernan et al. 1999). By accessing only statistical flare quantities, such as HXR end time, SXR maximum time, HXR fluence and SXR peak flux, we neglect part of the information contained in the light curves. However, such a statistical approach has the advantage that it can be applied to a large data set, i.e. it is not restricted to a selected sample of events, which intrinsically favors the analysis of large flares. Moreover, it enables us to define and investigate subsets of events, still representing statistically meaningful data sets.  The paper is structured in the following way. Section~2 contains a description of the soft X-ray and hard X-ray data used in the analysis together with the method of finding corresponding SXR/HXR events. In Section~3, we discuss, in the frame of the electron-beam-driven chromospheric evaporation model, the relationship between SXR and HXR emissions, and the associated thermal/nonthermal energies. In this respect, it is essential to clarify the question {\\it if\\,} and {\\it in which formulation} the Neupert effect is valid for the bulk of solar flares. In Section~4, our results are presented, comprising an investigation of the relative timing of the SXR and HXR emission as well as a detailed study of the HXR fluence~-- SXR peak flux relationship. The results are interpreted and discussed in Section~5, and the conclusions are drawn in Section~6. ",
        "conclusions": "In the following we briefly summarize the basic results of the analysis: \\begin{enumerate} \\item The distribution of the differences of the SXR peak times and HXR end times is strongly peaked at $\\Delta t = 0$. \\item Yet, a significant fraction of events ($\\sim$25\\%) shows strong deviations from the chosen timing criterion applied as a necessary condition for consistency with the Neupert effect (set~2). These events are characterized by increasing SXR emission beyond the end of the HXR emission. \\item Flares that satisfy the timing criterion (set~1), embracing about one half of the events, reveal a much higher correlation between the HXR fluence and the SXR peak flux than those of set~2. The strong correlations found for set~1 suggests that electron-beam-driven chromospheric evaporation plays an important role for these events. \\item Set~2 contains many fewer large events than set~1. For weak flares, on average, the events of set~2 have higher SXR peak fluxes at a given HXR fluence than those of set~1, suggesting that an additional energy  transport mechanism other than the HXR emitting electrons contributes to the SXR emission. \\item Events with negative $\\Delta t$, i.e. the SXR peak occurs before the end of the HXR emission, preferentially belong to set~1 and are of long duration. These events are compatible with the electron-beam-driven evaporation model. In the decay phase of long-duration flares, the instantaneous cooling of the hot plasma is likely to dominate over the evaporation-driven energy supply (Li et al. 1993). From the present data set we infer that, on average, this phase covers $\\sim$0.4~times of the HXR event duration. \\item For the events of set~1, the SXR peak flux -- HXR fluence relationship is not linear. However, for large HXR fluences the SXR peak flux -- HXR fluence relationship tends towards a linear function. Correspondingly, for the events of set~1, the factor $k$ is a decreasing function of the HXR fluence. Yet, for large HXR fluences, $k$ becomes approximately constant. \\item Although high correlations are found among the SXR peak fluxes and HXR fluences, the scatter in the SXR peak flux versus HXR fluence plot is larger than an order of magnitude (up to two orders of magnitude). \\end{enumerate} Finally, we stress that although the results presented show that in a statistical sense about half of the events show characteristics compatible with the Neupert effect, the scatter of the SXR peak flux versus HXR fluence indicates that a wide range of physical conditions are met in solar flares. The main outcomes of the analysis can be interpreted in the sense that the process of electron-beam-driven evaporation plays an important role in solar flares. On the other hand, the prolonged SXR emission found in a significant fraction of events also gives strong indications for the presence of an additional energy transport mechanism, probably thermal conduction, whereas the relative contribution of the different transport mechanisms shows a dependence on the flare importance. The energy provided by the additional agent may play a prominent role in weak flares, whereas in intense events its contribution is much less important than the electron-beam-driven component."
    },
    "0207/astro-ph0207167_arXiv.txt": {
        "abstract": "{Kink modes of solar coronal structures, perturbing the loop in the direction along the line-of-sight (LOS), can be observed as emission intensity disturbances propagating along the loop provided the angle between the LOS and the structure is not ninety degrees. The  effect is based upon the change of the column depth of the loop (along the LOS) by the wave. The observed amplitude of the emission intensity variations can be larger than the actual amplitude of the wave by a factor of two and there is an optimal angle maximizing the observed amplitude. For other angles this effect can also attenuate the observed wave amplitude. The observed amplitude depends upon the ratio of the wave length of kink perturbations to the width of the structure and on the angle between the LOS and the axis of the structure.  Sausage modes are always affected negatively from the observational point of view, as the observed amplitude is always less than the actual one. This  effect should be taken into account in the interpretation of wave phenomena observed in the corona with space-borne and ground-based imaging telescopes. ",
        "introduction": "In last few years, significant progress in the observational study of MHD wave activity of the solar corona has been achieved with SOHO/EIT and TRACE EUV imaging telescopes. Flare-generated decaying oscillations of coronal loops have been observed and interpreted as kink fast magnetoacoustic modes of the loops (Aschwanden et al. 1999; Nakariakov et al. 1999; Schrijver \\& Brown 2000; Aschwanden et al. 2002). Fast magnetoacoustic waves are possibly responsible for events such as coronal Moreton (or EIT) waves (Thompson et al. 1998, Ofman \\& Thompson 2002). Slow magnetoacoustic waves have been discovered in polar plumes (Ofman et al. 1997; DeForest \\& Gurman 1998; Ofman, Nakariakov \\& DeForest 1999) and in long loops (Berghmans \\& Clette 1999; De Moortel, Ireland, Walsh 2000; Nakariakov et al. 2000; De Moortel 2002 ). These observational breakthroughs give rise to the use of MHD coronal seismology (Nakariakov et al. 1999; Robbrecht et al. 2001; Nakariakov \\& Ofman 2001) and were interpreted to support the idea of wave-based theories of coronal heating (e.g., Tsiklauri \\& Nakariakov 2001), and the solar wind acceleration (e.g., Ofman, Nakariakov \\& Seghal 2000).  Slow and fast magnetoacoustic waves are compressive and cause perturbations of plasma density. As an emission depends upon the density, the waves can be detected as the emission variations by imaging telescopes. An important characteristic of the phenomenon is the angle between the direction of the wave propagation and the line of sight (LOS). Imaging telescopes allow us to observe magneto\\-acoustic waves propagating at sufficiently large angles to the LOS. In particular, this fact motivated the interpretation of the propagating EUV emission disturbances as  slow magnetoacoustic waves (see the references above). In addition, Alfv\\'en waves, which are linearly incompressible, as well as almost incompressible kink modes of coronal magnetic structures (e.g., Roberts 2000 and references therein), can also be detected with an imaging telescope with sufficient spatial and temporal resolution, if perturbations of the magnetic field have a component {\\it perpendicular} to the LOS. Indeed, as the magnetic field is frozen into the coronal plasma, the perpendicular displacement of the field can be highlighted by variation of emission intensity.  In this paper we discuss an alternative method for the observational detection of kink modes of coronal magnetic structures oscillating in the plane {\\it containing} the LOS. It is shown that this would lead to modulation of the intensity of the emission {\\it along the axis of the structure}, produced by the change of the LOS column depth of the loop. We also demonstrate that this phenomenon is important for sausage modes. ",
        "conclusions": "The phenomenon of the modulation of the emission intensity by kink modes polarized in the plane formed by the loop axis and the LOS provides  a possibility for the observational detection of the kink modes in coronal structures. According to the discussion above, the LOS effect {\\it at an optimal angle, $\\theta_{max}$,} can amplify the kink perturbations by a factor of 2 (cf.  Fig.~4). For example, if the boundary perturbation is produced by a kink mode of a relatively modest amplitude of about 5\\%, which corresponds to the typical coronal wave amplitudes detected by SOHO/EIT (DeForest \\& Gurman 1998) and by TRACE (e.g. De Moortel et al. 2002 and references therein), the observed perturbation of the intensity produced by the kink wave can reach 10\\%, which would make the wave easily observable. In contrast, the optimal observation angle for the sausage modes is simply $90\\degr $ (cf.  Fig.~8). Consulting figures \\ref{on_k} and \\ref{on_theta} we can see parameter regimes for both amplification and attenuation of the kink mode observations. Figures \\ref{on_kSaus} and \\ref{on_thetaSaus} demonstrate that relative to $90\\degr $ observations there is only attenuation of sausage modes. Additional observability constraints are connected with the wave period and length.  In the case of EUV imaging coronal telescopes, such as EIT and TRACE, the observability of the waves is limited by the telescope time resolution. For example, taking the kink speed inside a loop to be 1000~km/s, which corresponds to the estimations in Nakariakov et al. 1999, Nakariakov \\& Ofman 2001, the time resolution of about 30~s does not allow us to observe wavelengths shorter than 30~Mm.  In the case of ground-based observations, when the loop is observed in the green line bandpass, the observability is limited by the spatial resolution of the telescope, which is usually over a few arcsec, as the time resolution of such observations is usually less then 1 sec (e.g, the cadence time of SECIS is 2.25$\\times10^{-2}$~s and the pixel size is 4.07~arcsec, Williams et al. 2001). In particular, this effect can be responsible for propagating 6~second disturbances of the Fe\\,{\\sc xiv} green line, discovered by the stroboscopic method in a solar eclipse data and interpreted as fast magnetoacoustic waves (Williams et al. 2002, see also Williams et al. 2001). Indeed, estimating the wavelength of the perturbations at about 18~Mm (for the wave period of about 6~s and the propagation speed of about 2~Mm/s), we conclude that the waves are detectable with the time and spatial resolution of the telescope. For loop widths of about 3-6~Mm, the normalized wavelength is about 3-6, which gives the wavenumber $k$ of about 1-2. According to Fig.~3, the LOS effect discussed in this paper could amplify the observed amplitude of the wave, optimizing the observability of the modes.  According to figure \\ref{theta_k:fig}, for the parameters discussed, the optimal angle would be about $20\\degr -40\\degr $. This makes the phenomenon discussed here relevant to the interpretation of Williams et al. (2002)'s results. However,  the confident interpretation of the propagating disturbances in terms of the kink waves requires detailed comparison of the observed results and theoretical predictions. Also, interpretation of these observations in terms of propagating kink modes should be tested against another possible interpretation of the waves as sausage modes.  The study of the possible relevance of the LOS effect to the interpretation of short wavelength propagating waves observed in the corona by Williams et al. (2002) is now in progress and will be presented elsewhere.  Also, this phenomenon should be taken into account in the analysis of other examples of coronal wave activity, in particular the slow waves in loops and polar plumes, discussed in Introduction. However, direct application of the results presented in this paper to the coronal slow magnetoacoustic waves is not possible as the slow waves are essentially compressible, and the density variation should be taken into account in Eq.~(\\ref{inten}). This phenomenon should also be taken into account in interpretation of intensity oscillations observed in prominence fine structures (e.g. Joarder, Nakariakov \\& Roberts 1997; Diaz et al. 2001). This suggests another possible  development of this study."
    },
    "0207/astro-ph0207398_arXiv.txt": {
        "abstract": "We report the correct classification of an overlooked Fanaroff-Riley class II radio-loud quasar with broad absorption lines, only the second such object so identified.  The rare properties of this quasar, LBQS 1138$-$0126, are twice overlooked. First LBQS 1138$-$0126 was found in the Large Bright Quasar Survey but only noted as a possible broad absorption line quasar without additional follow-up. Later LBQS 1138$-$0126 was rediscovered and classified as a radio-loud broad absorption line quasar but not recognized as an FR II radio source. We describe the radio, absorption line, and optical polarization properties of LBQS 1138$-$0126 and place it in context with respect to related quasars.  In particular, spectropolarimetry shows that LBQS~1138$-$0126 has high continuum polarization increasing from 3\\% in the red (rest-frame 2400 \\AA) to over 4\\% in the blue (rest-frame 1650 \\AA), essentially confirming the intrinsic nature of the absorption.  The polarization position angle rotates from $\\sim -30\\arcdeg$\\ in the red to $\\sim 0\\arcdeg$\\ in the blue; the radio lobe position angle is $\\sim 52\\arcdeg$\\ for comparison. LBQS 1138$-$0126 is additionally notable for being one of the most radio-loud broad absorption line quasars, and for having low-ionization broad absorption lines as well. ",
        "introduction": "Some 10-20\\% of optically selected quasars display broad, blueshifted ultraviolet absorption lines from highly ionized species (e.g., C IV -- HiBAL quasars), with a smaller fraction also showing absorption from very low-ionization species (e.g., Mg II -- LoBAL quasars). A decade ago, no formally radio-loud broad absorption line (BAL) quasars were known despite searching for such objects (e.g., Stocke et al. 1992).  This strong anti-correlation seemed to be a clue to explaining the apparent bimodal distribution of radio-loudness in quasars (which is very much less apparent in new samples selected using deep radio surveys, see e.g., White et al. 2000).  In the mid 1990s, with the advent of large area, deep radio surveys like the NVSS (Condon et al. 1998) and FIRST (Becker et al. 1995), radio-loud BAL quasars began to be found.   Becker et al. (1997) found the first such source, FIRST J1556+3517, by matching optically red point sources with FIRST radio sources; corrections for dust reddening and the possibility of beaming at radio frequencies leave this quasar straddling the traditional borderline between radio-loud and radio-quiet quasars (Najita, Dey, \\& Brotherton 2000). Additional follow-up to the FIRST survey found more radio-loud BAL quasars by matching radio sources to the APM catalog (Becker et al. 2000; Becker et al. 2001) and the Sloan Digital Sky Survey (Menou et al. 2001).  Over 50 BAL quasars have now been identified using radio-selection techniques.  The largest and brightest sample of radio-selected BAL quasars is that of Becker et al. (2000), which comprises $\\sim$27 BAL quasars from the FIRST Bright Quasar Survey (White et al. 2000).  Their radio properties can tell us some things that cannot be learned from optically selected BAL quasars. First, a wide range of radio spectral indices are present, including both flat and steep spectra; unified radio models (e.g., Orr \\& Browne 1982) would indicate that therefore a range of orientations are present.  Second, the radio sources are almost all compact (90\\%), whereas a matched parent population from the FIRST Bright Quasar Survey (FBQS) consists of only 60\\% compact sources. This is similarly contrary to the idea that BAL quasars are simply normal quasars seen edge-on.  Properties at other wavebands are also inconsistent with BAL quasars being edge-on quasars (e.g., Brandt \\& Gallagher 2000).  Were it not for these results, the discovery of a powerful edge-on Fanaroff-Riley type II (FR II; Fanaraoff \\& Riley 1974) radio source that is also BAL quasar, FIRST J101614.3+520916 (hereafter FIRST J1016+5209, Gregg et al. 2000), might have been taken as supporting evidence for an edge-on geometry.  The once popular notion of such an edge-on geometry was motivated by spectropolarimetry (e.g., Ogle et al. 1999) and by theoretical expectations for winds arising from accretion disks (e.g., Murray et al. 1995; Konigl \\& Kartje 1994).  Instead, Becker et al. (2000) and Gregg et al. (2000) interpret BAL quasar radio properties to indicate that BAL quasars are young, evolving quasars, moving from dusty objects with high covering fractions and high accretion rates to becoming more typical quasars.  Brotherton et al. (1998) found five radio-loud BAL quasars by matching color-selected quasar candidates from the 2dF Quasar Survey (e.g., Croom et al. 2001) against radio sources in the NVSS Survey. These were of particular interest because their blue colors ensured minimal dust-reddening which can artificially enhance the radio-loudness by making an object appear optically weak relative to its radio emission. Also of interest was that the observed targets were taken from a large optically selected sample, making possible statistical comparison with the Large Bright Quasar Survey (LBQS; Foltz et al. 1989; Hewett et al. 1991; Chaffee et al. 1991; Morris et al. 1991; Francis et al. 1992). The 2dF-NVSS cross-match indicates that approximately 1 in 500 bright blue quasars should be seen in an observed-frame optical spectrum as a radio-loud BAL quasar.  With the LBQS coming in at $\\sim$1000 quasars it would be likely to find 1 or 2 such objects, but not too surprising to not find any.  An error labeling the coordinates of one of the five radio-loud BAL quasars from Brotherton et al. (1998) prevented identification with the correct radio source (Brotherton et al. 2002).  Likewise, this error prevented identification with a previously discovered quasar LBQS 1138$-$0126 (Hewett et al. 1991), which had only been noted as being a possible BAL quasar but had not been followed-up in either the optical or radio. We show in this paper that LBQS 1138$-$0126 is not only a BAL quasar, it is also a powerful FR II radio source, only the second such BAL quasar found. ",
        "conclusions": "It is clear that LBQS 1138$-$0126 is a powerful double-lobed FR~II radio source, and that the optical/UV spectrum contains both high-ionization and low-ionization BALs.  The high continuum polarization makes a very strong statistical argument that the absorption lines are broad and intrinsic and we are not being misled by the low-resolution of our spectra.  These properties lead to the conclusion that LBQS 1138$-$0126 is an FR II BAL quasar, only the second such object so identified after FIRST J1016+5209 (Gregg et al. 2000).  We also note the existence of the less luminous, low-redshift quasar PKS 1004+13 (Wills et al. 1999) that is a strong candidate for having BALs and is soon to have a definitive UV spectrum obtained with the Hubble Space Telescope. One other BAL quasar from Brotherton et al. (1998), UN J1053$-$0058, is seen to be a very core-dominated radio triple in FIRST Survey images.  The spatial exent of the triple is approximately 45\\arcsec, corresponding to a projected size of 550 kpc at $z=1.55$ for our adopted cosmology. Hutsem{\\'e}kers \\& Lamy, H. (2000) report an optical polarization of 1.89\\% with a position angle some 20-30\\arcdeg from being perpendicular to the radio axis.  The large radio core dominance (an observed core-to-extended 1.4GHz flux ratio of about 12) would indicate that the radio core, and the radio-loudness, is artificially enhanced by relativistic beaming.  Therefore UN J1053$-$0058 is perhaps best identified as a beamed radio-quiet quasar rather than as an intrinsically luminous radio-loud quasar (see Falcke, Patnaik, \\& Sherwood 1996; Falcke, Sherwood, \\& Patnaik 1996).  We focus the remaining discussion on quasars like LBQS 1138$-$0126 and FIRST J1016+5209.  How similar are LBQS 1138$-$0126 and FIRST J1016+5209? Table 1 compares their properties.  The similar small BIs, modest reddening, high polarization and radio luminosity in particular stand out.  Also similar is the fact that both have optical polarization position angles intermediate between those parallel or perpendicular to the large-scale radio structure.  The comparison of the angular size is strongly affected by adopted cosmology although both sources appear to be relatively large (at least several hundred kpc).  Significant differences include the presence of low-ionization BALs in LBQS 1138$-$0126 but not FIRST J1016+5209, and the smooth single trough structure of LBQS 1138$-$0126 compared to the jagged multiple trough structure in the BALs of FIRST J1016+5209 which also spans a much larger velocity range.  We note that among radio-quiet BAL quasars there exists a wide range of trough structures which have so far not been found to clearly correlate with other properties.  Gregg et al. (2000) proposed and discussed an evolutionary sequence for LoBAL quasars to evolve into HiBAL quasars and then radio-loud quasars with associated absorption that is the remnants of the once smooth BAL outflow. The formally radio-loud BAL quasars from Becker et al. (2000) with compact radio structures would be the youngest (most recently fueled) sources, perhaps themselves closely related to giga-Hertz peaked sources (GPS) and compact steep spectrum (CSS) radio-loud quasars (e.g., O'Dea 1998). Next in the progression is LBQS 1138$-$0126, in which the radio jet has escaped the nuclear regions but low-ionization BALs are still visible and dust is still present along the line of sight.  FIRST J1016+5209 is the next stage as the LoBAL features vanish and the HiBAL features begin to break up. The final stage before being regarded as a normal, unabsorbed radio-loud quasar is represented by PKS 1157+014, an object once put forward and rejected as a radio-loud BAL quasar for which variability has provided proof of the intrinsic nature of its rather narrow absorption features (Aldcroft, Bechtold, \\& Foltz 1998).  The polarization properties of LBQS 1138$-$0126 are similar to those of other highly polarized BAL quasars (e.g., Ogle et al. 1999).  The polarization mechanism is widely thought to arise from asymmetric scattering as other likely mechanisms can be ruled out with high-quality spectropolarimetry for individual objects.  In particular, the rise in the polarization toward shorter wavelengths is very common and can be attributed to a decreasing amount of dilution from direct light more reddened than the scattered light path or as a signature of dust scattering (e.g., Hines et al. 2001). Less commonly seen in polarized BAL quasars is a significant rotation of the polarization position angle, although it appear (e.g., QSO 2359$-$1241, Brotherton et al. 2001a).  Such a rotation can be ascribed to multiple scattering light paths with different amounts of reddening or different polarization efficiencies as a function of wavelength (e.g., in the case of dust scattering).  Such complex scattering geometries might be expected for young, highly accreting objects with large covering fractions, as opposed to a simpler scattering geometry as seen in edge-on Seyfert 2 and radio galaxies for which the polarization position angle is perpendicular to system (jet) axes (Antonucci 1993).  LBQS 1138$-$0126 is not especially red or faint and was easily found using optical selection techniques. The radio emissions are strong and detected in multiple surveys.  The key elements in making the classification of a radio-loud FR II BAL quasar are having an optical spectrum covering the C IV $\\lambda$1549 region and matching optical properties with radio properties. Given the statistics of Brotherton et al. (1998) and Menou et al. (2000) selecting BAL quasars by optical techniques and matching to radio surveys, the SDSS should turn up an entire population of quasars similar to LBQS 1138$-$0126.  The small number statistics and the possible existence of as yet unidentified biases make a precise estimate impossible, but we might expect 100-200 FR II BAL quasars to be identifiable from combining information from SDSS discovery spectra and the FIRST Survey.  This might seem like a large number of sources, but these still appear to be rare objects requiring very large surveys to discover in significant numbers. Such a sample would permit quantitative tests and development of the evolutionary hypothesis of Gregg et al. (2000)."
    },
    "0207/astro-ph0207437_arXiv.txt": {
        "abstract": "This paper is originally intended to give a comprehensive review of  the pulsar wind nebulae and magnetosphere, but it has been moved to a poster paper so that we have changed the aim of the paper and focused on the Crab Nebula problem to suggest that particle acceleration takes place not only at the inner shock but also over a larger region in the nebula with disordered magnetic field.  Kennel and Cornoniti (1984) constructed a spherically symmetric model of the Crab Nebula and concluded that the pulsar wind which excites the nebular is kinetic-energy dominant (KED) because the nebula flow induced by KED wind is favorable to explain the nebula spectrum and expansion speed. This is true even with new Chandra observation, which provides newly the spatially resolved spectra. We have shown below with 3D modelling and the Chandra image that pure toroidal magnetic field and KED wind are incompatible with the Chandra observation. ",
        "introduction": "KC model (Kennel and Coroniti 1984) assumes that a super fast MHD wind from the central pulsar terminates at a shock, and the shocked wind radiates in synchrotron radiation, which is observed as the nebula. The central cavity is identified as the wind region.  The pulsar wind is originally Poynting energy dominant deep in the pulsar magnetosphere, but by MHD acceleration the energy is converted into kinetic energy of the plasma outflow. How efficient the acceleration is a main problem in the theory of relativistic centrifugal wind. It is known that this problem is coupled with the jet-disc formation, which is clearly observed with Chandra (Fig. 4).  The acceleration efficiency is parameterized by so-called $\\sigma$-parameter which is the ratio of the Poynting energy to kinetic energy of the terminal flow just before the shock. KC find that $\\sigma$ determins the expansion speed of the nebula and in turn spectrum of the nebula. It was a great success that KC model reproduces the nebula spectrum (Fig. \\ref{fig:1}). KC found that $\\sigma = 0.003$ (kinetic energy dominant) and the Lorentz factor of the wind is $3 \\times 10^6$.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=60mm]{sshibatafig1.eps} \\end{center} \\caption{KC model well expains the over all spectrum of the nebula (after Atoyan \\& Aharonian, 1996).} \\label{fig:1} \\end{figure}  On the other hand, no wind theory explains such a high acceleration. It must be noted that the observed image suggest that the wind has jet-disc structure, which is not explained, either.  The aim of this paper is to apply the picture of KED wind to the new Chandra data and examine whether the model still describes the nebula well or it needs some modification. ",
        "conclusions": ""
    },
    "0207/astro-ph0207601_arXiv.txt": {
        "abstract": "We present optical spectra and photometry sampling the first six months after discovery of supernova (SN) 1999gi in NGC~3184.  SN~1999gi is shown to be a Type II-plateau event with a photometric plateau lasting until about 100 days after discovery.  The reddening values resulting from five independent techniques are all consistent with an upper bound of $\\Ebv < 0.45$ mag established by comparing the early-time color of SN~1999gi with that of an infinitely hot blackbody, and yield a probable reddening of $\\Ebv = 0.21 \\pm 0.09$ mag.  Using the expanding photosphere method (EPM), we derive a distance to SN~1999gi of $11.1^{+2.0}_{-1.8}$ Mpc and an explosion date of 1999 December $5.8^{+3.0}_{-3.1}$, or $4.1^{+3.0}_{-3.1}$ days prior to discovery.  This distance is consistent with a recent Tully-Fisher distance derived to NGC~3184 ($D \\approx 11.59$ Mpc), but is somewhat closer than the Cepheid distances derived to two galaxies that have generally been assumed to be members of a small group containing NGC~3184 (NGC 3319, $D = 13.30 \\pm 0.55$ Mpc, and NGC 3198, $D = 13.80 \\pm 0.51$ Mpc).  We reconsider the upper mass limit ($9^{+3}_{-2} {\\rm \\ M}_{\\odot}$) recently placed on the progenitor star of SN~1999gi by Smartt et al. (2001, 2002) in light of these results.  Following the same procedures, but using the new data presented here, we arrive at a less restrictive upper mass limit of $15^{+5}_{-3} {\\rm \\ M}_{\\odot}$ for the progenitor.  The increased upper limit results mainly from the larger distance derived through the EPM than was assumed by the Smartt et al. analyses, which relied on less precise (and less recent) distance measurements to NGC~3184.  Finally, we confirm the existence of ``complicated'' P-Cygni line profiles in early-time and later photospheric-phase spectra of SN~1999gi.  These features, first identified by Baron et al. (2000) in spectra of SN~1999em as high-velocity absorptions in addition to the ``normal'' lower-velocity component, are here verified to be true P-Cygni profiles consisting of both an absorption trough and an emission peak at early times.  In the earliest spectrum, taken less than a day after discovery, the features extend out to nearly $-30,000$ \\kms, indicating the existence of very high-velocity material in the outer envelope of SN~1999gi. ",
        "introduction": "\\label{sec:introduction}  Supernova (SN) 1999gi was discovered by Nakano et al. (1999) on 1999 December 9.82 (UT dates are used throughout this paper) at an unfiltered magnitude of $m \\approx 14.5$ in the nearly face-on ($i < 24^{\\circ}$, from the Lyon-Meudon Extragalactic Database [LEDA\\footnote{\\url{http://leda.univ-lyon1.fr}}]) SBc galaxy NGC~3184.  The identification of hydrogen in an early-time spectrum quickly defined it as a Type II event (Nakano et al. 1999; see Filippenko 1997 for a review of SN types), and the absence of the SN on CCD images of the same field taken 6.64 and 7.32 days earlier (Trondal et al. 1999, limiting unfiltered magnitude 18.5, and Nakano et al. 1999, limiting unfiltered magnitude 19.0, respectively) implies that it was discovered shortly after explosion.  There have been two previous investigations of SN~1999gi.  In the first, Leonard \\& Filippenko (2001) examine a single epoch of optical spectropolarimetry of SN~1999gi taken 107 days after discovery. They find an extraordinarily high degree of linear polarization, $p_{\\rm max} = 5.8\\%$, where $p_{\\rm max}$ is the highest level of polarization observed in the optical bandpass.  If intrinsic to SN~1999gi, such polarization implies an enormous departure from spherical symmetry (H\\\"{o}flich 1991).  However, Leonard \\& Filippenko (2001) conclude that the majority of the polarization is likely due to interstellar dust, and is not intrinsic to the SN.  From photometry reported in various IAU Circulars, Leonard \\& Filippenko (2001) tentatively classify SN~1999gi as a Type II-plateau supernova (SN II-P).  In addition, the total flux spectrum of SN~1999gi from day 107 and a comparable spectral epoch of SN~1999em, a classic SN II-P (Leonard et al. 2002 [hereafter L02]; Hamuy et al. 2001), show great spectral similarity, suggesting that they may have been quite similar events.  In the second study, Smartt et al. (2001; hereafter S01) examine pre-explosion archival {\\it Hubble Space Telescope (HST)} images of NGC~3184 and use the lack of a progenitor-star detection in the pre-discovery frame, along with an estimated distance of $D = 7.9$ Mpc, to set an upper limit on the absolute magnitude of the progenitor for SN~1999gi.  This is then translated into an upper mass limit of $9^{+3}_{-2}\\ {\\rm M}_{\\odot}$ for the progenitor of SN~1999gi through comparison with stellar evolution models.  A subsequent reanalysis of the same data by Smartt et al. (2002; hereafter S02), using improved models, confirms this limit.  Since stars with initial mass $\\lesssim 8\\ {\\rm M}_{\\odot}$ are not expected to undergo core collapse (e.g., Woosley \\& Weaver 1986, and references therein), this upper bound sets very tight constraints on the possible mass of the progenitor, a fact that has important implications for the nature of the progenitors of SNe II-P. Indeed, other than the Smartt et al. studies, progenitor masses for SNe II-P are virtually unconstrained by direct observation.\\footnote{Mass constraints on the progenitors of other types of core-collapse SNe have been obtained through studies of their environments; see, e.g., Van Dyk et al. (1999b).}  Although the progenitors of SN 1961V (Goodrich et al. 1989; Filippenko et al. 1995; Van Dyk, Filippenko, \\& Li 2002), SN 1978K (Ryder et al. 1993), SN 1987A (e.g., White \\& Malin 1987; Walborn et al. 1987), SN 1993J (Filippenko 1993; Aldering, Humphreys, \\& Richmond 1994; Cohen, Darling, \\& Porter 1995), and SN 1997bs (Van Dyk et al. 1999a) have been identified, all of these SNe II were peculiar.  In this paper, we present 15 optical spectra and 30 photometric epochs of SN~1999gi sampling the first 169 and 174 days since its discovery, respectively, and derive its distance through the expanding photosphere method (EPM).  We present and discuss our photometric and spectroscopic observations in \\S~\\ref{sec:photometry} and \\S~\\ref{sec:spectra}, respectively.  We estimate the reddening of SN~1999gi from a variety of techniques in \\S~\\ref{sec:reddening}.  We apply the EPM to SN~1999gi in \\S~\\ref{sec:sn1999giepmdistance}, and compare the derived distance to existing estimates of the distance to NGC~3184 in \\S~\\ref{sec:epmdistance}.  In \\S~\\ref{sec:progenitormass} we discuss the impact of our results on the progenitor mass limits previously determined by S01 and S02.  We summarize our main conclusions in \\S~\\ref{sec:conclusions}.  Note that much of the background material for the data and analysis presented in this paper, including the details of the photometric and spectral reductions as well as many specifics of the EPM technique itself, is thoroughly covered by earlier studies and therefore not repeated here.  In particular, the recent analysis by L02 of SN~1999em, a strikingly similar event to SN~1999gi, is frequently referenced. ",
        "conclusions": "\\label{sec:conclusions}  We present 15 optical spectra and 30 photometric epochs of SN~1999gi sampling the first 169 and 174 days since discovery, respectively, and derive its EPM distance.  Our main conclusions are as follows.  \\begin{enumerate}  \\item SN~1999gi is a Type II-P event with a photometric plateau lasting until about 100 days after discovery.  It reached $B$ maximum on 1999 December $13.7 \\pm 1.8$ (HJD 2,451,526.2 $\\pm 1.8$, or $3.9 \\pm 1.8$ days after discovery), and achieved peak $B$ and $V$ magnitudes of $\\sim 14.8$ and $\\sim 14.6$, respectively.  Overall, we find the photometric behavior of SN~1999gi to be extremely similar to that of SN~1999em.  \\item The very early-time spectra of SN~1999gi confirm the existence of the high-velocity absorption features in the profiles of \\hbeta\\ and \\ion{He}{1} $\\lambda 5876$ that were first identified by Baron et al. (2000) in spectra of SN~1999em.  The highest-velocity feature (\\hbeta, day 1) extends out to nearly $-30,000$ \\kms, implying the existence of very high velocity material in the outer envelope of SN~1999gi at early times. These features are verified to be true P-Cygni profiles, consisting of both an absorption trough and an emission peak in early-time spectra.  The high-velocity features are not seen, however, in \\halpha\\ at early times.  \\item By comparing the early-time spectral shape with blackbody functions we derive an upper limit on the reddening of SN~1999gi of $\\Ebv < 0.45$ mag; comparison with the color evolution of SN~1999em suggests an even lower limit, of $\\Ebv < 0.30$ mag.  Other reddening estimates are consistent with these limits, and imply a somewhat lower reddening, $\\Ebv = 0.21 \\pm 0.09$ mag, which we adopt as the preferred reddening value.   \\item Our best estimate for the EPM distance and explosion time of SN~1999gi is $D = 11.1^{+2.0}_{-1.8}$ Mpc and $t_{\\circ} = 4.1^{+3.0}_{-3.1}$ days prior to discovery.  This distance is consistent with some recent distance estimates to NGC~3184, but is $\\sim 20 \\%$ shorter than the average of the Cepheid distances derived to two putative group-member galaxies.  \\item The EPM distance implies an average plateau brightness of $\\overline{M}_V~{\\rm (plateau)} = -16.0 \\pm 0.4$ mag, which is very similar to the value found for SN~1999em and consistent with the average plateau brightness found by L02 of $\\overline{M}_V~{\\rm (plateau)} = -16.4\\pm{0.6}$ mag for 9 photometrically confirmed SNe II-P with EPM distances.  \\item Following the analysis methods of S01 and S02 we derive a new upper mass limit for the progenitor of SN~1999gi of $15^{+5}_{-3}\\ {\\rm M}_{\\odot}$, which is substantially less restrictive than the original limit of $9^{+3}_{-2}\\ {\\rm M}_{\\odot}$ found by S01 and S02.  The higher limit comes mainly from the longer distance derived through the EPM, than was assumed by the earlier analyses.  \\end{enumerate}"
    },
    "0207/astro-ph0207092_arXiv.txt": {
        "abstract": "We present the result of a Chandra ACIS observation of the pulsar PSR~B1853+01 and its associated pulsar wind nebula (PWN), embedded within the supernova remnant W44. A hard band ACIS map cleanly distinguishes the PWN from the thermal emission of W44. The nebula is extended in the north-south direction, with an extent about half that of the radio emission. Morphological differences between the X-ray and radio images are apparent. Spectral fitting reveals a clear difference in spectral index between the hard emission from PSR~B1853+01 ($\\Gamma \\sim$ 1.4) and the extended nebula ($\\Gamma \\sim$ 2.2). The more accurate values for the X-ray flux and spectral index are used refine estimates for PWN parameters, including magnetic field strength, the average Lorentz factor $\\gamma$ of the particles in the wind, the magnetization parameter $\\sigma$, and the ratio k of electrons to other particles. ",
        "introduction": "The remarkable pulsar wind nebula associated with PSR~B1853+01 and embedded in the evolved, mixed-morphology supernova remnant W44 is of particular interest for several reasons. It is the oldest known pulsar wind nebula (PWN), associated with an active pulsar (only the PWN in IC~443 is thought to be older -- Bocchino \\& Bychkov 2001).  Its age of 20,000 yr is estimated from the PSR~B1853+01 spindown (Wolszczan, Cordes \\& Dewey 1991). At the same time, PSR~B1853+01 is one of the 10 youngest known pulsars. Thus the nebula may allow the testing of hypotheses regarding PWN evolution, and serves as a bridge between the young, active pulsars in supernova remnants and the preponderant population of isolated, old radio pulsars. Additionally, because of its high proper motion, the pulsar leaves a record of its evolution embedded in the extended nebula. A combination of high resolution radio and X-ray observations can potentially disentangle this record. For example, radio measurements of the extent of the nebula perpendicular to the direction of motion provide means for setting an upper limit to the lifetime of the radio emitting electrons ($\\sim$ 6,000 yr) unavailable from observation of wind nebulae associated with stationary pulsars.  Finally, some of the properties of this PWN are similar enough to those of the much younger and more luminous Crab Nebula (Chevalier 2000) to invite comparisons and speculation regarding the reason for the similarities.  With the clearest view of the X-ray universe, especially above 3 keV, now available via the Chandra X-ray Observatory, more comprehensive studies of embedded PWNs become feasible. In this paper we use Chandra's Advanced CCD Imaging Spectrometer (ACIS) to reveal the X-ray structure and spectrum of the synchrotron nebula surrounding the W44 pulsar, and provide a more careful look at the X-ray spectrum and its spatial variation. As has been shown in numerous other works (e.g., Harrus et al. 1996; Frail et al. 1996; Torii et al. 2000), the study of the wind nebulae surrounding pulsars provides a means for understanding the energetics of pulsars, and in particular how they transfer their rotational spindown energy into a relativistic wind.  W44 is one of the first remnants for which hard band X-ray imaging was used to isolate a pulsar and its associated PWN from the brighter, softer thermal X-ray emission associated with the remnant's shock-heated gas. Neither the pulsar nor the nebula is apparent in low energy X-ray images, such as that from the ROSAT PSPC (Rho et al. 1994). Using ASCA and its broader band imaging, however, Harrus et al. (1996) showed that while the PWN is invisible in the broad band image, it becomes the dominant feature above 4 keV. The centroid of the X-ray emission is consistent with the location of the pulsar, PSR~B1853+01 (Wolszczan, Cordes \\& Dewey 1991). ASCA's modest angular resolution precluded spatially distinguishing the synchrotron nebula from the surrounding diffuse emission, but Harrus et al. showed that the spectrum of the region including the PWN has a hard continuum component not detected elsewhere in W44. The techniques pioneered in Harrus et al. (1996) have subsequently been used to identify stellar remnants or synchrotron nebulae in other remnants (e.g., IC~443 - Keohane et al. 1997; MSH~15-56 - Plucinsky 1998; G292.0+1.8 - Torii, Tsunemi, \\& Slane 1998). The most important consequence of this approach is the dramatic increase in the number of supernova remnants with identified synchrotron nebulae and/or compact stellar remnants.  The discovery of the X-ray synchrotron nebula occurred contemporaneously with the mapping of the pulsar wind nebula in the radio (Frail et al. 1996). At 1.4 GHz it appears cometary in shape with an extent of $\\sim$2.5 arc minutes. The pulsar is located at the southern extremity. The radio surface brightness peaks at the widest part of the tail, approximately 1 arc minute north of the pulsar. Frail et al. interpret this structure as the result of the pulsar's motion through the interior of the remnant. Using three independent techniques, they derive a velocity of the pulsar of approximately 375 km/s. The radio emission has a nonthermal spectrum with a spectral index of -0.12$\\pm$0.04, and it is 17$\\pm$4 percent polarized. The spectral index distinguishes the PWN from the surrounding emission associated with W44 ($\\Gamma\\sim$-0.33); the spectrum and degree of polarization are similar to other pulsar wind nebulae. Using a combination of the X-ray and radio properties, Frail et al. estimate some key pulsar wind nebula parameters, including magnetic field strength, and the Lorentz factor $\\gamma$ of the electrons near the spectral break between the radio and X-ray slopes.  Giacani et al. (1997) presented radio and X-ray images of W44 as a whole. The PWN is apparent but inconspicuous in the radio. A line of H$\\alpha$ filaments lies along the eastern edge of the PWN, but there is no general correspondence with radio features, and it is unclear whether this emission is associated with the PWN or with shock heated material near the PWN only in projection.  The distance to W44 has generally been taken to be around 3 kpc, based on measurements of H~I absorption and 1720 MHz maser lines (Caswell et al. 1975; Claussen et al. 1996).  The analytical modeling of Cox et al. (1999) refines this distance to be between 2.5 and 2.6 kpc.  We use a distance of 2.6 kpc, and scale parameters in terms of d$_{2.6}$.  The implications of using this refined value are minor.  A luminosity estimate, for instance, is reduced by 25 percent, which is probably well within the uncertainty of the estimate. ",
        "conclusions": "The ACIS observation of W44 has revealed the structure and spectrum of the PWN associated with the W44 pulsar. We find as in previous work (Harrus et al. 1996) that the pulsar and PWN stand out at energies above $\\sim$2 keV. In the ACIS data they are detected for the first time at all energies down to the $\\sim$1 keV interstellar cutoff (but only because of our ability to discern the PWN's true shape from the hard band image). The PWN is clearly extended.  It is also highly asymmetric, with significantly greater extent to the north, opposite the pulsar's apparent projected direction of motion. The spectra of the pulsar and PWN can be described as power laws. There is no evidence for a $\\sim$10$^6$~K thermal component that might be associated with surface emission from the pulsar.  Such emission is presumably not detectable as a consequence of the background thermal emission from W44 and/or the high column density. The nebular spectrum is softer than the pulsar, and hints at steepening with distance from the pulsar.  The X-ray morphology of the PWN dif fers crucially from the radio morphology in two ways. First, it is a factor of two smaller in extent along the pulsar direction of motion. This difference can be ascribed to the overall softening of the electron spectrum as the higher energy electrons lose their energy more rapidly via synchrotron radiation. Second, the X-ray surface brightness peaks at or very near the pulsar, and decreases monotonically with distance from it. In contrast, the PWN has a more complex radio structure.  The pulsar is embedded in a neck of emission 10$\\arcsec$ wide that connects with the larger cometary nebula $\\sim$15$\\arcsec$ to the north. There is little diffuse radio emission surrounding the pulsar itself. A bridge of emission connects the pulsar and the PWN to its north. The radio surface brightness peaks not at the pulsar but near the northern edge of the X-ray nebula.  It is unusual to see such pronounced structural differences in a PWN.  More usually the X-ray and radio surface brightness profiles resemble each other and peak near the pulsar.  It is possible that low surface brightness structure associated with the PWN in both the radio and the X-ray is lost as a consequence of low contrast with the surrounding emission from the thermal remnant.  Is is surprising that any extent of the X-ray nebula is found. Based on synchrotron lifetime arguments, electrons sufficiently energetic to produce X-rays were expected only in close proximity to the pulsar (Harrus et al. 1996). In principle, the combination of the X-ray and radio profiles along the direction of the pulsar motion provides a record of the history of the production of energetic particles in the pulsar magnetosphere.  \\begin{deluxetable}{rccccc} \\tablecolumns{6} \\tabletypesize{\\footnotesize} \\tablecaption{Parameters for Various Pulsar Wind Nebulae} \\label{tab:pwn} \\tablewidth{0pt} \\tablehead{ \\colhead{} & \\colhead{3C 58} & \\colhead{Crab} & \\colhead{G21.5-0.9} & \\colhead{W44} & \\colhead{IC 443}} \\startdata X-ray Size (pc) & 10x6 & 1.2 & 7 & 1x0.5  & 3.5x2 \\\\ Radio Size (pc) & 10x6 & 3.5x2.3 & 2.2x1.3 & 2x1 & 1.3x0.9 \\\\ Distance (kpc) & 3.2 & 2 & 5 & 2.6 & 1.5 \\\\ Age (yr)       & 820 & 950 & 3-6000 & 20,000 & 30,000 \\\\ L$_x$ (ergs s$^{-1}$) & 2.4$\\times$10$^{34}$ &  2.1$\\times$10$^{37}$ & 3.3$\\times$10$^{35}$ &  6.0$\\times$10$^{32}$ & 2.6$\\times$10$^{33}$ \\\\ \\.{E} (ergs s$^{-1}$) & 4.0$\\times$10$^{36}$ &  4.7$\\times$10$^{38}$ & 3-6$\\times$10$^{37}$ &  4.3$\\times$10$^{35}$ & 1.3$\\times$10$^{36}$ \\\\ L$_x$/\\.{E} & 0.006 & 0.05 & 0.005-0.01& 0.001 & 0.002 \\\\ $\\Gamma$ range & 1.85, 2.4 & 1.6-2.3 & 1.5-2.8 & 2.1-2.3 & 1.6-2.3 \\\\ \\.{E} (ergs s$^{-1}$) & 4.0$\\times$10$^{36}$ &  4.7$\\times$10$^{38}$ & 3-6$\\times$10$^{37}$ &  4.3$\\times$10$^{35}$ & 1.3$\\times$10$^{36}$ \\\\ $\\sigma$ & 2-15$\\times$10$^{-3}$ & 3$\\times$10$^{-3}$ & 0.4-1.1$\\times$10$^{-3}$ & 0.4-1.0$\\times$10$^{-3}$ & \\\\ Cutoff frequency (Hz) & 5$\\times$10$^{10}$ & 1$\\times$10$^{13}$ & & 8$\\times$10$^{12}$ & 1$\\times$10$^{11}$ \\\\ References\\tablenotemark{a}  & 1,2 & 3,4 & 5 & 6,7 & 8 \\\\ \\enddata \\tablenotetext{a}{REFERENCES -- (1) Torii et al. 2000; (2) Bocchino et al. 2001; (3) Kennel \\& Coroniti 1984a; (4) Willingale et al. 2001; (5) Safi-Harb et al. 2001; (6) Wolszczan, Cordes \\& Dewey 1991; (7) Frail et al. 1996; (8) Bocchino \\& Bychkov 2001.} \\end{deluxetable}  In Table 2, we compare the values of the various parameters for the W44 PWN with those of other PWNs, including the young Crab and 3C~58, the intermediate age G21.5-0.9, and the older IC~443 PWN.  In the discussion that follows, we describe how the values for the W44 parameters were inferred, and compare them with the values from the other objects.  The X-ray flux in the 2.2-8.2 keV band, where the spectrum was measured, is 2.7$\\times$10$^{-13}$~ergs~cm$^{-2}$~s$^{-1}$. This is considerably lower than that inferred by Harrus et al. (1996), but it can be expected that the considerably greater difficulty of extracting the PWN signal from the lower angular resolution ASCA data led to a less accurate flux estimate. (Note that the ASCA detection is barely significant.) Extrapolation of the unabsorbed flux to the Einstein band (0.2-4.0 keV) yields a value of 7$\\times$10$^{-13}$~ergs~cm$^{-2}$~s$^{-1}$, and a corresponding luminosity of 6$\\times$10$^{32}$ d$_{2.6}{^2}$ ergs~s$^{-1}$ (40,000 times less luminous than the Crab). This may be compared with the luminosity of 8$\\times$10$^{32}$~ergs~s$^{-1}$ predicted using the empirically derived relation for pulsar wind nebulae (Seward \\& Wang 1988) log(L$_x$)=1.39log(\\.{E})-16.6, where L$_x$ is the PWN X-ray luminosity and \\.{E} is the rate of rotational energy loss by the pulsar, which has a value of 4.3$\\times$10$^{35}$~erg s$^{-1}$ (Wolszczan, Cordes \\& Dewey 1991). The correspondence is remarkable, considering the nebula's atypical morphology and history.  The spectral index is not observed to change radically with distance from the pulsar, in contrast to all the other PWNs listed in Table 2. Most interesting is the contrast with the most similar object known, the PWN in IC~443, whose spectral index varies by 0.7 (Bocchino \\& Bychkov 2001). As pointed out above, the overall nebular X-ray spectral index is also similar to that of the Crab Nebula. Chevalier (2000) developed a model for the X-ray luminosity of PWNs in which he claimed that the X-ray spectral index is an indicator of the efficiency with which the particle energy is converted into X-ray emission.  In particular, he argues that in PWNs with Crab-like X-ray (and thus electron) spectra, the X-ray luminosity should be produced with high efficiency.  For simplicity, his model assumes a constant magnetic field.  The field in the W44 PWN might be nearly constant, given the lack of X-ray spectral index variation. According to this model, the value and constancy of the X-ray spectral index suggest that the W44 PWN should have an \\.{E}/L$_x$ that is much closer to the Crab than the observed factor-of-five difference.  Since the only other PWN that seems Crab-like in this regard according to Chevalier (the LMC remnant 0540-69.3) is, like the Crab, about 1,000 yr old, perhaps some additional factor, such as the age of the nebula, needs to be incorporated in the model.  In the widely accepted model of PWNs, the energy source is the spindown energy of the pulsar.  The energy is efficiently transferred into a relativistic wind with some characteristic Lorentz factor $\\gamma$.  This energy is divided between Poynting (magnetic and electric) flux and particle energy flux, the ratio of which is the parameter $\\sigma$. The initial interaction between the wind and the surrounding medium forms a termination shock, usually creating wisps and filaments identifiable in the radio and visible bands (and now in the X-ray, using Chandra).  In the pulsar's magnetosphere, where the particles are created by pair production, the wind is expected to be magnetically dominated ($\\sigma$~$\\ge$~1). By the time the wind reaches the termination shock, observations indicate it becomes particle dominated ($\\sigma$~$<$~1). No clear theoretical explanation for this transition has emerged (Arons 1998). Measurements of wind parameters in other PWNs typically show $\\sigma$ on the order of a few times 10$^{-3}$ (see Table 2).  The composition of the particle flux is also of interest.  This is characterized by the parameter k, the ratio between the energy density in electrons, and that in all other particles (positrons, protons and other nuclei).  While models commonly assume an electron-positron plasma with k~=~1, Frail et al. (1996) have shown that in the W44 wind nebula 5~$\\le$~k~$\\le$~30.  The Chandra observation allows us to refine estimates of these parameters characterizing the PWN, as well as the magnetic field strength B. Under the assumption that the break between the radio and the X-ray spectra arises from synchrotron losses, Frail et al. (1996) have shown how an estimate of the turnover frequency, $\\nu_B$, of the PWN spectrum leads to estimates of the nebular magnetic field strength and the Lorentz factor, $\\gamma$, for the electrons near $\\nu_B$. From the improved measurement of the X-ray flux and spectral index of the PWN, we find a best-fit cutoff frequency $\\nu_B$ = 8$\\times$10$^{12}$~Hz. The reduced flux value is the primary reason why the estimate of the synchrotron cutoff in the present work is substantially lower than that of Harrus et al. (1996). Our best-fit $\\nu_B$ is similar to the Crab's break frequency of 10$^{13}$~Hz. Using the equations reproduced by Frail et al. from Pacholczyk (1970) we then find:  B = 1040 $\\mu$G ($\\nu_B$/10$^{12}$ Hz)$^{-1/3}$(t$_{res}$/1000 yr)$^{-2/3}$ = 160 $\\mu$G; and  $\\gamma$ $\\sim$ 10$^5$ ($\\nu_B$/10$^{12}$ Hz)$^{1/2}$ (B/100 $\\mu$G)$^{-1/2}$ = 2.2$\\times$10$^5$  \\noindent Here t$_{res}$ represents the age of the nebula, which Frail et al. take to be 5,700 yr based on the nebular extent in the radio.  The new cutoff frequency estimate falls squarely between the values used by Frail et al. (1996) to bound the value of k. Using our best value of $\\nu_B$, we find k $\\sim$ 10, consistent with their conclusionthat the particle energy is electron dominated.  We can estimate the value of the magnetization parameter $\\sigma$, following the approach used by Torii et al. (2000) for 3C~58. They used the formalism developed by Kennel and Coroniti (1984a, b) for the Crab, who showed that $\\sigma$ is related to the velocity profile of the nebula by v(z)/c~$\\sim$~3$\\sigma$[1~$+$~(3z$^2$)$^{-1/3}$].  Here z = r/r$_s$, the ratio between radius and the distance between the pulsar and the termination shock. Torii et al. estimated $\\sigma$ by es timating r$_N$ and v$_N$, the values at the edge of the radio nebula. Unlike 3C 58, in which the pulsar shows little proper motion, care must be exercised here in estimating global nebular parameters like size in the presence of substantial pulsar motion.  For the nebular size r$_N$ we take the largest dimension in the radio nebula perpendicular to the direction of motion, which is 1 arc minute or 0.75~d$_{2.6}$~pc. For the nebular age, we cannot use the nominal spin down age of 20,000~yr. Given the estimated proper motion, the pulsar has moved $\\sim$8 arc minutes, an angular distance far larger than the size of the observed nebula. Instead we use the age of the radio nebula, 5,700 yr, which is based on synchrotron lifetime arguments (Frail et al. 1996). Assuming homologous expansion, we find v$_N$ (= 2/5 $\\times$ 0.75~pc/5,700~yr)~$\\sim$~50~km~s$^{-1}$; assuming constant expansion, we find v$_N$~$\\sim$~130~km~s$^{-1}$. Another estimate of the nebular lifetime can be derived from the time required by the pulsar to move from the location corresponding to the largest perpendicular dimension to its current location; 1.5 arc minutes divided by the proper angular motion of 25 mas/yr, or 3,600 yr. This simple calculation suggests that the lifetime used can be off by not more than a factor of $\\sim1.6$.  There is no evidence in the visible, radio or X-rays of the wisps observed in the Crab or 3C~58 that are interpreted as the termination shock. Since the termination shock represents the location where the ram pressure of the relativistic wind from the pulsar (\\.{E}/(4$\\pi$cr)) equals the pressure in the nebula (presumably dominated by magnetic pressure), we can estimate r$_s$ by equating these two quantities at the termination shock. Thus \\.{E}/(4$\\pi$c r$_s$)=1/3(B$^2$/(8$\\pi$)).  For B=160 $\\mu$G and \\.{E}=4.3$\\times$10$^{35}$~ergs~s$^{-1}$ (Wolszczan, Cordes \\& Dewey 1991), we find r$_s$~$\\sim$~0.02~pc, leading to a value of z$_N$~=~r$_N$/r$_s$~$\\sim$~80, which can be compared with z$_N$ of 20 for the much younger and more energetic Crab, and 15-100 for 3C~58 (Torii et al. 2000).  Combining the estimates for z$_N$ and v$_N$, we estimate $\\sigma~\\sim~$0.4-1.0$\\times$10$^{-3}$, depending upon which value of v$_N$ is assumed. As indicated in Table 2, this value of $\\sigma$ is lower than that found for the younger PWNs, including the Crab (3$\\times$10$^{-3}$ -- Kennel \\& Coroniti 1984a), 3C~58 (2-15$\\times$10$^{-3}$ -- Torii et al. 2000), but is consistent with that for the older G21.5-0.9 (4-11$\\times$10$^{-4}$ Safi-Harb et al. 2001). While the sample of measurements is small, it suggests that $\\sigma$ may decrease with PWN age, but it either stays near a value of $\\sim$10$^{-3}$ over many thousands of years or $\\sigma$ must have a value around 10$^{-3}$ if a PWN is to be observable.  We can use the refined magnetic field estimate to obtain a synchrotron lifetime estimate for the hard band X-ray emission pictured in Figs. 1 and 2. For a magnetic field strength B in Gauss, and a photon frequency $\\nu$ in Hz, the synchrotron lifetime $\\tau$ in seconds is $\\sim$6$\\times$10$^{11}$B$^{-3/2}\\nu^{-1/2}$.  For the estimated value of B (160 $\\mu$G) and a photon frequency of 5$\\times$10$^{17}$~Hz ($\\sim$2.5~keV), we find $\\tau~\\sim$3$\\times$10$^8$~s, or approximately 15 yr. This lifetime is very short compared with the $\\sim$20,000~yr pulsar age (Wolszczan, Cordes \\& Dewey 1991) or the apparent nebular age of 5,700 yr. The short lifetime, along with the $\\sim$1~d$_3$~pc extent of the X-ray nebula implies that the X-ray emitting electrons ($\\nu$~$\\ge$~5$\\times$10$^{17}$~Hz) have a high streaming velocity, $\\sim$1/3~c. The high streaming velocity in turn suggests that the magnetic field in the extended nebula is ordered, and oriented along the wake of the pulsar's motion. It is reasonable to expect an ordered field if one considers that the PWN magnetic field should be considerably stronger than any residual magnetic field in the remnant interior, even in the mixed-morphology W44 whose interior density is substantially higher than a typical shell-like SNR.  As a final note, one feature of potential future interest is the possible low surface brightness extended X-ray emission. If it is nonthermal emission associated with the PWN, then W44 would be the third PWN with a larger apparent extent in the X-ray than in the radio. Similar structures have been found in the plerionic remnant G21.5-0.9 (Slane et al. 2000; Warwick et al. 2001) and the PWN inside IC~443 (Bocchino \\& Bychkov 2001), with clearly nonthermal X-ray emission extending well beyond the radio nebula. If the emission in W44 is confirmed, it is possible that low surface brightness radio emission is also produced but is invisible against the foreground emission from the shocked gas in the supernova remnant. The existence of such low level emission has been speculated upon by Warwick et al. (2001) for G21.5-0.9, where it might be easier to detect. The reality of the X-ray emission should be straightforwardly demonstrated by XMM/Newton, whose substantially higher throughput will facilitate more accurate measurements of the PWN spectrum and its variations.  In summary, the ACIS observation of the pulsar wind nebula surrounding PSR~B1853+01 in W44 reveals an extended nebula, half the size of the radio PWN. Spectroscopy reveals a significant difference between the power law photon index of PSR~B1853+01 ($\\Gamma$~$\\sim$~1.4) and that of the nebula ($\\Gamma$~$\\sim$~2.2).  Variation of the photon index within the nebula has not been detected.  The X-ray size and spectrum of the PWN have allowed us to estimate key parameters, including magnetic field strength, the average $\\gamma$ of the particles in the wind, the magnetization parameter ($\\sigma$), and the ratio k of electrons to other particles. We find that despite the unusual morphology produced by the high velocity and age of the pulsar, the W44 PWN has properties similar typical of other PWN. A number of unresolved issues, such as the possible existence an extended nebula and spectral variation within the nebula, make this fascinating object deserving of more extensive study. It is hoped that deeper observations using either Chandra or XMM/Newton can provide more exact values of the observables and facilitate measurements of their variation with distance from PSR~B1853+01, leading in turn to more robust estimates of the PWN parameters."
    },
    "0207/hep-ph0207116_arXiv.txt": {
        "abstract": "We use the field theoretical model to perform relativistic calculations of neutrino energy losses caused by the direct Urca processes on nucleons in the degenerate baryon matter. By our analysis, the direct neutron decay in the superdense nuclear matter under beta equilibrium is open only due to the isovector meson fields, which create a large energy gap between protons and neutrons in the medium. Our expression for the neutrino energy losses, obtained in the mean field approximation, incorporates the effects of nucleon recoil, parity violation, weak magnetism, and pseudoscalar interaction. For numerical testing of our formula, we use a self-consistent relativistic model of the multicomponent baryon matter. The relativistic emissivity of the direct Urca reactions is found substantially larger than predicted in the non-relativistic approach. We found that, due to weak magnetism effects, relativistic emissivities increase by approximately 40-50\\%, while the pseudoscalar interaction only slightly suppresses the energy losses, approximately by 5\\%. ",
        "introduction": "Modern calculations \\cite{PBPELK97} based on relativistic equations of state indicate that the neutron star cores consist of neutrons with the admixture of protons, electrons, muons and some exotic particles (including hyperons, K-mesons, quarks and so on\\ldots ). The composition is governed by the charge neutrality and equilibrium of the medium under weak processes $B_{% {\\rm 1}}\\rightarrow B_{{\\rm 2}}+l+\\bar{\\nu}_{l}$, $\\ \\ \\ B_{{\\rm 2}% }+l\\rightarrow B_{{\\rm 1}}+\\nu _{l}$, where $B_{{\\rm 1}}$ and $B_{{\\rm 2}}$ are baryons (or quarks), and $l$ is a lepton, either an electron or a muon. These reactions, widely known as the direct Urca processes, are a central point of any modern scenarios of evolution of neutron stars. Neutrino energy losses caused by the direct Urca processes lead to a rapid cooling of degenerate neutron star cores \\cite{LRP94}. The corresponding reactions on nucleons, $n\\rightarrow p+l+\\bar{\\nu}_{l}$, $\\ \\ \\ p+l\\rightarrow n+\\nu _{l}$% , are the most powerful sources of neutrinos and antineutrinos in cooling neutron stars. In spite of widely adopted importance of these reactions, the corresponding neutrino energy losses are not well investigated yet. A simple formula suggested by Lattimer et al. \\cite{Lat91} more than ten years ago has been derived in a non-relativistic manner. Actually, however, the superthreshold proton fraction, necessary for the direct Urca processes to operate in the degenerate nuclear matter, appears at large densities, where the Fermi momenta of participating nucleons are comparable with their effective mass. Moreover, according to modern numerical simulations, the central density of the star can be up to eight times larger than the nuclear saturation density \\cite{PBPELK97}. This implies a substantially relativistic motion of nucleons in the superdense neutron star core. The appropriate equation of state for such a matter is actually derived in the relativistic approach, and the relevant neutrino energy losses must be consistent with the used relativistic equation of state. Some aspects of this problem was studied by Leinson and P\\'{e}rez \\cite{PLB}, who have estimated relativistic effects of baryon recoil and parity violation in the direct Urca processes including also effects of the baryon mass difference. This approach is useful, for example, for the direct Urca reactions changing the baryon strangeness because, in this case, the contribution of the weak magnetism into the matrix element of the beta-decay is not well known even for free particles. For nucleons, the relativistic regime should incorporate the effects of weak magnetism and pseudoscalar interaction, which drastically influence the neutrino emissivity of the corresponding reactions. To demonstrate this, in the present paper we consider the totally relativistic direct Urca process on nucleons. We utilize the Walecka-type relativistic model of baryon matter \\cite{Serot}, where the baryons interact via exchange of $\\sigma $, $\\omega $, and $\\rho $ mesons, and perform the calculation of the neutrino energy losses in the mean field approximation. This approximation is widely used in the theory of relativistic nuclear matter, and allows to calculate in a self-consistent way the composition of the matter together with energies, and effective masses of the baryons. In Section II we begin with considering the relativistic kinematics of the neutron beta decay in the medium under beta equilibrium. We consider a free gas model and some field theoretical models of nuclear matter to demonstrate that the direct neutron decay in neutron stars is open only due to strong interactions caused by isovector mesons. We shortly discuss the energies and the wave functions of nucleons in the mean field approximation and demonstrate the nonconservation of the charged vector current of nucleons in this model. The matrix element of the neutron beta decay is derived in Section III. In Section IV we calculate the neutrino energy losses caused by the direct Urca on nucleons in the degenerate nuclear matter under chemical and thermal equilibrium. In Section V we inspect the non-relativistic limit of the neutrino energy losses in order to compare this with the expressions earlier obtained in \\cite{Lat91} and \\cite{PLB}. Efficiency of the relativistic approach is numerically studied in Section VI. We evaluate the neutrino energy losses due to the direct Urca processes on nucleons in the multicomponent baryon matter under beta equilibrium and compare the relativistic result with that predicted by the known non-relativistic formula. We specially discuss the contributions of weak magnetism and pseudoscalar interaction. Summary and conclusion are in Section VII. In Appendix, we discuss some details of the model used for the nuclear matter.  In what follows we use the system of units $\\hbar =c=1$ and the Boltzmann constant $k_{B}=1$. Summation over repeated Greek indexes is assumed ",
        "conclusions": "In the mean field approximation, we have studied the neutrino energy losses caused by the direct Urca processes on nucleons in the degenerate baryon matter under beta equilibrium. We have shown that the direct Urca processes in a superdense matter of neutron star cores are kinematically allowed only due to isovector mesons, which differently interact with protons and neutrons. By creating the energy gap between proton and neutron spectrums the isovector mesons support a time-like total momentum transfer from the nucleon, as required by kinematics of the reaction. In the mean field approximation, we derived the matrix element of the nucleon transition current, which is found to be a function of the space-like kinetic momentum transfer. We have calculated the neutrino energy losses caused by the direct Urca processes on nucleons. Our Eq. (\\ref{QMFA}) for neutrino energy losses exactly incorporates the effects of nucleon recoil, parity violation, weak magnetism, and pseudoscalar interaction. To quantify the relativistic effects we consider a self-consistent relativistic model, widely used in the theory of relativistic nuclear matter. The relativistic energy losses are up to four times larger than those given by the non-relativistic approach. In our analysis, we pay special attention to the effects of weak magnetism and pseudoscalar interaction in the neutrino energy losses. We found that, due to weak magnetism effects, relativistic emissivities increase by approximately 40-50\\%, while the pseudoscalar interaction only slightly suppresses the energy losses, approximately by 5\\%. The mean field Eq. (\\ref% {QMFA}) may be considered as a starting point for studying of the correlation effects."
    },
    "0207/astro-ph0207279_arXiv.txt": {
        "abstract": "Contrary to theoretical expectations, observations with the {\\em Rossi X-ray Timing Explorer (RXTE)} show that in X-ray binaries timing properties are not uniquely correlated with X-ray luminosity. For instance, although the frequencies of the kilohertz quasi-periodic oscillations (kHz QPOs) correlate with X-ray flux on short ($\\sim$few hours) time scales, on time scales longer than a day the QPO appears at more or less the same frequency, whereas the luminosity may be a factor of a few different. The result is a set of almost parallel tracks in a QPO frequency vs. X-ray flux plot. Despite the ``parallel tracks'' are a common phenomenon among kHz QPO sources, until now, after five years of observations with RXTE, not a single transition between two of these tracks had been seen. Here I present the first detection of such a transition, in 4U\\,1636--53. ",
        "introduction": "Observations with the {\\em Rossi X-ray Timing Explorer (RXTE)} have revealed kilohertz quasi-periodic oscillations (kHz QPOs) in some 20 X-ray binaries. It is generally thought that these kHz QPOs reflect the motion of matter in orbit at some preferred radius in the accretion disk around the neutron star (see van der Klis 2000 for a review of the phenomenology of these QPOs, and for a description of the models so far proposed to explain them).  Calculations by Miller, Lamb, \\& Psaltis (1998) show that the inner radius of the disk is set by angular momentum losses to the radiation field: When mass flow through the disk increases, the inner radius of the disk decreases, and therefore the QPO frequency (the Keplerian frequency at this radius) increases. But because in accretion-powered systems luminosity is proportional to the mass accreted onto the compact object, from the above it follows that there should be a one-to-one relation between QPO frequency and bolometric flux.  To the extent that X-ray flux is a good measure of the bolometric flux (see, e.g., Ford et al. 2000), observations with RXTE seem to contradict the above expectations. Figure 1a shows a plot of QPO frequency vs. X-ray intensity for the transient source 4U\\,1608--52 (M\\'endez et al. 1999). Each segment there represents an uninterrupted observation lasting $\\simless$1 hour, whereas different tracks denote observations separated by intervals longer than a day.  From this Figure it is apparent that frequency and X-ray count rates are positively correlated during relatively short periods, but they are uncorrelated over longer time intervals. ",
        "conclusions": "Despite 5 years of observations with RXTE, so far no single transition between two tracks in a QPO frequency vs. X-ray intensity diagram had been observed. Here I report the first direct observation of one such transition, in an RXTE observation of 4U\\,1636--53 made in 1998 (see Figure 1b). The observation starts with the source at the upper end of ``Track 1'' (``BEGIN''), and ends at the upper end of ``Track 2'' (``END''). In between, 4U\\,1636--53 goes three times from one track to the other, as indicated by the arrows with numbers 1, 2, and 3. The transitions are very fast ($\\simless$300 s), consistent with upper limits inferred from the time intervals between two consecutive tracks in this and other sources, in diagrams similar to the one shown in Figure 1a (e.g., M\\'endez 2000).  \\begin{figure} \\plotfiddle{mendez_fig1.eps}{5.15cm}{0}{30}{30}{-160}{0} \\caption{(a) QPO frequency vs. X-ray intensity for 4U\\,1608--52, showing the so-called ``parallel tracks''. (b) Transition between two tracks (gray and black points) in 4U\\,1636--53.} \\end{figure}  During the transitions in 4U\\,1636--53, X-ray intensity changes by $\\simless$3\\,\\%. From Figure 1a we see that in 4U\\,1608--52 X-ray intensity differs by factors of a few between tracks (in 4U\\,1636--53 jumps of $\\sim$30\\,\\% are observed over longer time intervals than shown in Figure 1b), it is therefore not clear whether this larger differences in X-ray intensity are the cumulative effect of several small jumps, or if they have a completely different origin.  I will present these results in more detail elsewhere (M\\'endez \\& van der Klis, 2002, in preparation), and there I will discuss more extensively their implications upon models of the accretion flow in X-ray binaries."
    },
    "0207/astro-ph0207565_arXiv.txt": {
        "abstract": "We present high-resolution ($<$0\\farcs5) mid-infrared Keck II images of individual sources in the central region of NGC 6334 I. We compare these images to images at a variety of other wavelengths from the near infrared to cm radio continuum and speculate on the nature of the NGC 6334 I sources. We assert that the cometary shape of the UCHII region here, NGC 6334 F, is due to a champagne-like flow from a source on the edge of a molecular clump and not a due to a bow shock caused by the supersonic motion of the UCHII region through the interstellar medium. The mid-infrared emission in concentrated into an arc of dust that define the boundary between the UCHII region and the molecular clump. This dust arc contains a majority of the masers in the region. We discuss the nature of the four near-infrared sources associated with IRS-I 1, and suggest that one of the sources, IRS1E, is responsible for the heating and ionizing of the UCHII region and the mid-infrared dust arc. Infrared source IRS-I 2, which has been thought to be a circumstellar disk associated with a linear distribution of methanol masers, is found not to be directly coincident with the masers and elongated at a much different position angle. IRS-I 3 is found to be a extended source of mid-infrared emission coming from a cluster of young dusty sources seen in the near-infrared. ",
        "introduction": "NGC 6334 is a parsec long train of rich molecular clouds and HII regions located at galactic coordinates \\textit{l}=351\\arcdeg, \\textit{b}=0.\\arcdeg% 7. The complex lies at a distance of 1.74 kpc from the Sun \\citep{Nec78}, parallel to and located in the Carina-Sagittarius spiral arm. It is the site of possibly the largest number of recently formed OB stars observed in the Galaxy, which may have been triggered by the recent passage of a spiral density wave \\citep{hg83}.  NGC 6334 was first discovered in the far-infrared by \\citet{em73}. Later observations in the far-infrared by \\citet{mcb79}, revealed six centers of emission. They were labeled by increasing southern declination using Roman numerals I-VI. Our observations were of NGC 6334 I, the northernmost far-infrared region of NGC 6334, and the site of a well-studied ultracompact HII region, NGC 6334 F. Though heavily obscured at visual wavelengths, NGC 6334 I is the center of a wealth of activity in the infrared, millimeter, and radio, as well as the site of many molecular sources and masers. Over the decades, many authors have studied this region of NGC 6334, and each, it seems, used nomenclature of their own to describe it. NGC 6334 I is a large region that is identified with several significant sources summarized by % \\citet{KDJ99}. We will use the convention NGC 6334 F from the radio continuum observations of \\citet{rcm82} to describe the UCHII region we observed in the mid-infrared. The HII region is clearly cometary shaped in the radio \\citep{rcm82,drdg95,enm96}, millimeter \\citep{ckrdh97}, and mid-infrared \\citep{DPT00,ptf98}, with its head pointing to the northwest and the tail running to the southeast. The peak of the UCHII region lies near the infrared source IRS-I 1 of \\citet{bn74}, which has been presumed to be the ionizing source of the HII region. \\citet{hg83} also find another source $\\sim $6$\\arcsec$ to the northwest of IRS-I 1, designated IRS-I 2, and yet another $\\sim $18$\\arcsec$ to the east, designated IRS-I 3.  This region is very complex and is the site of a wide variety of activity. A near-infrared survey by \\citet{t96} found an embedded young cluster of 93 sources associated with NGC 6334 I, all within a radius of $\\sim $80$\\arcsec$% . This cluster, of which IRS-I 1, IRS-I 2, and IRS-I 3 are members, appears to only contain stars earlier than B3-B4 according to \\citet{t96}. In light of the complexity of the NGC 6334 I area, interpretation of data is not an easy task. In this paper we present high-resolution mid-infrared images of the sources within NGC 6334 I. In \\S 2 we will discuss the observations of NGC 6334 I, and explain the data reduction process in \\S 3. Interpretation of our data and a discussion of the phenomenology of each source will be presented in \\S 4. Finally, in \\S 5 we will present our conclusions. ",
        "conclusions": "High resolution mid-infrared observations of the central region of NGC 6334 I have revealed much about the nature and properties of the sources there. The UCHII region NGC 6334 F is composed of two sources, IRS-I 1 and G351.42+0.64:DPT00 2. The peak of IRS-I 1 appears to be coincident with the peak in the radio continuum. Ammonia observations of Kraemer et al. (1999) when registered properly with our mid-infrared data indicate that the shape of the UCHII region is not due to a bow-shock, but instead due to champagne-like flow from stellar source at the edge of a molecular clump. Maser emission is concentrated at this interface between the mid-infrared and ammonia emission, and may therefore be shock induced. There are two other strings of masers that lie near the two peaks in the ammonia emission and may be delineating the sites of hot molecular cores that are too young and/or embedded to be seen yet in the mid-infrared.  The mid-infrared emission from IRS-I 1 seems to be coming from an arc of dust at the interface between the molecular ammonia clump and the UCHII region, and may be material swept up by the expanding shock front of the UCHII\\ region. The color temperature peaks at a location interior to this mid-infrared arc, coincident with a near infrared source IRS1E. This source may be the stellar source responsible for the ionization and heating of the NGC\\ 6334 F region. Two other near infrared sources, IRS1W and IRS1SW, lie in the northern and southern parts of the mid-infrared arc and are associated with the majority of the masers in the region. There is no temperature peak at these locations, so the near infrared emission may just be reflected or shock excited emission. A fourth near infrared source (IRS1SE) seems to simply be reflected emission off the UCHII\\ region tail.  G351.42+0.64:DPT00 2 appears to be a clump of dust, perhaps swept up by the shock front of UCHII region. It displays a steep temperature gradient towards the color temperature peak. It also shows some signs of ionized emission in the higher resolution radio continuum images, but only on the hotter, southern side. For these reasons it may be that DPT00 2 has no central heating source but is simply heated and ionized by the same source heating and ionizing IRS-I 1 (IRS1E) .  IRS-I 2 was believed to be associated with the a linear structure of methanol masers and perhaps delineating a circumstellar disk. However, the thermal dust emission is elongated at a different position angle to the position angle of the maser distribution. The low-surface brightness, smooth color temperature distribution, and lack of a near-infrared component may indicate that there is no internal stellar source here at all. Furthermore, the masers are offset from the mid-infrared peak and could be associated with the secondary peak in the ammonia distribution.  Lower resolution mid-infrared images of IRS-I 3 showed it to be a double peaked source. However, the high resolution images presented here show that it has a complex and peculiar morphology. We find, using the near infrared data of Persi et al. (1996), that the large and extended mid-infrared sources are extended dust emission from a cluster of stellar sources seen in the near infrared. These stellar sources have directly influenced the morphology in the mid-infrared, and the structure of the source as seen in the color temperature map.  The reality of the 7 mm source in this region has been seriously called into question by the non-detection in the mid-infrared, and has recently been discovered to be an artifact of data reduction by follow-up observations of the original authors."
    },
    "0207/astro-ph0207086_arXiv.txt": {
        "abstract": "We present observations of 8 Galactic Bulge microlensing events taken with the 1.0m JKT  on La Palma during 2000 June and July. The JKT observing schedule was optimized using a prioritizing algorithm to automatically update the target list. For most of these events we have sampled the lightcurves at times where no information was available from the OGLE alert team. We assume a point-source point-lens (PSPL) model and perform a maximum likelihood fit to both our data and the OGLE data to constrain the event parameters of the fit. We then refit the data assuming a binary lens and proceed to calculate the probability of detecting planets with mass ratio $q=10^{-3}$. We have seen no clear signatures of planetary deviations on any of the 8 events and we quantify constraints on the presence of planetary companions to the lensing stars. For two well observed events, 2000BUL31 and 2000BUL33, our detection probabilities peak at $\\sim$30\\% and $\\sim$20\\% respectively for $q=10^{-3}$ and $a \\sim R_{\\mbox{E}}$ for a $\\Delta\\chi^2$ threshold value of 60. ",
        "introduction": "Microlensing alters the path followed by the photons emitted by a background stellar source as they come near the influence of the gravitational field of a massive foreground object which acts as a lens. The separation of the images created by the lensing effect ($\\sim 10^{-3}$arcsec) is too small to be resolved and only the combined flux is observed. The resulting lightcurve is symmetric in time with its maximum amplification at the time of closest approach between the projected position of the source on the lens plane and the lens itself. Its shape is well described by the formula: \\begin{equation} A(u) = \\frac{u^2 + 2}{u (u^2 + 4)^{1/2}} \\end{equation} where $A(u)$ is the total amplification and $u$ is the angular separation of source and lens in units of the angular Einstein ring radius $\\theta_{\\mbox{E}}$ \\cite{Pacz86}.  Even though most of the cases can be adequately described by this simple model, the shape of the lightcurve may not be symmetric and may be exhibiting significant deviations. These so-called anomalies of the lightcurve can be due to several factors and have been extensively examined in recent literature \\cite{Dominik99b,Wozniak97,Buchalter97,Alcock95b,Gaudi97}. Arguably the most interesting of these are the anomalies which can be attributed to the binary nature of the lens. The possibility of the secondary component being an object of planetary characteristics has spawned dedicated observing campaigns to reveal their presence \\cite{albrow98,rhie2000,bond01}. ",
        "conclusions": "We have followed 8 microlensing events using the JKT on La Palma for 2hrs per night from 6 June to 17 July 2000. We presented fits to the combined JKT-OGLE datasets and recalculated the event parameters. We searched the data for signatures of planetary companions with mass ratio $q=10^{-3}$ but have seen no indications of a planetary presence in the datasets. Finally we calculated the planetary detection probabilities on all the events for a mass ratio of $q=10^{-3}$. For events 2000BUL31 and 2000BUL33, our detection probabilities peak at $\\sim$30\\% and $\\sim$20\\% respectively for $q=10^{-3}$ and $a \\sim R_{\\mbox{E}}$ for ${\\Delta\\chi^2}_{\\mbox{thr}}$=60."
    },
    "0207/astro-ph0207615_arXiv.txt": {
        "abstract": "We calculate the Equivalent Widht of the Core and the centroid energy and relative flux of the 1st order Compton Shoulder of the iron K$\\alpha$ emission line from neutral matter. The calculations are performed with Monte Carlo simulations. We explore a large range of column densities for both transmitted and reflected spectra, and study the dependence on the iron abundance. The Compton Shoulder is now becoming observable in many objects thanks to the improved sensitivity and/or energy resolution of XMM--$Newton$ and $Chandra$ satellites, and the present work aims to provide a tool to derive informations on the geometry and element abundances of the line emitting matter from Compton Shoulder measurements. ",
        "introduction": "Iron K$\\alpha$ fluorescent lines emitted in neutral matter consists (in the matter rest frame) of a narrow core (corresponding to the line photons emerging unscattered from the emitting region) and several Compton Shoulders (CS; see Matt et al. 1991; George \\& Fabian 1991; Leahy \\& Creighton 1993; Sunyaev \\& Churazov 1996), corresponding to line photons emerging after one or more scatterings. While the higher order CS are expected to be very faint, the first order CS (hereinafter CS1) should now be observable in many objects thanks to the improved energy resolution and sensitivity of the instruments onboard $Chandra$ and XMM--Newton (see Kaspi et al. 2002; Bianchi et al. 2002; see also Iwasawa et al. 1997 for the only pre-$Chandra$ observation of a Compton Shoulder).  In this paper we calculate, by means of Monte Carlo simulations, the relative amount and centroid energy  of CS1, as well as the Equivalent Width (EW) of the unscattered line photons (Narrow Core, hereinafter NC),  in both transmitted and reflected spectra. In the former case, we adopt a spherical geometry, and explore a large interval of the column density. For the latter case we assume a plane--parallel geometry and calculate CS1 properties as a function of the inclination angle for different values of the column density in the perpendicular direction. The dependence of line properties on the iron abundance is also explored. ",
        "conclusions": "We have calculated the properties of the iron K$\\alpha$ Narrow Cores and of the first scattering Compton Shoulder in both transmitted and reflected spectra, refining and expanding previous works on the same subject (Matt et al. 1991; George \\& Fabian 1991; Leahy et al. 1993; Iwasawa et al. 1997; see Sunyaev \\& Churazov 1996 for a very detailed physical description of the process). The intensity of the Compton Shoulder is at most 30-40\\% of the Narrow Core of the line, with a centroid energy typically of about 6.3 keV or more.  Prior to the launch of $Chandra$, the Compton Shoulder was observed only in the ASCA spectrum of NGC~1068 (Iwasawa et al. 1997), with properties consistent with an origin in the Compton--thick torus of this source. Recently, thanks to the gratings onboard $Chandra$, two more cases have been reported. Bianchi et al. (2002) found a Compton shoulder in the HETG spectrum of the Circinus Galaxy. The value of $f$, i.e. about 20\\%, is consistent with reflection from the inner wall of the 4$\\times10^{24}$ cm$^{-2}$ torus.  Kaspi et al. (2002) measured a Compton shoulder in their 900 ks $Chandra$/HETG observation of the Seyfert 1 galaxy NGC~3783. The flux in CS1 is 14$\\pm$4 percent of the narrow core, which in turn has an EW of 90$\\pm$11 eV. The line is clearly due to reflecting matter, as no evidence for strong cold absorption is apparent. Looking at the values of $f$ in Figs.~\\ref{refl} and \\ref{refl_thin_all}, it is possible to conclude that the reflecting matter should have a column density of at least $\\sim10^{23}$ cm$^{-2}$ (slightly lower values are allowed only if the iron is overabundant). The value of the EW is less directly usable, as the geometry of the reflector is unknown (the values in this paper refer to a 2$\\pi$ solid angle of the matter, subtended to an isotropic primary source). However, as it seems unlikely that the solid angle is much larger than 2$\\pi$, low values of the iron abundance are disfavoured."
    },
    "0207/astro-ph0207423_arXiv.txt": {
        "abstract": "We formulate a Zel'dovich-like approximation for the Chaplygin gas equation of state $P = - A/\\rho$, and sketch how this model unifies dark matter with dark energy in a geometric setting reminiscent of $M$-theory. ",
        "introduction": "In the last few years improved observations \\cite{mel1} have forced a shift in our cosmological paradigm: the $\\Omega_{\\rm M} = 1 $ dust model has been swept aside and, in its place, we are faced with the problem of understanding a universe with an equation of state $\\bar{W} = \\bar{P}/\\bar{\\rho} < - 1/3$. That is to say, on average, pressure is comparable with density and, moreover, negative.  Of course, parametrically, this is readily accommodated by a cosmological constant $\\Lambda$ \\cite{peeb2} with $\\Omega_{\\Lambda} \\simeq$ 0.7 and $\\Omega_{\\rm DM} = 1 - \\Omega_{\\Lambda} \\simeq$ 0.3 (throughout the paper we neglect the small baryonic contribution). The well-known difficulty with $\\Lambda$ is that {\\it a priori} it seems an incredible accident that $\\Omega_{\\Lambda} \\simeq \\Omega_{\\rm DM}$ since $\\rho_{\\Lambda}/\\rho_{\\rm DM} \\sim a^{3}$, $a$ being the scale factor. Hence, much attention has been devoted to quintessence \\cite{wett3}, involving a real scalar field which tracks \\cite{zlat4} the background component until recently becoming dominant. However, simple tracking quintessence does not work \\cite{blud5}   and spintessence \\cite{boyl6}, where the scalar field is complex, suffers instabilities against the decay of dark energy into dark matter \\cite{kasu7}.  It is natural to conjecture that some of the aforementioned problems derive from treating dark matter and dark energy as separate issues. As an example, Barr and Seckel \\cite{barr8} have pointed out that in axion dark matter models quantum gravity effects break the Pecci-Quinn symmetry leading to a universe trapped in a false vacuum with an effective $\\Lambda$ of the correct magnitude. In another approach, Wetterich \\cite{wett9} has suggested that traditional WIMP dark matter should be replaced by quintessence lumps, thus unifying dark matter and dark energy. However, the radiation-matter transition and structure formation remain  open questions in this scenario.  Herein we present a dark matter-energy unification model suggested by the observation of Kamenshchik et al. \\cite{kamen10} that a perfect fluid obeying the Chaplygin gas equation of state \\begin{equation} P = - \\frac{A}{\\rho} \\label{eq01} \\end{equation} should lead to a homogenous cosmology with \\begin{equation} \\bar{\\rho} (a) = \\sqrt{A + \\frac{B}{a^{6}}}  , \\label{eq02} \\end{equation} with $B$ being an integration constant, thus interpolating between dark matter, $\\bar{\\rho} (a \\rightarrow 0) \\simeq \\sqrt{B}/a^{3}$ and dark energy $\\bar{\\rho} (a \\rightarrow \\infty) \\simeq \\sqrt{A}$. Before doing so, we must first show why Eq. (\\ref{eq01}), aside from its interesting mathematical features \\cite{jach11}, might describe reality. ",
        "conclusions": "A shortcoming of the Zel'dovich approximation is that at the caustic matter flows through unimpeded so that structures quickly dissolve \\cite{pauls18}. This may be circumvented via the truncated Zel'dovich approximation \\cite{pauls18}. A preferable alternative would be an extension of the adhesion approximation \\cite{gurb19} which also allows the extraction of mass functions.  Approximation technicalities aside, the case is made that the Chaplygin gas offers a realistic unified model of dark matter and dark energy. That this is achieved in a geometric (brane world) setting rooted in string/$M$ theory makes this model all the more remarkable."
    },
    "0207/astro-ph0207109_arXiv.txt": {
        "abstract": "{We present and discuss new long-slit Echelle spectra of the LMC LBV candidate \\sk\\ and put them in context with previous images and spectra. While at first glance a simple spherically expanding symmetric shell, we find a considerably more complex morphology and kinematics. The spectra indicate that morphologically identified deviations from sphericity are outflows of faster material out of the main body of \\sk. The morphological as well as the kinematic similarity with other LBV nebulae makes it likely that \\sk\\ is an LBV candidate, indeed, and poses the question in how far outflows out of expanding LBV nebulae are a general property of such nebulae---at least during some phases of their evolutions. ",
        "introduction": "Stars are known up to masses around 100\\,M$_{\\sun}$ with main sequence luminosities of $10^{5-6}$\\,L$_{\\sun}$. As O stars they are located in the upper left part of the {\\it Hertzsprung-Russell Diagram} (HRD). In their later evolution, they leave the main-sequence at almost constant luminosity towards redder spectral types, i.e, they quickly cool and become supergiants. However, instead of completing their evolution towards the red, the most massive ones among them enter a phase of very high mass loss (up to 10$^{-4}$\\,M$_{\\sun}$yr$^{-1}$) and reverse the direction of their evolution, i.e., the stars become hotter again (e.g., Schaller 1992; Langer et al. 1994; Schaerer 1996a,b). These stars are known as {\\it Luminous Blue Variables\\/} (LBVs). The location of the turning points of LBVs in the HRD is called the {\\it Humphreys-Davidson limit} (HD-limit; Humphreys \\& Davidson 1979, 1994, Langer et al.\\ 1994). It is a function of the stellar luminosity. The position of most LBVs and LBV candidates known is illustrated in Fig\\footnote{ This figure is nearly identical to Fig.\\ 9 from Humphreys \\& Davidson (1994), and was copied with the permission of the authors.}.\\ \\ref{fig:hrdlbv}, where a solid line marks the empirical HD-limit. Since LBVs show also spectral variability both the hottest and coolest temperature (e.g., their position in the non-eruptive and eruptive states) are indicated with filled and open dots. The regime of the hypergiants, red supergiants and the location of the main-sequence is shown for reference, as is the position of the precursor of SN1987A. Presently, only a few LBVs are known (roughly 40, including candidate objects, see, e.g., Humphreys \\& Davidson 1994), of which 9 are in our Galaxy and 10 in the LMC.  The strong stellar winds and possible so-called giant eruptions, lead to the formation of nebulae around LBVs, the {\\it LBV nebulae\\/} (e.g., Nota et al.\\ 1995, Weis 2001). These LBV nebulae are typically up to 2\\,pc in diameter. Because of this small size, they can be studied only in our galaxy and---with the high-resolution of the {\\it Hubble Space Telescope\\/} (HST)---in a few neighboring galaxies like, for instance, the {\\it Large Magellanic Cloud\\/} (LMC). For a better understanding of the evolution of LBVs and especially the formation of LBV nebulae a good knowledge of the parameters of the nebulae around LBVs is of great interest.  \\begin{figure} \\begin{center} {\\resizebox{\\hsize}{!}{\\includegraphics{f1.eps}}} \\end{center} \\caption{HRD with the position of known LBVs and LBV candidates. The position of \\sk\\ is indicated and similar to that of S119 or WRA 751, two LBV candidates. For reference the main sequence, the regime of the hypergiants and red supergiants and the position of the SN1987A precursor are plotted in addition. \\sk\\ was positioned according to the temperature and luminosity determined by Thompson et al.\\ (1982).} \\label{fig:hrdlbv} \\end{figure} ",
        "conclusions": "The images of the nebula around \\sk\\ as well as the spectra show several cases of deviations from the morphology and kinematics of a spherical nebula with symmetric expansion.  Due to the lack of sufficiently resolved images of the nebula, one can only speculate---with the help of spectra---about the three dimensional structure of the nebula. While part of its main body may reasonably well be approximated by a sphere with a symmetric expansion, the nebula is expanding faster in the north-eastern hemisphere comparing, for instance the higher expansion velocities detected in Slits 3W (17.5\\,\\kms) and Slit 4N (16.9\\,\\kms) with those in the nebula's center (typically 15\\,\\kms, Slit 3E or Slit Center NS) or the southern part of the nebula which expands even slower than the center (see $pv$-diagrams in Fig. \\ref{fig:radialrun}). Even though the difference of the expansion velocities is marginal, it is significantly higher than the expected errors, and therefore most likely real.  Knot E shows an extension of the nebula of about 3\\farcs2 (0.8\\,pc) and moves faster than the shell (red-shifted). Here material seems to surpass the shell and move away from the observer. Similarly, the material of Knot N approaches us faster than most of the shell (see blue-shifted extension in Slit Center NS). In addition to this faster moving material of Knot N, there gas is streaming out. This out-streaming gas---as traced in the spectra---corresponds to the morphologically identified Fil N. Fil N has a length of roughly half the diameter of the shell and moves up to 15\\,\\kms\\ faster along the line-of-sight. Throughout the nebula---including Fil N---we find a \\NH\\ ratio in the range of $0.65 - 0.7$ and thus further evidence for Fil N being a part of \\sk's nebula.  The spectra also indicate that to the north-east the shell of the nebula around \\sk\\ is not closed. Together with the faster moving Fil N it is therefore most likely that the nebula around \\sk\\ shows an outflow. The north-eastern part of the nebula has opened up and material is streaming out evidenced in particular by Fil N. Both properties, open Doppler ellipses and an increasing velocity extension, were also found in the nebula around the LMC LBV candidate S\\,119 (Weis et al.\\ 2002). The kinematics and morphology of the nebula around S\\,119 are very similar to those found here. In particular, it is worthwhile noting that this outflow from \\sk's nebula, shows the same linear increase in velocity as the outflow in S\\,119. The out-streaming material moves the faster the further it is away from the star. Comparing both objects we conclude that the nebulae of these stars are quite similar and represent (candidate) LBV nebulae with outflows.  It is not clear whether Knot E is also part of an outflow or if this feature resembles more an asymmetry in the nebula comparable to the Caps in WRA\\,751 (Weis 2000) or R\\,127 (Weis 2002, in prep). An interpretation as an outflow or as caps is supported by the faster moving extension in Slit 3W.  The dynamic age of a nebula is defined by $\\tau = \\eta (r/v_{\\rm exp})$, with $\\eta$ depending on whether we model a steady fast wind swept up bubble ($\\eta=1$, Garc{\\'\\i}a-Segura \\& Mac Low 1995), an energy-conserving bubble ($\\eta = 0.6$, Weaver et al.\\ 1977) or a momentum-conserving bubble ($\\eta=0.5$, Steigman et al.\\ 1975). For \\sk\\ this would yield a range for the dynamic age between $6.3\\,10^4$ (for a maximum expansion velocity of 17.5\\,\\kms\\ and a momentum conserving bubble) to $1.5\\,10^5$ years (with the slowest expansion of 14\\,\\kms\\ and steady sweep-up). Compared to dynamic ages of other LBV nebula (ranging typically between $0.01-2\\,10^4$, see, e.g., Nota et al. 1995) this value is up to a factor of 10 higher than the average. With the estimated duration of the LBV phase of about 25\\,000 years (Maeder \\& Meynet 1987, Humphreys \\& Davidson 1994, Bohannan 1997) the dynamic age would be too high for the nebula being created during the stars LBV phase. In this context, however, one has to keep in mind that the duration of the LBV phase is mainly estimated by determining the ratio of LBVs stars to Wolf-Rayet stars, i.e., a method with comparatively large inherent errors, due to incompleteness of the respective samples. In the context of \\sk\\ beeing a nebula with an outflow, as proposed here, these age determinations with the help of the nebula's size and expansion velocity might well be overestimated. If the nebula was disrupted by the outflow, most likely the expansion velocity of the nebula decreased since the pressure dropped. Therefore the dynamic age determination yields a higher age, as a result from the lower expansion velocity.  The dynamic age for S\\,119 another LBV candidate showing outflow, for instance, also seems to be slightly higher compared to other LBV nebulae, reach a  maximum of $\\sim 4\\,10^4$ years (Weis et al.\\ 2002).  While from the kinematics alone it is not clear yet clear in each case which features are part of an outflow and which are cap-like extension of the shell, the structure of the expansion ellipses as well as Fil N show that at least these features are due to an outflow. The similarity with other LBVs and their nebulae makes it likely that \\sk\\ is at least a good candidate LBV, indeed, and poses the question in how far outflows out of expanding LBV nebulae are a general property of such nebulae---at least during some phases of their evolutions---and thus connects directly to the question for their origin and evolution."
    },
    "0207/astro-ph0207345_arXiv.txt": {
        "abstract": "We study the motions of small solids, ranging from micron-sized dust grains to 100-m objects, in the vicinity of a local density enhancement of an isothermal gaseous solar nebula. Being interested in possible application of the results to the formation of clumps and spiral arms in a circumstellar disk, we numerically integrate the equations of motion of such solids and study their migration for different values of their sizes and masses and also for different physical properties of the gas, such as its density and temperature. We show that, considering the drag force of the gas and also the gravitational attraction of the nebula, it is possible for solids, within a certain range of size and mass, to migrate rapidly (i.e. within $\\sim$1000 years) toward the location of a local maximum density where collisions and coagulation may result in an accelerated rate of planetesimal formation. ",
        "introduction": "It is generally believed that planet formation starts as a secondary process to star formation by coalescence of small bodies in circumstellar disks. With regard to our solar system, two mechanisms have been proposed for the formation of the giant planets in such a disk around our Sun; the widely accepted core accretion model \\citep{Pol96} and the disk instability scenario \\citep{Bos00}. It has recently been noted that a solar nebula massive enough to possibly form giant planets via the core accretion model is likely marginally-gravitationally unstable \\citep{Pol96,Bos00,Inab02}. The alternative approach, namely the disk instability mechanism, however, implies that such an instability could lead to rapid formation of gas giant planets. It is, therefore, of great importance to study how the dynamics of small solids will be affected in such an unstable environment, and what implications there will be on the collision and coagulation process.  In general, in a non-turbulent, freely rotating gaseous disk at hydrostatic equilibrium, there is a radial gradient associated with the gas pressure. This pressure gradient counteracts the gravitational attraction of the central star and causes the gas molecules to have slightly different velocities than Keplerian circular. When the pressure gradient is positive, the velocity of a gas molecule is greater than the local Keplerian velocity. A solid in the gas, in this case, feels an acceleration by the gas along its orbit and, consequently, the increase in its orbital angular momentum forces the solid to a larger orbit. In this case, we say that the solid feels a ``tail wind.'' The opposite is true  when the pressure gradient is negative. That is, a solid body will be subject to a ``head wind'' and will migrate toward smaller orbits.  One of the features of a rotating gravitationally unstable disk is the appearance of spiral arms or clumps where the density of the medium is locally enhanced. In the vicinity of such density enhancements, the pressure of the gas may change radially and cause the particles in the disk to migrate toward the location of the maximum gas density. We are interested in studying the dynamics of solids that undergo such migration and in exploring the possibility of applying the results to the formation of planetesimals in a marginally- gravitationally unstable disk. As the first stage of our project, we present here the results of a systematic study of the migration of solids subject to gas drag and the gravitational force of a circumstellar disk, around the location of its maximum density. To focus attention on the dynamics of the solids and its association with parameters such as the temperature of the gas, the sizes of the objects, and also the values of their densities, we consider a hypothetical solar nebula with a circularly symmetric density function.  Studies of the motions of solids in  gaseous mediums have been presented by many authors. In a detailed analytical analysis in 1962, Kiang studied the dynamical evolution of solids in elliptical orbits subject to resistive forces proportional to arbitrary powers of their relative velocities and their distances to the star. In his study, Kiang considered three cases of stationary, uniformly rotating, and also freely rotating gaseous mediums. However, he did not consider the pressure gradient of the gas. It was  \\citet{Whip64} who first mentioned that the rotation of the solar nebula deviates from Keplerian because of counterbalancing the gravity of the Sun by the internal pressure of the gas which in turn results in in/outward migration of small solids. In 1972, Whipple studied the dynamics of such solids in the solar nebula where, following an approximation by \\citet{ProbFas69}, he also included the resistive effect of the gas. Whipple's work was subsequently expanded upon and generalized by \\citet{Weid77} for a variety of model nebulae and different sizes of solids.  A comprehensive study of the effect of gas drag on the motions of solid bodies can also be found in the classic work of \\citet{Ada76}. In their paper, Adachi et al. studied the motion of a solid on an elliptical orbit in a solar nebula whose density and temperature vary inversely with different powers of the distance from the Sun. They also presented a detailed analysis of the form of the gas drag for different relative velocities and relative sizes of solids, and also different values of the gas Reynolds number.  Among the more recent studies of the dynamical evolution of solids in a gaseous disk, one can name the work of \\citet{Rano93} on resonance capture of planetesimals subject to a drag force proportional to their relative velocities, as a barrier for the inward flow of solids to the accretion zone of a planetary embryo in the solar nebula, a paper by \\citet{Sup00} on the formation of icy planetesimals subject to a linear combination of Stokes and Epstein drags in an azimuthally symmetric, turbulent and thin solar nebula with a polytropic equation of state, and also a paper by \\citet{Iwas01} on the stability/instability of protoplanets subject to gas drag.  In this paper, we study the dynamics of solid bodies in an inhomogeneous gaseous disk. Our model nebula consists of a Sun-like star at its center and non-interacting collisionless bodies scattered on its midplane. We consider the effect of drag force  and also include the gravitational attraction of the nebula.  The outline of this paper is as follows. Section 2 introduces the equations of motion and also the basic relations concerning drag and the gravitational force of the gas. Section 3 defines the system of interest, and section 4 presents the results of our numerical simulations. Section 5 concludes this study by reviewing the results and discussing their applications. ",
        "conclusions": ""
    },
    "0207/astro-ph0207459_arXiv.txt": {
        "abstract": "{ Two independent groups (Giveon et al. \\cite{giveon}; Mart\\'{\\i}n-Hern\\'{a}ndez et al. \\cite{martin}) have recently investigated the Galactic metallicty gradient as probed by ISO observations of mid-infrared emission-lines from HII regions. We show that the different gradients inferred by the two groups are due to differing source selection and differing extinction corrections. We show that both data sets in fact provide consistent results if identical assumptions are made in the analysis. We present a consistent set of gradients in which we account for extinction and variation in electron temperature across the Galactic disk. ",
        "introduction": "In two recent studies Giveon et al. (\\cite{giveon}) and Mart\\'{\\i}n-Hern\\'{a}ndez et al. (\\cite{martin}) analyzed mid-infrared fine-structure emission lines, as observed by the {\\it Infrared Space Observatory} (ISO), to study the excitation and metallicity of HII regions across the Galactic disk. Discrepancies appear to be present in the gradients published in these two papers. In this note we re-analyze the reduced data compiled by Giveon et al. and  Martin-Hernandez et al. and we demonstrate that similar results are obtained if the same source samples are chosen, and no assumptions are made for the extinction correction. We then apply extinction and electron temperature corrections to both data sets and infer consistent Galactic abundance gradients from the two independent ISO studies. ",
        "conclusions": ""
    },
    "0207/astro-ph0207490_arXiv.txt": {
        "abstract": "The cooling of neutron stars by URCA processes in the kaon-condensed neutron star matter for various forms of nuclear symmetry energy is investigated. The kaon-nucleon interactions are described by a chiral lagrangian. Nuclear matter energy is parametrized in terms of the isoscalar contribution  and the nuclear symmetry energy in the isovector sector. High density behaviour of nuclear symmetry energy  plays an essential role in determining the composition of the kaon-condensed neutron star matter which in turn affects the cooling properties. We find that the symmetry energy which decreases at higher densities makes the kaon-condensed neutron star matter fully protonized. This effect inhibits strongly direct URCA processes resulting in  slower cooling of neutron stars as only kaon-induced URCA cycles are present. In contrast, for increasing symmetry energy direct URCA processes are allowed  in the almost whole density range where the kaon condensation exists. ",
        "introduction": "The possibility of the kaon condensation in dense nuclear matter, proposed by Kaplan and Nelson \\cite{Kaplan:1986yq}, is recently a subject of intensive theoretical   and experimental research \\cite{Lee:1996ef,Li:1997zb}. Such a kaon-condensed phase of dense matter is of direct astrophysical relevance as it would form in neutron stars and could affect their properties. In this paper we study how sensitive is the cooling of neutron stars to the presence of the kaon condensates. A particularly important problem in this regard, which we address here, is to assess the influence of the uncertainty of the high density behaviour of nuclear symmetry energy on the charged kaon condensation itself. {Other strange hadrons, like neutral kaons and hyperons could also appear in neutron star matter. Their inclusion here, however, would blur the role of the nuclear symmetry energy as it would introduce additional uncertainty due to poorly known nucleon-hyperon interactions. The analysis reported here will be extended to account for hyperons in a forthcoming paper. Model calculations show that hyperons are likely to influence  the results concerning kaon condensation \\cite{Knorren:1995ds}.}  Among theoretical approaches we may distinguish those based on the chiral  theory \\cite{Kaplan:1986yq} and the ones using the meson exchange picture. The former ones have their roots in the spontaneously broken $ SU(3)\\!\\times\\!SU(3) $ symmetry. It gives a solid base to study kaon-nucleon interactions where kaons are treated as pseudo-Goldstone bosons. In the meson exchange picture, the kaon couplings  to nucleons are realized indirectly through the exchange of vector and scalar mesons \\cite{Glendenning:1999ak}. The predictive power of chiral theory is lost in such approaches. In order to study the kaon condensation in the neutron star matter a realistic model of nucleon-nucleon interactions should be considered.  The nucleon-nucleon part of the interaction Hamiltonian is usually taken from other theories of nuclear matter, such as e.g. many-body models with {\\em realistic potentials} or relativistic mean field (RMF) theory. Realistic potential approach is firmly based on experiment. Two-nucleon potentials are fitted to the scattering data. Three-body potentials must be added to get agreement with the binding energy of light nuclei ($^3$H, $^4$He) and the saturation properties of nuclear matter. The  energy of nuclear matter is usually derived in a variational approach  \\cite{Wiringa:1988tp} or by using other techniques as e.g. some version of the  Brueckner theory.  The RMF approach starts from postulating the Lagrangian for nucleons and mesons which takes into account relevant internal symmetries (as e.g. isospin). The meson-nucleon interactions are assumed to be of a Yukawa type with coupling constants treated as free parameters which are fitted to the saturation point properties obtained from the mass formula and from breathing modes of giant resonances \\cite{glend-book}.  {Unfortunately, up to now there is no consistent description of the ground state properties of nuclear matter in terms of the chiral theory with its free-space parameters. Recently some novel approaches  were proposed. One of them starts from the chiral Lagrangian \\cite{Mao:1998nv} but still treats its crucial parameters in the spirit of RMF model -- fitting them to the saturation point properties. A different idea based on the power counting  scheme was presented in \\cite{Kaiser:2001jx}, but its extrapolation to higher densities seems not to be well justified.}  In the next section we discuss the high density behaviour of the nuclear symmetry energy. One should stress that the symmetry energy governs the chemical composition of  nuclear matter in the beta equilibrium and hence it determines the particle species present in  neutron stars. In Sect.\\ref{sec_model} the chiral model kaon-nucleon interactions is derived. In Sect.\\ref{sec_thermod} the kaon condensate properties and the role of nuclear symmetry energy in  a kaon-condensed matter are studied. In particular we assess how uncertainty of the behaviour of the symmetry energy at high densities influences conclusions regarding the formation and properties of the kaon-condensed neutron star matter. Implications  for the cooling of neutron stars are investigated in Sect.\\ref{sec_cooling}. ",
        "conclusions": "We have shown that the formation and properties of the kaon-condensed phase of neutron star matter are quite sensitive to the high density behaviour of the nuclear symmetry energy, $E_s(n)$, which still remains the most poorly known property of dense matter. In particular, for $E_s$ decreasing to negative values the formation of the kaon condensate can be inhibited for the lowest absolute value of $a_3m_s$. This is in stark contrast to the results  for monotonically increasing $E_s(n)$, when the condensation occurs for any value of $a_3m_s$. The properties of the kaon-condensed neutron star matter and the abundance of particle species are sensitive to the high density form of the nuclear symmetry energy, as shown in Figs.\\ref{xxallav.eps}-\\ref{xxalltni.eps}. A surprising finding is that the direct URCA process is often not allowed in the kaon-condensed matter for decreasing $E_s(n)$.  The above results allow us to assess how great is the influence of nuclear force models on the total neutrino emissivity of kaon-condensed neutron star matter. In case of increasing symmetry energy as in the RMF model $\\mathrm{(E_1)}$ dURCA operates almost always when kaon condensate is formed, except of a narrow zone around density where the chemical potential $\\mu$ is zero. For realistic potential models dURCA is systematically suppressed  to more narrow density range with decreasing symmetry energy. For an extreme case of the UV14+TNI interactions, the allowed zone shrinks eventually to a very narrow vicinity of some density close to the threshold value for condensation $n_c$. The kaon-induced URCA branch is always present but at the level of about 10 times smaller, as shown in Fig.\\ref{intallmu.eps}.  Another interesting  question concerns the role of $\\bar{K}^0$. As was shown in \\cite{Pal:2000pb} neutral kaons are easily produced just after the onset of $K^-$ condensation. Then the kaon-condensed matter becomes even more effectively isospin-symmetrized than in the absence of $\\bar{K}^0$. We would like to stress that this effect is sensitive to the form of the nuclear symmetry energy. The effect occurring for increasing $E_s(n)$ is not observed in the case of decreasing $E_s(n)$. High values of $x$ which we obtain in our model make the $\\bar{K^0}$ condensation  more difficult. The dispersion relations for anti-kaons are: \\beq \\omega_{K^-} = - \\frac{n(1+x)}{4 f^2} + \\left(\\left(\\frac{n(1+x)}{4 f^2}\\right)^2 + m_K^2 + (2a_1 x + 2a_2 + 4 a_3) m_s \\frac{n}{2f^2}  \\right)^{1/2} \\eeq \\beq \\omega_{\\bar{K}^0} = - \\frac{n(2-x)}{4 f^2} + \\left(\\left(\\frac{n(2-x)}{4 f^2}\\right)^2 + m_K^2 + (2a_1 (1-x) + 2a_2 + 4 a_3) m_s \\frac{n}{2f^2}  \\right)^{1/2} \\label{omega_k0} \\eeq The first term in the expression (\\ref{omega_k0}) for $\\omega_{\\bar{K}^0}$, which is the leading term, includes positive contribution from $x$, so high values of $x$  make the slope of $\\omega_{\\bar{K}^0}$, as a function of $n$, more flat after the charged kaons appear. Fig.\\ref{omega_rys} presents this effect clearly. We leave more detailed discussion of the neutral kaon condensate for future work.  \\begin{figure}[h] \\center{\\includegraphics[height=3in]{omegax_av1.eps}} \\caption{The proton fraction $x$ and anti-kaons energy in dense matter. The leveling off of $\\omega_K$ for ${\\bar{K}^0}$ corresponds to the onset of $K^-$ condensation which drives the strong increase of the proton fraction shown in the upper panel.} \\label{omega_rys} \\end{figure}  Note added: After completion of this work dr. T. Muto brought to our attention his research \\cite{Fujii:ua} concernig  kaon condensation and cooling of neutron stars. \\newpage  \\appendix"
    },
    "0207/astro-ph0207203_arXiv.txt": {
        "abstract": "We have obtained an HI absorption spectrum of the relativistic binary PSR~J1141$-$6545 and used it to constrain the distance to the system.  The spectrum suggests that the pulsar is at, or beyond, the tangent point, estimated to be at 3.7 kpc.  PSR~J1141$-$6545 offers the promise of stringent tests of General Relativity (GR) by comparing its observed orbital period derivative with that derived from other relativistic observables.  At the distance of PSR~J1141$-$6545 it should be possible to verify GR to an accuracy of just a few percent, as contributions to the observed orbital period derivative from kinematic terms will be a small fraction of that induced by the emission of gravitational radiation.  PSR~J1141$-$6545 will thus make an exceptional gravitational laboratory. ",
        "introduction": "An independent confirmation of General Relativity (GR) has recently been obtained using a measurement of the annual orbital parallax of the binary millisecond pulsar PSR~J0437$-$4715 to derive the inclination angle of the binary, thus predicting the shape of the Shapiro delay \\cite{vbb+01}. In this test, as with those using self-consistency checks, GR has been shown to be a remarkably accurate description of gravity.  Measurement of post-Keplerian ($PK$) orbital parameters in relativistic binary pulsars \\cite{dt92} have provided the most stringent tests of GR to date.  PSR~B1913+16 displays a measurable orbital period derivative ($\\dot{P_b}$), time dilation and gravitational redshift parameter ($\\gamma$) and an advance of periastron ($\\dot{\\omega}$), which are all consistent with GR \\cite{tw89,dt91,tay94}. Unfortunately, at the $\\sim$1\\% level, these tests cannot be authoritative because of our ignorance of the distance and exact contribution of the Galactic potential to the observed orbital period derivative of the binary system. PSR~B1534+12 is a nearby relativistic pulsar which demonstrates all of the above $PK$ phenomena, together with the range ($r$) and shape ($s$) of Shapiro delay which have been shown to be remarkably self-consistent with GR \\cite{sac+98}. As predicted by Bell and Bailes (1996)\\nocite{bb96}, this pulsar could not be used to verify the emission of gravitational radiation at better than the 15\\% level because of the uncertain distance to the pulsar. Instead, Stairs et al.~(1998) have used GR to obtain the pulsar distance to a high degree of accuracy.  The relativistic binary pulsar PSR~J1141$-$6545 was recently discovered in the Parkes multibeam pulsar survey \\cite{klm+00a}. It has a spin period ($P$) of 394~ms, a very narrow pulse ($0.01P$) and is in an eccentric, 4.7~hr orbit. The measured $\\dot{\\omega}$ is 5.3 degrees~yr$^{-1}$ and it has a predicted $\\dot{P_b}$ of $-3.85 \\times 10^{-13}$, which is twice that observed for PSR~B1534+12, but nearly an order of magnitude less than that observed for PSR~B1913+16. Tests of GR using PSR~J1141$-$6545 would nicely complement those already made using the other relativistic systems because it is thought to contain a neutron star and a white dwarf, rather than two neutron stars, as in the PSR~B1534+12 and PSR~B1913+16 systems. Within a few years, the orientation and component masses should be known to high accuracy from measurements of the advance of periastron and $\\gamma$ terms. This will completely specify the orbital geometry of the system and make a specific prediction about the magnitude of the orbital period derivative.  A dispersion measure ($DM$) of 116 pc~cm$^{-3}$ places this pulsar at a distance of 3.2 kpc using the Taylor and Cordes (1993)\\nocite{tc93} model.  These estimates can be significantly in errorr. In individual cases the error in the DM--distance can be as great as a factor of 2. If this distance is correct the kinematic ``contamination'' of the observed orbital period derivative would be small.  Hence, PSR~J1141$-$6545 could be an extremely good gravitational laboratory.   The detection of neutral hydrogen absorption features in the spectrum of a pulsar, together with a detailed neutral hydrogen emission spectrum in the same direction, can be used to place constraints upon its distance.  A discussion of the neutral hydrogen distance determination method is given by Frail and Weisberg (1990)\\nocite{fw90} and the most recent review is by Weisberg (1996)\\nocite{wei96}. In this paper we use neutral hydrogen measurements to constrain the distance of PSR J1141$-$6545 and show that it should be an exceptional laboratory for testing GR. ",
        "conclusions": "The measurement of the absorption spectrum of this pulsar has allowed a lower limit to be placed upon its distance of 3.7 kpc, the tangent point distance predicted by the Galactic rotation curve. The mean electron density in the direction of this pulsar is therefore at most 0.03 cm$^{-3}$. We have also examined the Galactic and kinematic contribution to the observed $\\dot{P_b}$ and found that it should only be at the few percent level for reasonable pulsar transverse velocities making PSR~J1141$-$6545 an excellent gravitational laboratory. Future measurements of the system's proper motion will help to further define the level at which this system can be used to test GR."
    },
    "0207/astro-ph0207035_arXiv.txt": {
        "abstract": "We have developed a model for molecular hydrogen formation under astrophysically relevant conditions. This model takes fully into account the presence of both physisorbed and chemisorbed sites on the surface, allows quantum mechanical diffusion as well as thermal hopping for absorbed H-atoms, and has been benchmarked versus recent laboratory experiments on \\hm\\ formation on silicate surfaces.  The results show that \\hm\\ formation on grain surface is efficient in the interstellar medium up to some 300K. At low temperatures ($\\leq$100K), \\hm\\ formation is governed by the reaction of a physisorbed H with a chemisorbed H. At higher temperatures, \\hm\\ formation proceeds through reaction between two chemisorbed H atoms. We present simple analytical expressions for \\hm\\ formation which can be adopted to a wide variety of surfaces once their surfaces characteristics have been determined experimentally. ",
        "introduction": "Molecular hydrogen is the most abundant molecule in the universe and dominates the mass budget of gas in regions of star formation. \\hm\\ is also an important chemical intermediate in the formation of larger species and can be an important gas coolant in some conditions, particularly in the early universe. Yet, despite its importance, the \\hm\\ formation process is still not well understood. Observationally, it has been shown that \\hm\\ can be efficiently formed over a wide range of temperatures ( Jura 1974; Tielens \\& Hollenbach 1985a, 1985b; Hollenbach \\& McKee 1979). Theoretically, Gould and Salpeter (1963) showed the inefficiency of \\hm\\ formation in the gas phase and concluded that recombination of physisorbed H on ``dirty'' ice surfaces is efficient between 10 and 20K. Hollenbach and Salpeter (1970; 1971), recognizing that this small temperature range presents a problem, considered reactions involving H atoms bound to dislocations and impurities with energy exceeding normal physisorbed energies, and obtained a recombination efficiency $\\simeq$ 1 up to a critical temperature between 25 and 50K. Goodman (1978) calculated the quantum and thermal mobility of the atoms on graphite grains assuming that these atoms could only be physisorbed. Many studies focussed for various reasons on icy surfaces where H is physisorbed (Buch \\& Zhang 1991; Takahashi et al. 1999). However, most astrophysically relevant material (e.g., silicates, graphite) can bind H also in chemisorbed sites (Barlow \\& Silk, 1976, Aronowitz \\& Chang 1980; Leitch-Devlin \\& Williams 1984; Klose 1992; Fromherz et al. 1993). On these surfaces, binding can occur in a physisorption layer (E$\\sim$500K) at a distance of some Z$\\sim$3$\\AA$ as well as in a chemisorption layer (E$\\sim$10000K) deeper into the surface Z$\\sim$0.5$\\AA$ (Barlow \\& Silk, 1976; Zangwill, 1988). Recently, Katz et al. (1999) developed a model for \\hm\\ formation benchmarked by two sets of experiments. This model considers the atoms bound only in physisorbed sites and diffusing only thermally on the surface, colliding and recombining to form molecules. As for the ice models studied earlier, this allows molecule formation only below 15K for olivine grains and 20K for carbon grains which contradicts the ISM observations. Perusing these models, it is clear that the process of the \\hm\\ formation is governed by the binding of atomic H to the surface and the concomitant mobility of these atoms (Leitch-Devlin \\& Williams 1984; Tielens \\& Allamandola 1987). Surface diffusion can occur through quantum mechanical tunneling (at low temperatures) and through thermal hopping (at high temperatures). For a proper description of molecular hydrogen formation in the ISM both types of binding diffusion processes have to be taken into account (Cazaux \\& Tielens 2002). ",
        "conclusions": "Our study reveals the presence of two distinct regimes of \\hm\\ formation, which reflect directly the presence of two types of atomic H binding sites. At low temperature (T$\\leq$100K), \\hm\\ formation involves migration of physisorbed H atoms. At higher temperatures (T$\\geq$100K), \\hm\\ formation results from chemisorbed H recombination.  The presence of these two types of binding sites allow \\hm\\ formation to proceed relatively efficiently even at elevated temperatures. The study of Hollenbach and Salpeter (1971) focused on icy surfaces on which H can only physisorb. As a result, \\hm\\ formation ceased at temperatures in excess of $\\sim$20K. Recognizing this problem, Hollenbach and Salpeter involved the presence of enhanced binding sites on the ice with ill-determined parameters. These sites allowed \\hm\\ formation to proceed up to some 75K. Since their study, it has become abundantly clear that interstellar grains are not covered by ice in the diffuse interstellar medium (Whittet et al. 1983, 1988). Silicate and graphitic surfaces are now widely accepted as astrophysically relevant grain surfaces (Mathis 1991) and those surfaces intrinsically possess enhanced binding sites; e.g., chemisorbed sites. The parameters of these chemisorbed sites have not yet been well determined because experiments have focused on low temperature \\hm\\ formation (Pirronello et al. 1997; Katz et al. 1999). The values adopted in this study are however quite reasonable and illustrate the efficiency of \\hm\\ formation at elevated temperatures well. When future experiments determine the values of the parameters involved ($E_{H_C}$, $E_{H_P}$, $\\mu$), the results can be directly adjusted. Inside dense clouds, interstellar grains are covered by ice. Of course, in such environments, almost all hydrogen is already in molecular form and \\hm\\ formation is perhaps only of academic interest. Nevertheless, we note that in such environments, molecular hydrogen formation may proceed mainly through H-abstraction from molecules such as H$_{2}$S and N$_{2}$H$_{2}$ (Tielens and Hagen 1982).  In a sense, these species act as ``chemisorption'' sites for hydrogen. Migrating atomic H may tunnel through the reaction barriers involved and form \\hm. Eventually, these ice covered grains are transported back into the diffuse ISM, when the molecular cloud is disrupted. Photodesorption and sputtering in strong shocks quickly remove their ice on a timescale of some $10^6$ yr (Tielens and Hagen 1982, Jones et al. 1994, Draine and Salpeter 1979, Barlow 1978). At that point, molecular hydrogen formation is again governed by the properties of bare grain surfaces. Similarly, any (thin) layer accreted in the diffuse interstellar medium will be quickly sputtered in even a modest velocity ($\\sim$30km s$^{-1}$) shock (Jones et al. 1994).\\\\Finally, the formation efficiency of molecular hydrogen will also depend on the sticking coefficient of H atoms colliding with the grain. In our model and the formulae derived, the sticking coefficient is subsumed in the incident flux, F. Astrophysical studies of sticking of H on grain surfaces have concentrated on physisorbed interactions and the sticking coefficient is $\\sim$1 at low temperatures and decreases with increasing temperature to about 0.3 at T=300K ( Hollenbach \\& Salpeter 1970; Hollenbach \\& McKee 1979; Burke \\& Hollenbach 1983; Leitch-Devlin \\& Williams 1985). However, if the interaction occurs through much stronger chemisorption, then the sticking coefficient might be large even at high temperatures (Tielens \\& Allamandola 1987; Duley \\& Williams 1984).   \\subsection{Summary and conclusions} Recently, we have modelled molecular hydrogen formation on grain surfaces. This model consider hydrogen atoms bound to the surface at two energy levels (i.e. chemisorption and physisorption). The H mobility from one site to another is a combination of tunelling effect and thermal diffusion. This model has been experimentally benchmarked (Pirronello et al. 1997a, 1997b, 1999) and the relevant surface characteristics have been determined. These characteristics allow us to extend our model for \\hm\\ formation under astrophysically relevant conditions. The results show efficient \\hm\\ formation from $\\sim$6K to $\\sim$300K. The different processes involved in \\hm\\ formation at different temperatures has been discussed.  Until about 100K, \\hm\\ forms by recombination of a physisorbed H with a chemisorbed H and is highly efficient. At higher temperatures, when physisorbed atoms evaporate quickly, the recombination of two chemisorbed atoms is required to form \\hm. \\hm\\ formation is then less efficient, $\\epsilon_{H_2}$$\\sim$0.2. The parameters involved in H chemisorption and \\hm\\ formation at high temperatures are presently not well known. The adopted values are very reasonable and the gross characteristics -- \\hm\\ formation at high temperatures -- are undoubtedly correct. Nevertheless, future experiments are very important to determine the maximum temperature to which \\hm\\ formation in the ISM can occur."
    },
    "0207/astro-ph0207567.txt": {
        "abstract": "We have modeled the infrared (IR) Spectral Energy Distribution (SED) of circumstellar disks embedded in a HII region and photoevaporated by the external ultraviolet radiation. The model applies to the photoevaporated disks (proplyds) in the Orion Nebula, most of them illuminated by the O6.5 star \\thC.  First we calculate the IR emission of a Pre-Main-Sequence star surrounded by a dusty globule that is immersed within an HII region. The globule is assumed to be spherical, homogeneous, optically thin at IR wavelengths and photoevaporated according to the Dyson (1968) model. Second, we consider the IR emission of a disk directly exposed to the nebular environment. The reprocessing disk is passive and treated according to the Chiang and Goldreich (1997, CG97) model. We improve over the CG97 treatment by tracing the propagation of the various radiative fluxes (from the star exciting the HII region, nebular, and grazing from the disk central star) through the disk superheated atmosphere. Since the opposite disk sides receive different amounts of radiation, the flaring angle and the surface temperature distributions are different, resulting in well distinguished SEDs for the two disk faces. % and possibly disk warping. Finally, we combine the globule and disk models to estimate the IR emission of proplyds. The energy input from the central star and the nebular environment increase the disk flaring angle, and therefore also the amount of stellar radiation intercepted by the disk. %due to the dust in the evaporated envelope and in the disk's interior. The relative intensity of the disk vs. envelope emission varies with the tilt angle relative to the directions of \\thC\\, and of the Earth. We explore the dependence of the SEDs upon the tilt angle with respect to the Earth, the distance from \\thC, the size on the envelope, the inner disk radius and the temperature of the central star. The resulting SEDs are characterized by a broad peak of emission at 30-60\\,\\micron\\, and are in general significantly different from those of isolated disks in low-mass star forming regions like Taurus-Auriga. Our model indicates that in the presence of an external radiation field, relatively evolved \"Class 2\" objects may display a SED peaking at mid-IR and far-IR wavelengths. Also, the model can account for  the strong mid-IR excess we have recently detected at 10\\,\\micron\\, from embedded disks in the Orion Nebula. ",
        "introduction": "Circumstellar disks are thought to be the birthsites of planetary systems, and it is primarily for this reason that their properties have been studied extensively since the early 1980\u0092s (Shu, Adams, and Lizano 1987; Beckwith \\& Sargent 1996; Hartmann 1998).  The discovery of several tens of extra-solar planetary systems (e.g. Marcy, Cochran \\& Mayor 2000) added weight to the argument that disks create planets, although there is some debate about how often planetary systems really occur around stars.  Nevertheless, an understanding of the evolutionary properties of disks around young stars is thought to be an essential part of the understanding of our own origins.  Disks can also affect the early evolution of stars themselves.  Several authors have suggested that stellar rotation is locked to the orbital rates of the inner disks at early times allowing the disk to regulate the stellar angular momentum (Edwards et al. 1993, Choi \\& Herbst 1996), although this idea is still controversial (Stassun et al. 1999, Rebull 2001). Typical disk masses and accretion rates appear to be too small to greatly change the mass of the stars through accretion (Palla \\& Stahler, 1999). A number of observable characteristics of young star/disk systems play an important role in determining the stellar properties %rotation, for example, so and therefore understanding even subtle changes is important to ensure that the stellar properties are correctly interpreted.  Most of the stars in the Galaxy are thought to be born in dense OB clusters such as the Orion Nebula (Bally et al. 1998a; McCaughrean \\& Stauffer 1994).  However, our understanding of circumstellar disks comes mainly from studies of nearby dark clouds such as the Taurus/Auriga complex, and it is not clear that the extant knowledge of circumstellar disks applies to the majority of young star/disk systems in the Galaxy.  The study of the disks in Orion presents several problems that has slowed progress: they are seen against the Orion HII region, a bright, highly nonuniform source of radiation at wavelengths from the ultraviolet through millimeter; they appear to be fainter on average at millimeter wavelengths (Bally et al. 1998b); and the Orion nebula is three times farther from Earth than the well-studied dark clouds. Many of the disks are embedded within the HII region, meaning that they are surrounded by ionized halos of gas and dust that complicate the observation of disk properties alone.  To address these difficulties, we present in this paper a series of models of the emission from disks embedded within the Orion HII region, under the influence of other radiation sources in addition to the stars at their center. Section 2 is a discussion of various sources of radiation in the HII region to prepare for the calculations of emission from discrete sources.  In the following sections, we show how the different radiation affects the IR emission arising from three type of sources: a spherically symmetric, dusty globule surrounding a star (section 3), a star/disk system directly exposed to the ionized environment (section 4), and finally a star/disk system surrounded by a dusty globule (section 5).  Section 6 contains a discussion of our findings with an exploration of the critical parameters, comparison with observations, and remarks on the limit of our treatment. ",
        "conclusions": "We have explored the IR emission of circumstellar disks in the environment where star formation most typically occurs, i.e. a HII region powered by massive OB stars. This scenario applies in particular to the Orion Nebula, where the interaction of circumstellar disk with the environment has  been directly resolved by the HST. We have build our model considering the four types of radiation relevant to the dust heating in a HII region, e.g ionizing (EUV) and non-ionizing (FUV) flux from the exciting star (\\thC), resonant \\Lya\\, radiation and the remaining nebular radiation. We evaluated how these radiation sources affect the IR SED arising from: \\begin{enumerate} \\item A spherical, homogeneous, optically thin, circumstellar globule, photoevaporated by the UV radiation of the HII region exciting star. The globule has a neutral core of uniform density, and the photoevaporated atmosphere is treated following the Dyson (1968) model.  With our assumptions, the most important parameter constraining the IR emission, the radial optical depth to the UV radiation, depends only on the distance from the ionizing star and on the globule size. The thermal emission peaks in the range $10 - 30$\\,\\micron, depending mostly on the dust composition. \\item A non-accreting disk directly exposed to the nebular radiation. The disk is in hydrostatic and radiative equilibrium, and treated following the prescription of Chiang and Goldreich (1997). We modify the CG97 scheme to account for the fact that the nebular radiation hits the disk surface with large angles. We follow the propagation of the various radiative fluxes  through the disk superheated atmosphere, deriving the vertical temperature profile of the disk atmosphere. The disk faces receive unequal amount of radiation and present different flaring angles. In particular, the flux emitted from the face opposite to the ionizing star provides a model to the IR emission of the dark silhouette disks observed in the Orion Nebula. \\item A combined system composed by the disk and the photoevaporated envelope, i.e. a photoionized proplyd of the type observed in the immediate surroundings of  \\thC\\, in the Orion Nebula. \\end{enumerate} Depending on the distance and on the tilt angle of the disk with respect to the ionizing star, the disk flaring may be substantially higher than in the case of isolated disks. The tilt angle with respect to the Earth plays also a major role by hiding the central parts of the disk. The relative intensity of the disk vs. envelope emission varies with the tilt angle to the direction of the Earth. The high temperatures reached by the dust either at the disk atmosphere or within the envelope produce a SED peaking at $30-60$\\,\\micron.  We explore the dependency of the SEDs upon the tilt angle with respect to the Earth, the distance from the ionizing star, the size on the envelope, the inner disk radius, and the temperature of the disk's central star. The resulting SEDs are characterized by a broad peak of emission at 30-60\\,\\micron\\, and are in general significantly different from those of isolated disks in low-mass star forming regions like Taurus-Auriga. Our model indicates that in the presence of an external radiation field, relatively evolved ``Class 2\" objects may display a SED peaking at mid-IR and far-IR wavelengths. The model explains the strong mid-IR excess we have recently detected on several sources in a 10\\,\\micron\\, survey of the Orion Nebula."
    },
    "0207/astro-ph0207639_arXiv.txt": {
        "abstract": "Several of the inconsistencies plaguing the field of novae are resolved once we consider novae  to be steady state super-Eddington objects. In particular, we show that the super-Eddington shell burning state is a natural consequence of the equations of stellar structure, and that the predicted mass loss in the super-Eddington state agrees with nova observations. We also find that the transition phase of novae can be naturally explained as ``stagnating\" winds. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207313_arXiv.txt": {
        "abstract": "We present a high-resolution ($\\sim$5'') image of the nucleus of M\\,82 showing the presence of widespread emission of the formyl radical (HCO). The HCO map, the first obtained in an external galaxy, reveals the existence of a structured disk of $\\sim$650\\,pc full diameter. The HCO distribution in the plane mimics the ring morphology displayed by other molecular/ionized gas tracers in M\\,82. More precisely, rings traced by HCO, CO and HII regions are nested, with the HCO ring lying in the outer edge of the molecular torus. Observations of HCO in galactic clouds indicate that the abundance of HCO is strongly enhanced in the interfaces between the ionized and molecular gas. The surprisingly high overall abundance of HCO measured in M\\,82 (X(HCO)$\\sim$4\\,10$^{-10}$) indicates that its nuclear disk can be viewed as a {\\it giant} Photon Dominated Region (PDR) of $\\sim$650\\,pc size. The existence of various nested gas rings, with the highest HCO abundance  occurring at the outer ring (X(HCO)$\\sim$0.8\\,10$^{-9}$), suggests that PDR chemistry is {\\it propagating} in the disk. We discuss the inferred large abundances of HCO in M\\,82 in the context of a starburst evolutionary scenario, picturing the M\\,82 nucleus as an evolved starburst. ",
        "introduction": "M\\,82 is the closest galaxy experiencing a massive star formation episode \\citep{rie80, wil99}. Its nuclear starburst, located in the central 1\\,kpc, has been the subject of continuum and line observations made in virtually all wavelengths from X-rays to the radio domain. These studies indicate that the high rate of supernova explosions and the strong UV radiation fields have heavily influenced the physical properties and kinematics of the interstellar medium in M\\,82. The high supernova rate has created a biconical outflow of hot gas (Bregman, Schulman, \\& Tomisaka 1995; Shopbell \\& Bland-Hawthorn 1998) also observed in the cold gas and dust (Alton, Davis, \\& Bianchi 1999; Seaquist \\& Clark 2001). The discovery of a $\\sim$500\\,pc molecular gas chimney and a giant supershell in M\\,82, detected in SiO, indicates the occurrence of large-scale shocks in the disk-halo interface of the starburst \\citep{gbu01}. Furthermore, there are evidences that the strong UV-fields have created a particular physical environment in the molecular gas reservoir of M\\,82 \\citep{stu97,mao00,wei01}. The common picture emerging from these studies is that the bulk of CO emission in the nuclear disk of M\\,82 comes from moderately dense (with n(H$_2$)$\\sim$10$^{3}$-10$^{4}$cm$^{-3}$) Photon Dominated Regions (PDR). However, these conclusions are model-dependent and not free from internal inconsistencies \\citep{mao00}.  Observational evidence supports that the emission of the formyl radical (HCO) mainly arises from PDR at the interfaces between the ionized and the molecular gas in our Galaxy. After the first detection of HCO  by Snyder, Hollis, \\& Ulich (1976) ulterior searches have confirmed that this radical is associated with regions where chemistry is driven by an enhanced UV radiation field (Hollis \\& Churchwell 1983; Snyder, Schenewerk, \\& Hollis 1985; Schenewerk et al.1988; Schilke et al. 2001). On the theoretical side, several chemical models have been put forward to account for the observed abundances of HCO in PDR (de Jong, Boland, \\& Dalgarno 1980; Leung, Herbst, \\& Huebner 1984; Schilke et al. 2001).     We have chosen HCO, a privileged tracer of PDR chemistry, to investigate the influence of the M\\,82 starburst on its molecular gas reservoir. \\citet{sag95} tentatively detected the emission of HCO at 3mm in M\\,82, using single-dish observations. However, the HCO lines appear blended with SiO and H$^{13}$CO$^{+}$ lines in their spectrum. Furthermore, the low spatial resolution (50'') of their single-pointed map did not allow to infer the spatial distribution of HCO. In this Letter we present a high-resolution ($\\sim$5'') image of the emission of HCO in the nucleus of M\\,82. The interferometer HCO map of M\\,82, the first obtained in an external galaxy, shows unambiguous evidence that the whole nuclear disk has become a {\\it giant} PDR of $\\sim$650\\,pc size with a total HCO abundance of $\\sim$4\\,10$^{-10}$. The enhancement of HCO presents spatial variations which depend on the distance to the most prominent HII regions of the M\\,82's starburst. ",
        "conclusions": "The enhancement of HCO in PDR is mostly an observational fact, reproduced with uneven success by models. PDR models published so far have tried to account for the derived abundance of HCO using either gas-phase schemes \\citep{jon80,leu84} or incorporating dust grains to the chemistry \\citep{sch01}. According to \\citet{jon80}, the C$^+$ ion, highly abundant in UV-processed cloud envelopes, would start the chain of reactions leading to CH$_2$ and finally, to HCO. Confirming these expectations, \\citet{hol83} found an empirical correlation between HCO and C$^+$ radio recombination line emission in two galactic clouds. \\citet{sch01} have proposed a model where HCO is the final by-product of the photodesorption or evaporation of solid formaldehyde on dust grain mantles \\citep{wes95}.  Although the values of X(HCO) predicted by \\citet{sch01} are still one order of magnitude below what is typically found in PDR, these results are encouraging. Based on a 1-dimensional model made for the Orion molecular cloud, \\citet{sch01} predict that the large photodissociation rate of HCO \\citep{dis88} could be counterbalanced only for extinctions A$_v>$5-6. For a population of clouds uniformly bathed in the pervasive UV field of a starburst galaxy, the A$_v$ limit should be risen by a factor of $\\sim$a few, making the equivalent condition on column density close to N(H$_2$)$>$10$^{22}$cm$^{-2}$.   The HCO map of M\\,82 provides direct evidence that the starburst event has heavily processed the bulk of molecular gas in this galaxy. The global fractional abundance of HCO ($\\sim$4\\,10$^{-10}$, averaged over $\\sim$650\\,pc) can be accounted within a PDR scenario. Paradoxically, the spatial correlation between HCO-enriched clouds and HII regions is poor however. This is noticeable both at large scales (HCO/CO/HII nested ring morphology) but also at small scales (peaks of HCO/HII emission avoid each other). The poor correlation suggests that UV fields coming from the strongest HII regions of M\\,82, lying in the inner $\\sim$400\\,pc, have photodissociated the bulk of HCO in nearby molecular cloud envelopes. Previous observational evidence pointed out to a disrupted physical environment for the dense ISM in M\\,82. Observations of the atomic carbon [CI]\\,$^{3}$P$_{2}\\rightarrow^{3}$P$_{1}$ and $^{3}$P$_{1}\\rightarrow^{3}$P$_{0}$ lines \\citep{sch93,stu97} confirmed that the M\\,82 [CI]/CO abundance ratio ($\\sim$0.5) is higher than observed in non-starburst disks (it is $\\sim$.15 in our Galaxy). Moreover, the measured J=2--1/1--0 ratio can only be reconciled with clouds being small (with sizes $\\sim$1\\,pc), only moderately dense (with densities $\\sim$10$^{3}$-10$^{4}$cm$^{-3}$), and hot (with temperatures $\\sim$50-100K) \\citep{stu97}. Similar conclusions, based on Large Velocity Gradient (LVG) and PDR modelling of $^{12}$CO, $^{13}$CO, and C$^{18}$O line ratios have been reported by \\citet{mao00} and \\citet{wei01}. The bulk of molecular clouds in the inner $\\sim$400\\,pc of M\\,82 might have low densities (at some places in the disk as low as $\\sim$10$^{3}$cm$^{-3}$) and small sizes ($\\sim$1\\,pc). Assuming this scenario, the survival of HCO against photodissociation, requiring column densities N(H$_2$)$>$10$^{22}$cm$^{-2}$ might be difficult. Most remarkably, the location of HCO-enriched clouds in our map, identified by their largest $N(HCO)/N(H^{13}CO^{+})$ ratio (the outer ring and the central clump), coincides with Weiss et al.'s model predictions on where individual molecular clouds may have the largest column densities (where volume densities reach a few $\\sim$10$^{4}$cm$^{-3}$). Furthermore, the intensity of UV fields is there comparatively lower (see Figure 3).   The reported differences between HCO and H$^{13}$CO$^{+}$ maps in M\\,82 cannot be easily explained with the present PDR models however, as both species are expected to be closely associated in PDR. Other unexplored scenarios cannot be excluded as plausible explanations for the enhancement of HCO: X-Ray or Cosmic Ray Dominated Regions chemistries might account for it in the case of M\\,82 (Suchkov, Allen, \\& Heckman 1993). Moreover, we can not exclude the presence of dense molecular gas (with n(H$_2$)$>$10$^{5}$cm$^{-3}$) in the central 1\\,kpc of M\\,82 as already pointed out by \\citet{mao00} and \\citet{wei01}, or shown by the detection of tracers of dense molecular gas (see HCN map of Brouillet \\& Schilke 1993). However, our results suggest that evaporation/destruction of molecular clouds by UV fields is highly efficient in the inner $\\sim$400\\,pc of the galaxy. This efficiency shows spatial trends within the M\\,82's disk but the reasons behind these variations remain to be understood.    M\\,82 is the only galaxy where emission of HCO has been detected so far. HCO emission was also searched for in the starburst galaxy NGC\\,253 by \\citet{gbu00}. Their data, giving no detection, allowed to set a 3-$\\sigma$ upper limit for the  HCO/H$^{13}$CO$^{+}$  ratio of $\\le$ 0.12, i.e. a factor of 4-5 smaller than the value measured in the M\\,82's ring. HCO is the only {\\it complex} molecule showing a larger fractional abundance in M\\,82 than in NGC\\,253 \\citep{mau93}. Furthermore, the remarkably different abundances and spatial distributions of SiO gas in NGC\\,253 and M\\,82 are suggestive of an evolutionary link between these starbursts \\citep{gbu00,gbu01}. The large HCO abundance in M82 may fit within this scenario, which considers the M\\,82 starburst episode as more evolved than in NGC\\,253. The chemistry of molecular gas in NGC\\,253 is heavily influenced by the large-scale shocks and heating induced by a burst of pre-main sequence massive stars. This explains the high abundance of SiO and the significantly low abundance of HCO. The chemistry of molecular gas in M\\,82 is dominated by the action of UV fields produced by more evolved massive stars giving rise to HII regions. The remnant of the starburst presents a PDR-like chemistry favoring the presence of HCO, but forcing a low abundance for other complex molecules."
    },
    "0207/astro-ph0207125_arXiv.txt": {
        "abstract": "We investigate the prospects of detecting weakly interacting massive particle (WIMP) dark matter by measuring the contribution to the extragalactic gamma-ray radiation induced, in any dark matter halo and at all redshifts, by WIMP pair annihilations into high-energy photons. We perform a detailed analysis of the very distinctive spectral features of this signal, recently proposed in a short letter by three of the authors: The gamma-ray flux which arises from the decay of $\\pi^0$ mesons produced in the fragmentation of annihilation final states shows a severe cutoff close to the value of the WIMP mass. An even more spectacular signature appears for the monochromatic gamma-ray components, generated by WIMP annihilations into two-body final states containing a photon: the combined effect of cosmological redshift and absorption along the line of sight produces sharp bumps, peaked at the rest frame energy of the lines and asymmetrically smeared to lower energies. The level of the flux depends both on the particle physics scenario for WIMP dark matter (we consider, as our template case, the lightest supersymmetric particle in a few supersymmetry breaking schemes), and on the question of how dark matter clusters. Uncertainties introduced by the latter are thoroughly discussed implementing a realistic model inspired by results of the state-of-the-art N-body simulations and semi-analytic modeling in the cold dark matter structure formation theory. We also address the question of the potential gamma-ray background originating from active galaxies, presenting a novel calculation and critically discussing the assumptions involved and the induced uncertainties. Furthermore, we apply a realistic model for the absorption of gamma-rays on the optical and near-IR intergalactic radiation field to derive predictions for both the signal and background. Comparing the two, we find that there are viable configurations, in the combined parameter space defined by the particle physics setup and the structure formation scenario, for which the WIMP induced extragalactic gamma-ray signal will be detectable in the new generation of gamma-ray telescopes such as GLAST. ",
        "introduction": "\\label{sec:intro}  The accumulated evidence for the existence of large amounts of nonbaryonic dark matter in the Universe is by now compelling (for a review, see e.g.\\ \\cite{lbreview}). Data on the cosmic microwave background (CMB)\\cite{cmbr} and supernova observations \\cite{supernovae} jointly fix the energy fraction in the form matter and cosmological constant (or something similar) to $\\Omega_M\\sim 0.3$ and $\\Omega_{\\Lambda}\\sim 0.7$, respectively. At the same time, the CMB measurements limit the contribution from ordinary baryons to less than $\\Omega_B\\sim 0.05$, which is in excellent agreement with big bang nucleosynthesis. This means that non-baryonic matter has to make up most of the matter in the Universe, $\\Omega_{DM}\\simeq \\Omega_M$. Incidentally, recent measurements of the large-scale distribution of galaxies independently confirm $\\Omega_M=0.27\\pm 0.06$ \\cite{licia}, giving further credence to these conclusions. The currently best estimate of $\\Omega_M$ comes from a joint analysis of CMB and large scale structure data \\cite{2dF} and gives $\\Omega_Mh^2=0.115\\pm0.009$ where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$.  When it comes to the question of how the dark matter is distributed on small, galactic and sub-galactic, scales the situation is much less clear, however (for a review, see, e.g., \\cite{moorerev}). After being subject to an extensive debate, with both theoretical and observational controversies, it seems that the Cold Dark Matter (CDM) model, with dark matter made of, e.g., weakly interacting massive particles (WIMPs), or the $\\Lambda$CDM model, with a contribution from the cosmological constant, are in fair agreement with current observations, so that drastic modifications like strong self-interaction are not urgently called for (see, e.g., \\cite{primack}).  Focusing on the CDM model with WIMPs as dark matter candidates, there is an obvious interest to use as much as possible of the knowledge of the distribution of CDM given through state-of-the-art N-body simulations. In particular, the distribution of dark matter plays a crucial role in most WIMP detection methods, and determines therefore the possibility of identifying the dark matter and pinning down its particle properties.  In this vein, we recently presented a short note \\cite{beu} (hereafter BEU) where, contrary to previous predictions~\\cite{previous}, it was shown that in the hierarchical picture found in CDM simulations the cosmic $\\gamma$-ray signal from WIMP annihilations may be at the level of current estimates of the extragalactic $\\gamma$-ray flux. In this paper we deal more carefully with the issues of the formation of structure in a CDM or, rather, $\\Lambda$CDM universe, investigating the sensitivity of the expected gamma-ray flux to different treatments of the structure formation process. We also address the question of the diffuse background flux expected from various types of active galaxies and compare its spectral features with those of the signal from WIMP annihilations. We consider several sample cases in a theoretically favored WIMP scenario, that of supersymmetric dark matter, and highlight the possibility to disentangle such signals from the background in future measurements of the the extragalactic $\\gamma$-ray flux, in particular with the GLAST detector. Results for both the signal and background components are presented implementing a careful treatment of the absorption of high energy gamma-rays in the intergalactic space caused by pair production on the optical and infrared photon background.   The paper is organized as follows. In Sec.~\\ref{sec:dm} we set up the general formalism for computing the dark matter induced gamma-ray flux. In Sec.~\\ref{sec:halo} we investigate the properties of dark matter halos, on all scales of relevance to our problem, in various semi-analytical and numerical simulation scenarios. Implications for the WIMP induced gamma-ray flux are discussed in Sec.~\\ref{sec:flux}, while in Sec.~\\ref{sec:bkg}, we investigate the background problem, including the effects of varying within present observational limits the slope of the energy spectrum of the gamma-ray emission from active galaxies. We also check the effects of removing some more resolved point sources, as may be expected for the next generation of experiments. In Sec.~\\ref{sec:applications} we show a few examples of what signals can be expected for one of the favored WIMP candidates, the neutralino, and give an estimate of sensitivity curves for the GLAST detector. Sec.~\\ref{sec:conclusions} contains our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have studied predictions and the observability of the diffuse gamma ray signal from WIMP pair annihilations in external halos. We have found configurations that imply signals at a detectable level for GLAST, the upcoming gamma-ray space telescope, both for non-thermal dark matter neutralinos, such as in the anomaly-mediated supersymmetry breaking model, and for thermal relic neutralinos in the MSSM. The key ingredient to show that detectable fluxes may arise, which was neglected in early estimates, is the picture, inspired by the current theory for structure formation and by N-body simulation results, that dark matter clusters hierarchically in larger and larger halos, with light structures more concentrated than more massive ones. For dark matter candidates in the AMSB scenario our conclusion holds independently of further assumptions on the dark matter distribution inside halos. Pair annihilation cross sections for thermal relic WIMPs are generally smaller; this however can be compensated by the enhancement in the flux one finds if, as suggested by results of simulations, we assume that dark matter profiles are singular and contain small dense substructures.   If the branching ratio for WIMP annihilation into monochromatic gamma rays is significant (about a few times $10^{-4}$ or larger), the induced extragalactic flux shows a very distinctive feature, the asymmetric distortion of the line due to the cosmological redshift and its sudden drop at the value of the WIMP mass. The component with continuum energy spectrum can be at the level of background components but has less distinctive features: the flux is characterized by the ``$\\pi^0$ bump'', rather than by a spectral index, with the peak shifted to energies lower than $M_{\\pi}/2$ and the width set by the WIMP mass. Once a better measurement of the background will be available, it will be possible to address the question of whether or not this signal can be disentangled from other eventual components.  We have discussed in detail how our predictions depend on assumptions on the Cosmological model and the structure formation picture applied. Unless one introduces drastic changes, such as a large warm dark matter component, the cosmological parameters in the CDM setup do not play a major role; results are mainly sensitive to $\\sigma_8$ with about a factor of two uncertainty. Larger indeterminations, of the order of a factor of a few, are introduced when estimating the scaling of the concentration parameter with halo mass, as extrapolations with toy models out of the mass range of N-body simulation results are needed. The functional form of the dark matter profile in single halos introduces a factor of 10 uncertainty, going from the case of a $1/r^{1.5}$ cusp in the Moore profile to the case of non-singular profiles; that uncertainty is much smaller than, e.g., the one induced on the estimate of the flux from the center of our own Galaxy. The presence of substructures inside halos may provide a factor of a few enhancement in the flux, but this effect is more difficult to address: we have presented a simple and rather generic setup, which will be possible to refine when further information on halos will be provided by higher resolution numerical simulations.   Issues related to the estimate of the background are very important as well. We have presented here a novel estimate of the expected background from unresolved blazars in GLAST, exploiting recent data and discussing critically the uncertainties involved, including the role played by gamma absorption.   Concluding, the present analysis has been devoted to examining in detail an idea that three of the authors have recently presented in a short letter~\\cite{beu}. This work provides further support for such proposal, with exciting perspectives for upcoming measurements.  P.U. was supported by the RTN project under grant HPRN-CT-2000-00152. J.E. and L.B. wish to thank the Swedish Research Council for support."
    },
    "0207/astro-ph0207255_arXiv.txt": {
        "abstract": "Short time scale variations have been reported for a few Miras, including sudden 0.2 magnitude or more brightening in the visual lasting from a few hours to a few days.  Are these flashes real?  Are they hot?  What are they?  Almost all the natural time scales for variation in the atmosphere and wind of a Mira variable are months to years long.  What could cause such short-term variations?  One intriguing possibility is that the variations are associated with the interaction of a Jovian planet with the time-dependent outflow of a Mira wind.  Here we discuss observable features of such an interaction in terms of order-of-magnitude estimates and phenomenology, and make one clear prediction requiring observational followup.  Future work needs to include both theoretical calculations and the development of systematic methods for searching for such events. ",
        "introduction": "The light curves of Miras are characterized by large amplitudes (2.5 visual magnitudes or more), long periods (typically a year and ranging from about 100 to over 600 days), and relatively stable periodicity, but with erratic cycle-to-cycle variations and small period changes on decade-long time scales.  The typical light curve has a more rapid rise than its decline and shows nearly straight-line variation in magnitude with time on the declining branch (corresponding to an exponential decay in brightness).  There are often extended ``bumps'' occurring on the rising or descending branches at about the same phase every cycle. The timing and magnitudes of the bumps and of maximum light are typically variable on a scale of $10-20 \\%$, both phenomena most likely associated with the emergence of shocks through the visible-light photosphere (see e.g., \\citet{maf95}) and quite unrelated to the topic of this paper.  These normal variations can be understood, at least qualitatively, based on models such as those by G. H. Bowen (\\citet{bow88}; reviewed by \\citet{wil00}).  The star pulsates radially in a manner similar to that of the Cepheids and RR Lyrae stars, except that in this case it is a combination of hydrogen and helium (first) ionization that provides the ``valve'' for the pulsation (\\citet{woo74}, \\citet{kee70}).  In the atmosphere, these waves build up to become shock waves that propagate out through the atmosphere and into the wind. There are typically two shocks formed per cycle, providing a natural explanation for the ``bumps'' and perhaps also for those objects showing double maxima.  An outflow is generated, either by the pulsation alone or, in some stars, also by radiation pressure on dust grains formed in ``refrigerated'' zones in the atmosphere.  Mira variables have pulsation periods around one year, and this is also about how long it takes material and shocks to travel, at typical speeds of $5-10 \\ km/s$, one to two AU, or about one stellar radius.  This is then the natural time scale for variation.  Other mechanisms that have been suggested as possibly influencing the visual brightness include large-scale convective cells and solar-type flares.  The natural time scale for variations linked to convection will also be of the same order as that for the pulsation. Solar-type flares, although quick, would need to be very bright - orders of magnitude brighter than flares observed on any other single star - to contrast with these very luminous stars. There is no evidence for the kinds of magnetic fields that would be required. (For a discussion of this point, see \\citet{sok02}.)  Not seeing a plausible mechanism for short-time-scale high-contrast events, we have been very slow to accept reports of these, at least until some systematic effects could be established in a homogeneous data set. Such data are now emerging, and with further reflection we are also seeing some interesting possibilities for what the data may be telling us. ",
        "conclusions": ""
    },
    "0207/astro-ph0207549_arXiv.txt": {
        "abstract": "We obtained optical and infrared spectra of Cha H$\\alpha$ 5/cc 1, a faint possibly sub-stellar companion candidate next to the M6-type brown dwarf candidate Cha H$\\alpha$ 1 in Cha I, using FORS1 and ISAAC at the VLT. The VRIJHK colors of Cha H$\\alpha$ 5/cc 1 are consistent with either an L-type companion or a K-type background giant. Our spectra show that the companion candidate actually is a background star. ",
        "introduction": "Twelve M6- to M8-type objects called Cha H$\\alpha$ 1 to 12 were found in deep infrared, H$\\alpha$, and X-ray surveys (Comer\\'on et al. 2000). To search for faint visual companions around these bona-fide and candidate brown dwarfs, we have taken deep images with HST WFPC2 (R, I, H$\\alpha$), VLT FORS1 (VRI), and NTT SofI (JHK$_{\\rm s}$) and detected one particulary promising companion candidate: Cha H$\\alpha$ 5/cc 1, a companion candidate 1.5 arc sec off Cha H$\\alpha$ 1, is 3.8 to 4.7 mag fainter than the primary and its colors are consistent with an early- to mid-L spectral type. Assuming the same distance, absorption, and age as for the primary, the faint object would have a mass of 3 to 15 Jupiters according to Burrows et al. (1997) and Chabrier \\& Baraffe (2000) models. The probability for this companion candidate to be an unrelated fore- or background object is $\\le 0.7\\%$, its VRIJHK colors are marginally consistent with a strongly reddened background K giant. These results are published in Neuh\\\"auser et al. (2002).  Even with the best currently achievable astrometric precision (few mas), we would have to wait several years for checking whether this visual pair is a common proper motion pair. Alternativelly, and faster, one can check by spectroscopy whether the faint companion candidate is either an L-type companion or a reddened K-type background giant. ",
        "conclusions": ""
    },
    "0207/astro-ph0207580_arXiv.txt": {
        "abstract": "In this letter  we report  on the detection   of a new feature  in the complex structure  of  the  horizontal branch   (HB) of the   Galactic globular cluster  NGC~6752.  In the  $U$ {\\it vs.}  $(U-V)$ plane, the HB shows a discontinuity  (``jump'') at $U-V\\simeq-1.0$ (corresponding $T_{\\rm   e}\\sim23,000$K).   This  ``second  $U$-jump'' adds    to the $u$-jump identified by Grundahl et al.\\  (1999) in a dozen of clusters at $T_{\\rm e}\\sim11,500$K.  We show that  this new discontinuity might be due to the combination of post zero age  HB evolution and diffusion effects.  We  identify 11 AGB-manqu\\`e   stars. The comparison between post-HB  star counts and   evolutionary   lifetimes, as predicted   by canonical  stellar models,  shows  good  agreement, at  variance  with similar estimates for NGC 6752 available in the literature. ",
        "introduction": "Although   the global properties of  horizontal   branch (HB) stars in Galactic  globular clusters (GCs) are  rather  well known, our current understanding of these objects is still challenged by several puzzling features which came out from a number of observations.  Among these we mention: (a) a non-monotonic correlation between the HB morphology and metal   abundance, i.e.\\ besides   metallicity (first parameter) there must  be a ``second''    parameter  (Sandage \\&   Wildey 1967),  or  a combination  of various parameters (Fusi   Pecci et al.\\ 1993),  which determine  the observed   HB  morphologies;   (b)  the    HB  is   not homogeneously populated, and, in  particular,  all the HBs  with  blue tails (BT) show the  presence  of  gaps, i.e.\\  regions  significantly underpopulated   by stars (Sosin et al.\\   1997, Ferraro et al.\\ 1998; Piotto et al.\\ 1999); (c) a fraction of the HBs with BTs are populated by the so called extreme (or extended)  HB (EHB) objects, i.e.\\ hot HB stars  reaching temperatures of $30,000$K  or more (D'Cruz et al 1996, Brown  et  al.\\  2001, B01), both   in metal-poor  and  in  metal-rich clusters (Rich et al.\\ 1997).  These objects are the GC counterpart of the field blue subdwarf population (Newell 1973);  (d) all the HB with BTs present a   jump,  i.e.\\ a  discontinuity  around   $T_{\\rm e}\\sim 11,500$K in the  Str\\\"omgren $u$, $u-y$  (Grundhal et al.\\ 1999,  G99) and  Johnson  $U$,  $U-V$ (Bedin et   al.\\  2000, B00) color magnitude diagrams (CMD).  Moreover, in  some of the  BTs there is  evidence of: (i) a discontinuity in the relative abundance  of heavy elements (Behr et al.\\ 1999), (ii) a discontinuity in  the surface gravities (Moehler et  al.\\ 2000),  and,  finally (iii)  a discontinuity   in the stellar rotation  velocities (Behr et al.\\  2000,  Recio-Blanco et al.\\ 2002). The abundance anomalies, the  $u$-jump,  and the discontinuity  in the (log $g$, log $T_{\\rm e}$)   plane have been interpreted as  different manifestations of  the same physical  phenomenon, i.e.\\ the appearance of  radiative metal  levitation (G99), though   Moehler et al.\\ (2000) point  out  that  this  mechanism  can only    partly account for  the low-gravity problem.  Recio-Blanco et al.\\ (2002) have also shown that the discontinuity in the rotation velocity might be related to the G99 jump.   Considerable attention  (Landsman et al.\\   1996, L96; Sweigart  1997; D'Cruz et  al.\\ 2000;  B01) has been   devoted also to  the EHB stars, because  they  still represent  a  challenge  to canonical   HB models (D'Cruz et al.\\ 1996, B01), and also  because they are suspected to be the  principal sources of  the  ultraviolet emission in the  so-called $UV$-upturn galaxies (Greggio \\& Renzini 1990).  To shed more light on the origin of EHB stars, our group has  undertaken a long-term project to   obtain  multiwavelength data  of  both  the   inner core  and the outskirts of a number of EHB clusters.  Our main  goal is twofold: (i) to constrain the main properties of EHB stars; and (ii) to investigate the role of the  dynamical  evolution on the  origin of  these objects (B00).  In this letter, we present preliminary results on one of the prototype EHB clusters,    namely    NGC~6752.  This    object  is  a     nearby [($m-M$)$_{\\circ}=13.05$,   Renzini    et al.\\   1996],  low reddening [$E_{B-V}=0.04$,  Penny   \\&  Dickens  1986], intermediate metallicity ([Fe/H]$=-1.64$) cluster, with   a  complex HB extending   to  $T_{\\rm e}\\sim32,000$K.  We show   that the HB   of  NGC~6752 presents   a new interesting feature, located at $T_{\\rm e}\\sim23,000$K, resembling the G99 $u$-jump.   We    discuss whether current  theoretical   framework accounts  for this ``second jump''.  We  also identify  post HB stars, and compare their number with current evolutionary timescales. ",
        "conclusions": ""
    },
    "0207/astro-ph0207063_arXiv.txt": {
        "abstract": "Analysis of a 30,000 s X-ray observation of the Abell 3266 galaxy cluster with the ACIS on board the {\\it Chandra} Observatory has produced  several new insights into the cluster merger. The intracluster medium has a non-monotonically decreasing radial abundance profile. We argue that the most plausible origin for the abundance enhancement is unmixed, high abundance subcluster gas from the merger. The enrichment consists of two stages: off-center deposition of a higher abundance material during a subcluster merger followed by a strong, localized intracluster wind that acts to drive out the light elements, producing the observed abundance enhancement. The wind is needed to account for both an increase in the heavy element abundance and the lack of an enhancement in the gas density. Dynamical evidence for the wind includes: (1) a large scale, low surface brightness feature perpendicular to the merger axis that appears to be an asymmetric pattern of gas flow to the northwest, away from the center of the main cluster, (2) compressed gas in the opposite direction (toward the cluster center), and (3), the hottest regions visible in the temperature map coincide with the proposed merger geometry and the resultant gas flow.  The Chandra data for the central region of the main cluster shows a slightly cooler, filamentary region that is centered on the central cD galaxy and is aligned with the merger axis directly linking the dynamical  state of the cD to the merger.  Overall, the high spectral/spatial resolution Chandra observations  support our earlier hypothesis (Henriksen, Donnelly, \\& Davis 1999) that we are viewing a  minor merger in the plane of the sky. ",
        "introduction": "Galaxy clusters show a wide range of complex phenomenon in the X-ray regime including radial abundance gradients, radial temperature gradients, cooling flows, two-phase gas in the central region (Makishima et al. 2001), non-thermal emission, and two dimension features such as shock fronts.  Observations with {\\it ROSAT}, {\\it ASCA}, and {\\it Chandra} have shown that cluster evolution proceeds through mergers and that complex physical processes account for the presence or absence of many of these observed phenomenon. With the data sets now becoming available, we are in a position to address the physics of cluster evolution. {\\it ASCA} observations of cluster mergers, such as Abell 754 (Henriksen \\& Markevitch 1996) and Coma (Watanabe et al. 1999), offer a coarse description of the pattern of shocked gas that generally matches the merger scenarios played out in N-body hydrodynamical simulations. However, with increased resolution, {\\it Newton-XMM} and {\\it Chandra} have begun to provide results that go beyond the current simulations of cluster evolution and may perhaps require additional physical processes. For example, the distribution of metals in the intracluster medium (ICM) is an important diagnostic for galaxy evolution within clusters and for the cluster gas dynamics. Two-dimensional abundance maps are within the capabilities of both {\\it Newton-XMM} and {\\it Chandra}, yet the distribution of metals in the intracluster medium has not been simulated.  Abell 3266 has become a well studied cluster merger first shown to have evidence for a merger (in addition to optical substructure) based on detection of shocked gas through detailed modeling of {\\it ASCA} observations (Henriksen, Donnelly, \\& Davis 1999). Mergers are not always apparent in the optical data and clusters may even appear to have a ``relaxed\" morphology, either spatially or dynamically. For example, though A3266 has spatial substructure, it looks dynamically ``quiet\" in the optical in that it has a velocity distribution that is statistically consistent with a single Gaussian. On the other hand, the {\\it ASCA} temperature map and the {\\it Chandra} temperature map unambiguously show high temperature regions that may be associated with shocked gas. With a detailed statistical analysis of the galaxy positions and redshifts obtained by Quintana et al. (1996), in an earlier paper we found that the cluster is composed of a secondary component comprised of approximately 30 galaxies ($\\sim$10\\% of the total number of galaxies with measured redshifts). The velocity dispersion is significantly higher, $\\sim$1300 km s$^{-1}$ in the center compared to the global value of 1000 km s$^{-1}$. This was interpreted as evidence of a merger in progress. Further evidence of a merger comes from the co-alignment of the primary and secondary galaxy condensations with the elongation in the X-ray contours first noted by Mohr, Fabricant, \\& Geller, (1993). The elongation extends from the arc minute scale in the {\\it ROSAT} PSPC image to the location of the secondary galaxy condensation, 15 $\\arcmin$ to the NE, and is visible in the {\\it ASCA} GIS image. This, together with the Gaussian velocity distribution, suggests a merger in the plane of the sky along this axis of elongation. The temperature discontinuity in the {\\it ASCA} temperature map is perpendicular to the axis connecting the subclusters suggesting that the subcluster to the NE has already traversed the core and partially shockheated the ICM. Application of the pre- and post-shock gas temperatures derived from the {\\it ASCA} temperature map to the shock equations give a post-shock gas velocity of 1400 km sec$^{-1}$. Simulations of off center cluster mergers occurring into the plane of the sky (Takizawa 2000), show a similar temperature morphology to the {\\it ASCA} map of A3266 but require a larger secondary component relative to the primary (1:4) than results from our galaxy component separation ($\\sim$1:10). Thus, the determination of merger geometry is somewhat ambiguous using only the 2-dimensional temperature map. However, a secondary component this large falling into the plane of the sky should appear as a spike in the velocity distribution and significantly distort the Gaussianity of the velocity dispersion. This is not seen in our statistical analysis that shows the inner region, outer region, and entire cluster are all consistent with a single Gaussian profile with a similar heliocentric velocity peak (Henriksen, Donnelly, \\& Davis 1999). Additionally, there is a narrow angle tail (NAT) and a wide angle tail (WAT) radio source with velocities well within the cluster velocity dispersion and located, in 2-dimensions, within the shocked gas. They are oriented with their tail aligned with the direction of the gas flow. Therefore, a major merger into the plane of the sky seems less likely than a minor merger in the plane of the sky. This point is discussed in more detail in light of the new results from {\\it Chandra} in Section 4.2  {\\it BeppoSax} observations showed that the cluster, overall, exhibits a radially decreasing temperature profile (De Grandi \\& Molendi, 1999). {\\it Chandra} provides a more detailed description of the central region and gives a temperature map that is not subject to the PSF deconvolution techniques used with {\\it ASCA} data (Section 3.1). The {\\it BeppoSAX} data did not show any departures from constant metallicity out to 10$\\arcmin$. Using Chandra, we are able to resolve the abundance within the central resolution element of the BeppoSax observations, $\\sim$2 arcmin (Section 4.1), and uncover an enhancement. These results support the merger hypothesis in Abell 3266 and give a more detailed picture of the dynamics of the ICM during a merger. Finally, the cD galaxy in Abell 3266 has multiple nuclei. Multiple nuclei are found in 28\\% of the first rank galaxies in rich clusters (Hoessel 1980). The {\\it Chandra} observations provide unique evidence that the formation of the multiple nucleus cD in Abell 3266 may be linked to the merger (Section 4.3). All parameters are quoted with 90\\% confidence errors. A Hubble constant of 60 km s$^{-1}$ Mpc$^{-1}$ is used throughout. ",
        "conclusions": "The {\\it Chandra} observations of Abell 3266 provide a detailed picture of the cluster merger. Evidence was given to support the hypothesis that the merger is in the plane of the sky and is a relatively minor one. These observations show that the merger is undoubtedly linked to the dynamical state of the dumbbell galaxy morphology. An off center abundance enhancement was found that we suggest was formed by the merger. The formation process consists of deposition of higher metallicity material during the merger followed by a cluster wind that reduces the gas density while preferentially retaining the metals. Future observation with ASTRO-E2 XRS will be especially interesting for Abell 3266 since they will provide gas velocities that will directly test our conclusions about the merger geometry and formation of the abundance profile."
    },
    "0207/gr-qc0207097_arXiv.txt": {
        "abstract": "A self-consistent field method is developed, which can be used to construct models of differentially rotating stars to first post-Newtonian order. The rotation law is specified by the specific angular momentum distribution $j(m_{\\varpi})$, where $m_{\\varpi}$ is the baryonic mass fraction inside the surface of constant specific angular momentum. The method is then used to compute models of the nascent neutron stars resulting from the accretion induced collapse of white dwarfs. The result shows that the ratios of kinetic energy to gravitational binding energy, $\\beta$, of the relativistic models are slightly smaller than the corresponding values of the Newtonian models. ",
        "introduction": "In a recent paper~\\cite{liu02} (hereafter Paper~II), we demonstrate that the accretion induced collapse (AIC) of a rapidly rotating white dwarf can result in a rapidly rotating neutron star that is dynamically unstable to the bar-mode instability, which is the instability resulting from non-axisymmetric perturbations with angular dependence $e^{\\pm 2i\\varphi}$. Here $\\varphi$ is the azimuthal angle. This instability could emit a substantial amount of gravitational radiation that could be detectable by gravitational wave interferometers, such as LIGO, VIRGO, GEO and TAMA.  However, for this instability to occur, the neutron star must have a $\\beta = T/|W|$ greater than a critical value $\\beta_d\\approx 0.25$ (Paper~II). Here $T$ is the rotational kinetic energy and $|W|$ is the gravitational binding energy. Only the AIC of white dwarfs that are composed of oxygen, neon and magnesium (O-Ne-Mg white dwarfs) with $\\Omega > 0.93 \\Omega_m$ can produce neutron stars with such a high value of $\\beta$. Here $\\Omega$ is the angular velocity of the white dwarf and $\\Omega_m$ is the angular velocity at which mass shedding occurs on the equatorial surface. This type of source will not be promising for LIGO~II because its event rate is not expected to be very high.  Neutron stars are compact objects. General relativistic effects have a significant influence on both the structure and dynamical stability of the stars. Recently, Shibata, Baumgarte, Saijo and Shapiro studied the dynamical stability of differentially rotating polytropes in full general relativity~\\cite{shibata00} and in the post-Newtonian approximation~\\cite{saijo00}. They performed numerical simulations on the differentially rotating polytropes with some specified rotation law. They found that as the star becomes more compact, the critical value $\\beta_d$ slightly decreases from the Newtonian value $0.26$ to $0.24$ for their chosen rotation law. It is not clear, however, whether their result implies that relativistic effects would destabilize the stars we are studying, for the equilibrium structure of the star will also be changed by relativistic effects. The value of $\\beta$ of a relativistic star will not be the same as that of a Newtonian star with the same baryon mass and total angular momentum.  The objective of this paper is twofold. First, we develop a new numerical technique to construct the equilibrium structure of a rotating star with a specified specific angular momentum distribution to first post-Newtonian (1PN) order [i.e., including terms of order $c^{-2}$ higher than the Newtonian terms in the equations of motion]. Then we use this new technique to construct models of neutron stars corresponding to the collapse of the white dwarfs we studied in Paper~II and Ref.~\\cite{liu01} (hereafter Paper~I) and compare them with the Newtonian models.  Equilibrium models of neutron stars in full general relativity have been built by many authors~\\cite{wilson72,bonazzola74,butterworth75,butterworth76,butterworth7679,friedman86,komatsu89}. The neutron stars studied in the literature are either rigidly rotating or rotating with an \\textit{ad hoc} rotation law. New-born neutron stars resulting from core collapse of massive stars or accretion induced collapse of massive white dwarfs are differentially rotating~\\cite{monchmeyer88,janka89,fryer01,liu01,liu02}. It seems plausible that the rotation laws of these neutron stars could be approximated by the specific angular momentum distribution $j(m_{\\varpi})$ of the pre-collapse stars (see Paper~I and Sec.~\\ref{sec:fullGR}). Here $m_{\\varpi}$ is the baryonic mass fraction inside the surface of constant specific angular momentum. Equilibrium models of Newtonian stars with a specified $j(m_{\\varpi})$ have been constructed by many authors~\\cite{ostriker68,bodenheimer73,pickett96,new01}. However, none of these studies, to our knowledge, has been generalized to include the relativistic effects.  If a rotating axisymmetric star is described by a barotropic equation of state, i.e., the total energy density $\\epsilon$ is a function of pressure only, then there is a constraint on the rotation law (see Section~\\ref{sec:fullGR}). This rotational constraint is usually written in the form $u^0 u_{\\varphi}=F(\\Omega)$~\\cite{bardeen70,butterworth76}, where $F$ is an arbitrary function. Here $\\Omega$ is the angular velocity of the fluid with respect to an inertial observer at infinity; $u^0$ is the time component of the four-velocity and $u_{\\varphi}=u_{\\mu} \\varphi^{\\mu}$, where $\\varphi^{\\alpha}$ is the axial Killing vector field of the spacetime. In the Newtonian limit, this constraint reduces to the well-known result that $\\Omega$ is constant in the direction parallel to the rotation axis. The major obstacle in the construction of differentially rotating relativistic stars is that it is not clear what function $F$ should be used to produce the desired specific angular momentum distribution $j(m_{\\varpi})$. In the next section, we will reformulate the rotational constraint in a way that can be used to impose the rotation law $j(m_{\\varpi})$, at least in the 1PN calculations.  The structure of this paper is as follows. In Section~\\ref{sec:form}, we give a brief review on the full relativistic treatment of rotating relativistic stars and then reformulate the rotational constraint imposed by the barotropic equation of state. Next, we use the standard 1PN metric and show that the rotational constraint can be solved analytically. We then derive the equations of motion determining the structure of a star to 1PN order. In Section~\\ref{sec:num}, we generalize the self-consistent field method of Smith and Centrella~\\cite{smith92} so that it can be used to compute the structure of a star to 1PN order. In Section~\\ref{sec:results}, we apply the numerical method to construct neutron star models resulting from the collapse of the O-Ne-Mg white dwarfs we studied in Paper~II and compare them with the corresponding Newtonian models. Finally, we summarize our conclusions in Section~\\ref{sec:con}. ",
        "conclusions": ""
    },
    "0207/astro-ph0207533_arXiv.txt": {
        "abstract": "I review results from, and future prospects for, microlensing searches for extrasolar planets.  Analyses of well-sampled microlensing light curves by several collaborations have demonstrated that current searches are sensitive to $\\ga 1\\mjup$ planets with few AU separations from M dwarfs in the Galactic bulge.  To date, however, no unambiguous planetary detections have been made.  Detailed analysis has shown that this null result implies that $<33\\%$ of typical stars (i.e.\\ M dwarfs) in the Galactic bulge have Jupiter-mass companions with separations between $1.5$ and $4~\\au$, and $<45\\%$ have $3\\mjup$ companions between $1$ and $7~\\au$. The recent dramatic increase in the number of alerts per year will allow ongoing microlensing searches to probe companion fractions of a few percent within a few years. ",
        "introduction": "Ultra-precise radial velocity (RV) surveys have revealed over 100 planetary companions to nearby FGKM main-sequence stars.\\footnote{See http://cfa-www.harvard.edu/planets/catalog.html for a catalog and discovery references}  The minimum masses $\\mp$ and semi-major axes $a$ of these companions are shown in Figure 3.  The number of known planetary companions is now reaching the point where robust statistical inferences about trends within the sample can be made, and thus the study of extrasolar planets is in transition from the discovery phase to the characterization phase.  The `classical' methods of detecting extrasolar planets (RV, astrometry, transits, direct detection) are generally complementary to each other both in the range of semi-major axes $a$ they probe (see Fig.\\ 3), and in the parameters they measure.  They also generally suffer from the same set of drawbacks.  First, because they rely on light from either the parent star or the planet itself, they are generally limited to nearby systems.  Second, they are not currently sensitive to very low-mass planets.  For example, the systematic floor of RV surveys is thought to be $\\sim 1~{\\rm m~s^{-1}}$. This implies that Earth, Uranus, and Neptune analogs are probably inaccessible to RV surveys. Finally, they generally require that the system be monitored for at least one full period of the companion.  Thus, although RV surveys have been ongoing for over a decade, they are only now becoming sensitive to Jupiter-analogs (Marcy et~al.\\ 2003).  Microlensing is an alternative method of detecting planetary companions that overcomes many of the difficulties inherent in the classical methods.  Mao \\& Paczy\\'nski (1991) first proposed that microlensing could be used to detect planets; their ideas were subsequently expanded on by Gould \\& Loeb (1992).  Soon after these first two seminal theoretical papers, several microlensing planet searches were initiated (Pratt et~al.\\ 1995, Albrow et~al.\\ 2000), and microlensing searches have now been ongoing, in some form, for nearly a decade.  Here I review the basic theoretical concepts behind planetary microlensing, and describe how microlensing planet searches work in practice.  I then give an overview of what the analyses of actual datasets have taught us about the practicality of the method itself, and about planetary companions in general. Finally, I briefly speculate on future prospects for microlensing planet searches. ",
        "conclusions": ""
    },
    "0207/astro-ph0207643_arXiv.txt": {
        "abstract": "We have analyzed nearly eight years of MACHO optical photometry of the supersoft X-ray binary RX~J0513.9$-$6951 and derived a revised orbital period and ephemeris.  Previously published velocities are reinterpreted using the new ephemeris.  We show that the spectroscopic characteristics of the system depend strongly on whether the system is in a high or low optical state.  We also discuss the properties of the source's high/low optical states and its long-term light curve.  Evidence for a 83.3-day periodicity in the photometry is presented. ",
        "introduction": "The supersoft X-ray source RX~J0513.9$-$6951 (hereafter called X0513$-$69) was discovered with the $ROSAT$ satellite (Schaeidt, Hasinger, \\& Trumper 1993) and identified with a $\\sim$16.7 mag emission-line star in the Large Magellanic Cloud (Pakull et al. 1993; Cowley et al. 1993).  This star appears to be the same as the variable HV 5682 found by Leavitt (1908). Its systemic velocity places it in the LMC, hence implying an absolute magnitude of M$_V\\sim-2.0$, the most luminous of all the supersoft X-ray binaries (e.g. Cowley et al. 1998).  Long-term monitoring of X0513$-$69 shows that it is normally in a high optical state, but that it fades by about a magnitude every 100-200 days, remaining in a low state for about a month each time (e.g. Alcock et al. 1996).  The optical spectrum shows extremely strong He II 4686\\AA\\ emission with a narrow peak and extremely broad wings which are also weakly seen in the Balmer and Pickering lines.  Weaker emission lines of O VI, N V, C IV, and C III are also present.  Highly red- and blue-shifted ($\\pm\\sim$4000 km s$^{-1}$) components of the strongest emission lines appear to come from bi-polar jets (Crampton et al. 1996, Southwell et al. 1996).  The bright absolute magnitude and the presence of jet emission lines suggest that the accretion disk is very luminous and the system is undergoing rapid mass transfer.  In a spectroscopic study of X0513$-$69, Crampton et al.\\ (1996) derived an orbital period of $\\sim$0.76 days from small radial velocity variations (K$\\sim$11 km s$^{-1}$) of the emission lines.  A similar photometric period was found by Motch \\& Pakull (1996) who showed the orbital light curve has a single minimum and a full range of $\\sim$0.06 mag.  An improved photometric period was derived from three years of MACHO data by Alcock et al.\\ (1996) who found P=0.76278 days with a full amplitude of $\\sim$0.04 mag during the high optical state.  The very small photometric and velocity amplitudes suggest the system is seen at a low orbital inclination.  If the optical emission lines are assumed to show the orbital motion of the compact star, the velocities imply the secondary is an evolved low-mass ($<0.3M_{\\odot}$) star.  However, inconsistencies in the published ephemerides given by different authors prompted us to use the entire MACHO data set to derive an improved orbital period so that the existing optical data and new far-ultraviolet observations (Hutchings et al. 2002) could be intercompared.  This new period analysis is based on nearly 8 years of MACHO data.  Properties of the orbital light curve, long-term photometric variations, and a rediscussion of published spectroscopy are given in this paper. ",
        "conclusions": ""
    },
    "0207/astro-ph0207369_arXiv.txt": {
        "abstract": "We present an analytic method for rapidly forecasting the accuracy of gravitational potential reconstruction possible from measurement of radial peculiar velocities of every galaxy cluster with $M> M_{\\rm th}$ in solid angle $\\theta^2$ and over redshift range $z_{\\rm min}<z <z_{\\rm max}$.  These radial velocities can be determined from measurement of the kinetic and thermal Sunyaev--Zeldovich effects. For a shallow survey with $0.2 < z < 0.4$, coincident with the SDSS photometric survey, one mode of the gravitational potential (on length scales $\\ga 60\\mpc$) can be reconstructed for every $\\sim 8$ cluster velocity determinations. Deeper surveys require measurement of more clusters per $S/N > 1$ mode. Accuracy is limited by the ``undersampling noise'' due to our non--observation of the large fraction of mass that is not in galaxy clusters. Determining the gravitational potential will allow for detailed study of the relationship between galaxies and their surrounding large--scale density fields over a wide range of redshifts, and test the gravitational instability paradigm on very large scales. Observation of weak lensing by large--scale structure provides complementary information since lensing is sensitive to the tangential modes that do not affect the velocity. ",
        "introduction": "Although there are a variety of techniques for measuring the statistical properties of cosmological density fluctuations, we know of only three types of observations from which the density field itself may be reconstructed: weak lensing \\citep{mellier99,bartelmann01}, peculiar velocities \\citep{strauss95} and galaxy rotation vectors \\citep{lee00}. A map of the density field, combined with galaxy surveys, would be a highly valuable aid to understanding the formation of galaxies and clusters of galaxies. Numerical simulations could be performed with realizations of initial conditions constrained by our knowledge of the density field, allowing object--by--object comparison between theory and observation, rather than solely statistical comparison.  A map could also provide a guide to observers who may, for example, wish to search for the luminous tracers of the filamentary density structures.  Reconstruction of the density field from galaxy peculiar velocities was pioneered by \\citet{dekel90}. The radial component of peculiar velocities of galaxies can be determined by inferring the distance and then subtracting off the Hubble flow contribution to the redshift. Galaxy peculiar velocity determinations, because they rely on distance determinations, are only useful at $z \\la 0.1$.%  Peculiar velocities of galaxy {\\em clusters} determined from observations of the Sunyaev--Zeldovich (SZ) effects \\citep{sunyaev80} do not rely on a distance determination. The peculiar velocity signal arises from the cluster's radial motion with respect to the cosmic microwave background (CMB). High--resolution ($\\la 1'$), multi--frequency observations of galaxy clusters can be used to determine the radial velocities of galaxy clusters with an accuracy that is independent of the distance to the galaxy cluster\\footnote{This independence only fails for clusters at $z \\la 0.2$ which suffer significant confusion with primary CMB anisotropy due to their large angular sizes.}. Errors as low as $100\\kms$ may be achievable \\citep{holder02}. In addition to potential reconstruction, velocity measurements will be useful for cosmological parameter estimation \\citep{peel02}.  Here we present a method for rapidly forecasting the accuracy of density field reconstruction from peculiar velocity measurements and apply it to surveys with various redshift ranges. We also show how weak lensing observations can provide complementary information \\citep{mellier99,bartelmann01}.  Our analytic method for forecasting results assumes a uniform and continuous velocity field map.  Even if this map were derived from noiseless peculiar velocity measurements, there would still be an effective noise contribution from the fact that most of the mass in the Universe is not in galaxy clusters.  Fortunately, as we quantify below, the contribution to the peculiar velocity variance from each logarithmic interval in wavenumber $k$ drops as $k^{-1}$ for relevant scales, so this noise from under--sampling is not overwhelmingly large.  We consider three different survey types labeled ``SDSS'', ``DEEP/VIRMOS''  and ``SZ'' with redshift ranges $0.2<z<0.4$, $0.7<z<1.4$ and $0.2<z<2$ respectively.  Galaxy cluster redshifts are required to convert SZ measurements into gravitational potential maps. Thus, the redshift ranges of two survey types are subsets of those of optical surveys SDSS (Sloan Digital Sky Survey), DEEP II (Deep Extragalactic Evolutionary Probe; \\citet{davis02})  and VIRMOS (Visible-Infrared Multi-Object Spectrographs; \\citet{lefevre02}). The peculiar velocity surveys we imagine could be realized by targeted multi--frequency SZ follow--up on galaxy clusters identified in those optical surveys. The ambitious, deep ``SZ\" survey is modelled after proposed SZ surveys which can locate galaxy clusters over a large $z$ range. Again, the velocities can be determined from multi--frequency follow up (if necessary).  The redshift determinations for this last survey type are left as a currently unsolved challenge.  Radial peculiar velocities are directly sensitive to radial gradients in the gravitational potential.  Gravitational lensing is caused by tangential gradients in the gravitational potential; thus, in principle, lensing provides complementary information for potential reconstruction. However, we show below that lensing is actually unlikely to add much to the density field reconstruction. The reason is simple: lensing observations are two--dimensional, whereas velocity observations are three--dimensional.  The conclusion may be different if one considers lensing of many source populations with differing redshift distributions \\citep{hu02b}. ",
        "conclusions": "We have shown that galaxy cluster peculiar velocities can be used to make high signal--to--noise maps of the gravitational potential, and therefore matter density, on very large scales. A limiting source of uncertainty in these maps is what we have called the under--sampling noise due to the fact that only a small fraction of mass is in clusters.  The under--sampling noise can be reduced by lowering the mass threshold but the lower the mass of the cluster, the harder the measurement of its peculiar velocity. Much more work is needed to understand the demand on experimental resources and to optimize observing strategies. We have made several arguments that favor shallower surveys over deeper ones.  The $\\Phi$ and density maps will have numerous applications. Cross--correlations with various tracers of the density field (red galaxies, blue galaxies, infrared galaxies, lensing of galaxies, lensing of the CMB, integrated Sachs--Wolfe effect on CMB, ...) will be of great interest.  The greatest drawback of the maps is their inability to probe scales below the typical separation between clusters."
    },
    "0207/astro-ph0207196_arXiv.txt": {
        "abstract": "We report the discovery of X-ray broad absorption lines (BALs) from the BALQSO \\apm\\ originating from material moving at relativistic velocities with respect to the central source. The large flux magnification by a factor of $\\sim$ 100 provided by the gravitational lens effect combined with the large redshift ($z = 3.91$) of the quasar have facilitated the acquisition of the first high signal-to-noise X-ray spectrum of a quasar containing X-ray BALs. Our analysis of the X-ray spectrum of \\apm\\ places the rest-frame energies of the two observed absorption lines at 8.1 and 9.8~keV. The detection of each of these lines is significant at the $ > $ 99.9\\% confidence level based on the $F$-test. Assuming that the absorption lines are from Fe~{\\sc xxv} K${\\alpha}$, the implied bulk velocities of the X-ray BALs are $\\sim$ $0.2c$ and $\\sim$ $0.4c$, respectively. The observed high bulk velocities of the X-ray BALs combined with % the relatively short recombination time-scales of the X-ray absorbing gas imply that the  absorbers responsible for the X-ray BALs  are located at radii of \\simlt 2 $\\times$ 10$^{17}$~cm, within the expected location of the UV absorber. With this implied geometry the X-ray gas could provide the necessary shielding to prevent the UV absorber from being completely ionized by the central X-ray source, consistent with hydrodynamical simulations of line-driven disk winds. Estimated mass-outflow rates for the gas creating the X-ray BALs are typically less than a solar mass per year. Our spectral analysis also indicates that the continuum X-ray emission of \\apm\\ is consistent with that of a typical radio-quiet quasar with a spectral slope of $\\Gamma$~=~1.72$_{-0.05}^{+0.06}$. ",
        "introduction": "It is commonly accepted that most quasars contain energetic outflows of ionized gas emerging from their accretion disks at speeds ranging from $\\approx$ 5,000--30,000~km s$^{-1}$ (e.g., Turnshek et al. 1988; Weymann et al. 1991). These outflows imprint broad absorption features bluewards of the resonant UV emission lines of C~{\\sc iv}, Si~{\\sc iv}, N~{\\sc v}, and O~{\\sc vi}. Broad absorption features are expected to be observed only for lines of sight that traverse the outflow. The outflow is thought to be driven by radiation pressure on spectral lines from UV photons of the central source (e.g., Arav et al. 1995; Murray et al. 1995; Proga, Stone, \\& Kallman 2000; Srianand et al. 2002). An estimate of the mass-outflow rate may result from the study of the properties of the outflowing winds of quasars. This quantity may be used to evaluate the contribution of outflowing winds in distributing accretion-disk material into the vicinity of the quasar central engine and into the host galaxy over a typical life time of a quasar. Constraining this rate is also important for understanding the connection between black hole and bulge growth in the host galaxy (e.g., Fabian 1999).  An estimate of the mass-outflow rate requires knowledge of the velocities and locations of the various ions that contribute to the wind. Broad absorption features in the UV band often show multiple detached components with different velocities, column densities, and ionization states. One needs to include the contributions of all components to obtain an accurate value of the total mass-outflow rate. Present estimates of mass-outflow rates are based mostly on the contributions from ions absorbing in the rest-frame UV band. The present X-ray data for BALQSOs are sparse, and only poor to moderate signal-to-noise ratio (S/N) spectra are available (e.g., Chartas et al. 2001; Gallagher et al. 2001; Green et al. 2001; Oshima et al. 2001; Brinkmann et al. 2002; Gallagher et al. 2002). The few moderate S/N X-ray spectra of BALQSOs available show that their X-ray faintness is due to absorption with typical hydrogen column densities ranging from $\\sim$ 10$^{22}$--10$^{24}$~cm$^{-2}$. The ionization, kinematic and spatial properties of the X-ray absorbing material are not well constrained. The physical relationship between the UV and X-ray absorbers is unclear. Is the X-ray absorber part of the outflow? Another unresolved issue is how moderately ionized species can survive the extreme UV and soft X-rays produced by the central source. To account for this, theoretical studies have postulated the presence of an optically thick layer of shielding gas between the central source and the outflow that prevents the outflow from becoming completely ionized (e.g., Murray et al. 1995). Recent simulations by Proga et al. (2000) indicate that the outflow is self-shielding, i.e., the shielding gas is an integral component of the outflow. The observed X-ray absorbing gas in BALQSOs has been suggested as a candidate for the shielding gas; however, there has been little direct observational evidence for this association. Recently, narrow absorption lines (NALs) in the X-ray band have been detected in \\chandra\\ and \\xmm\\ observations of bright Seyfert 1 galaxies (e.g., Kaspi et al. 2002). The NALs are blueshifted relative to the systemic velocity suggesting that the NAL material is part of an outflow with a mean outflow velocity of a few hundred km~s$^{-1}$.  To improve our understanding of the X-ray absorption in BALQSOs, we performed a long \\chandra\\ observation of the bright gravitationally lensed BAL quasar \\apm. The flux magnification of \\apm, estimated to be $\\sim$ 100 (Egami et al. 2000; Mu{\\~ n}oz et al. 2001) in the X-ray band, and its high redshift of $z = 3.91$ allowed us to study the kinematic and ionization properties of a BAL quasar in the X-ray band. Specifically, the high redshift of \\apm\\ places the strong Fe~K features at energies where the telescope effective area is much larger. Here we present the results of these observations. ",
        "conclusions": ""
    },
    "0207/astro-ph0207657_arXiv.txt": {
        "abstract": "Relative ages for the globular cluster (GC) subpopulations in the Virgo giant elliptical galaxy M87 (NGC~4486) have been determined from Str\\\"omgren photometry obtained with WFPC2 on board \\textit{HST}.  Using a variety of population synthesis models, and assuming the GC mass at the turnover of the luminosity function is the same for both subpopulations, differential ages have been determined from the observed magnitudes at the turnover of the globular cluster luminosity function and from the mean color of each subpopulation.  We measure an age difference between the two subpopulations of $0.2 \\pm 1.5$ (systematic) $\\pm 2$ (random) Gyr, in the sense that the blue GCs are formally older. Thus, to within our measurement errors, the two subpopulations are found to be coeval. Combined with previous spectroscopic age determinations for M87 GCs, our results favor a picture in which the GCs associated with this galaxy are formed at high redshift, and within a period of a few Gyr. ",
        "introduction": "Globular clusters (GCs) are among the oldest stellar systems in the Universe.  Their large numbers around early type galaxies, their ease of detection and their simplicity as single-age, mono-metallic stellar populations, make them ideal tracers of the evolutionary history of their host galaxies. A key result that has emerged from observations of GC systems is that a large fraction of early-type galaxies show a bimodal distribution in broadband color (e.g. Gebhardt \\& Kissler-Patig 1999). As broadband colors of old stellar populations are much more sensitive to metallicity than age, this bimodality indicates the presence of two GC subpopulations, a metal-rich and a metal-poor, but due to the age-metallicity degeneracy (e.g. Worthey 1994) no firm conclusion can be drawn on the ages of the two components based on broadband colors alone. Several models have attempted to explain the bimodality in terms of a particular galaxy formation process: major mergers of late-type galaxies (Ashman \\& Zepf 1992), two bursts of \\textit{in-situ} star formation (Forbes, Brodie \\& Grillmair 1997) and hierarchical formation (C\\^ot\\'e, Marzke \\& West 1998). The first two models form the metal-rich GCs out of a second burst of star formation while the third assumes that the GC system is assembled via dissipationless mergers.  A key observational constraint on those models is the relative ages of the GC subpopulations. Multiple metallicity populations do not necessarily require separate bursts of star formation but may also reflect differences in the environment where the GC subpopulations formed. Elucidating the formation histories of the GC subpopulations on the basis of age can help discriminate between the proposed models.   Most previous determinations of the ages of the GC subpopulations in undisturbed early-type galaxies seem to suggest coeval subpopulations (Kissler-Patig et~al 1998; Cohen, Blakeslee \\& Ryzhov 1998; Puzia et~al. 1999; Larsen et~al. 2002), albeit usually with large uncertainties.  For M87, there have been conflicting claims. Cohen et~al.  (1998)  obtained Keck spectroscopy of $150$ GCs in M87 and found no sign of a variation in age with metallicity, excluding the possibility that one population is half as old as the other at the $99\\%$ confidence level. On the other hand, Kundu et~al. (1999) estimate using $V$ and $I$ photometry from \\textit{HST}, that the metal-rich clusters are $3-6$ Gyr younger than their metal-poor counterparts.  We have obtained Str\\\"omgren photometry of the GC system in the inner part of M87. By combining a metallicity-sensitive color index with the turnover $u$-band luminosity of the GC luminosity function (GCLF) of the two subpopulations, we have measured their relative age by comparing our observations with population synthesis models. We find the two GC subpopulations to be coeval within our measurement errors, in agreement with the spectroscopic estimates. ",
        "conclusions": "\\label{sec:resdis}  The turnover magnitudes for the metal-rich and metal-poor subpopulations of M87 and their respective mean colors were compared to the predictions of population synthesis models. We used the latest versions of the models of Bruzual \\& Charlot (1993), Maraston (1998)  and Worthey (1994) to compute grids for every available metallicity. Because a particular value for the mass at the turnover is not assumed, the \\textit{absolute} ages for the subpopulations are unknown.  Instead, the model grid was shifted in such a way that one of the subpopulations lies on the $14$ Gyr isochrone, which is consistent with the age of the oldest MW GCs (VandenBerg 2000), and then the relative age was obtained by linearly interpolating from the model grid.  Choosing another age for the oldest isochrone does not change our results significantly. A graphical example of the procedure can be seen in Figure~\\ref{fig:grid}. For each set of models we tried two different IMFs: a Salpeter (1955) and a Miller \\& Scalo (1979)  for the Bruzual \\& Charlot (1993) and Worthey (1994) models and a Salpeter (1955) and a Gould, Bahcall \\& Flynn (1997)  for the Maraston (1998) models.   \\begin{figure} \\plotone{f2.eps} \\caption{Determination of the relative ages of the GC populations using Maraston's (1998) models with a Gould et~al. (1997) IMF. The short dashed lines are isometallicity tracks of $Z=0.0001,0.001,0.01,0.02,0.04$ and the long dashed lines are isochrones of $3$, $6$, $9$ and $15$ Gyr. The points are at the estimated turnover luminosities and mean $s_2$ color for both populations; the error bars are $1\\sigma$ uncertainties. The dotted line is the isochrone of $14$ Gyr where the blue population has been assumed to lie. The age for the red population, and thus the relative age, is then interpolated from the grid. (Some isochrones have been omitted for clarity).} \\label{fig:grid} \\end{figure}  \\begin{deluxetable}{c c c c c c c c c} \\tablecaption{Observed and Derived Properties of M87 Globular Clusters\\label{tab:res}} \\tablehead{ \\colhead{Index} & \\multicolumn{2}{c}{Metal-poor}& \\multicolumn{2}{c}{Metal-rich} & \\colhead{IMF\\tablenotemark{a}} & \\multicolumn{3}{c}{Relative Age\\tablenotemark{b} (Gyr)}\\\\ \\cline{2-5}  \\cline{7-9}\\\\ \\colhead{} & \\colhead{$m_{col}$}& \\colhead{$m_{TO}$}& \\colhead{$m_{col}$}& \\colhead{$m_{TO}$}& \\colhead{} & \\colhead{M98} & \\colhead{W94} & \\colhead{BC93} \\\\ } \\tablecolumns{9} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\startdata $s_1$ & $0.67\\pm0.015$ & $24.04\\pm0.09$ & $1.27\\pm0.017$ & $25.09\\pm0.15$ & S55 & $1.3\\pm3.1$ & $-1.9\\pm3.0$ & $1.8\\pm3.2$\\\\ &  &  &  &  & MS79 & \\nodata &$-1.3\\pm2.3$ & $-0.5\\pm2.2$ \\\\ &  &  &  &  & GBF97 & $0.0\\pm3.2$ & \\nodata & \\nodata \\\\ $s_2$ & $1.07\\pm0.024$ & $23.99\\pm0.09$ & $1.95\\pm0.027$ & $25.00\\pm0.14$ & S55 & $1.5\\pm2.7$ & $-1.4\\pm2.7$ & $2.5\\pm2.9$\\\\ &  &  &  &  & MS79 & \\nodata & $-0.4\\pm2.2$ & $-0.1\\pm2.1$ \\\\ &  &  &  &  & GBF97 & $0.4\\pm2.7$ & \\nodata & \\nodata \\\\ \\enddata \\tablenotetext{a}{S55 = Salpeter 1955; MS79 = Miller \\& Scalo 1979; GBF97 = Gould, Bahcall \\& Flynn 1997.} \\tablenotetext{b}{M98 = Maraston 1998; W94 = Worthey 1994; MC93 = Bruzual \\& Charlot 1993.} \\end{deluxetable}  The results are summarized in Table~\\ref{tab:res}, where the reported errors come from measurement uncertainties. We take the dispersion of the age differences for each color index, $\\sigma \\sim 1.5$ Gyr, as an indication of the systematic errors arising from the models.  As the typical measurement error is 2-3 Gyr, and we have measurements using two color indices, the mean of our results gives a formal age difference of ${\\Delta}t \\equiv {\\rm age}_{blue} - {\\rm age}_{red} = 0.2 \\pm 1.5$ (systematic) $\\pm 2$ (random) Gyr. Note that the quoted systematic uncertainty does not include possible differences in turnover masses; a 10\\% difference in the mass at the turnover would change the inferred age difference by $\\sim$ 1.5 Gyr.  \\textit{Within the errors, all models are consistent with the two populations being coeval}.  The average metallicities obtained from the models for the subpopulations are [Fe/H]$_{blue}=-1.58$ and [Fe/H]$_{red}=-0.30$, in reasonable agreement with the values [Fe/H]$_{blue}=-1.41$ and [Fe/H]$_{red}=-0.23$ obtained by Kundu et~al. (1999). Cohen et~al.  (1998) found no variation of age with metallicity from their Keck spectra of M87 GCs, for which they find a mean age of $13.2$ Gyr. Dividing the results presented in their Table 4 into two populations and averaging the results gives a formal age difference of $\\Delta t \\sim 1$ Gyr with a typical uncertainty in their age estimates of $\\sim 2$ Gyr. Our results are in good agreement with theirs, and both are consistent with the subpopulations being coeval within the uncertainties. This is gratifying as their sample of GCs is completely independent of ours and they are using a different set of diagnostics from the models. Indeed, as they determine absolute ages for each GC without making additional assumptions about the GC masses, the agreement may be taken as indirect evidence in support of our basic assumption that the GCLF turnovers for the different subpopulations correspond to approximately the same mass.   Although our observations were carried out in the same field as Kundu et~al. (1999), we do not find support for their claim that the metal-rich subpopulation is $3-6$ Gyr younger than the metal-poor.  As no uncertainty is given for their estimate we are not able to assess the significance of the disagreement, but none of the combinations of color indices, models and IMFs that we have explored gives a result which falls within their estimated range.  Our conclusions depend rather strongly on the assumption of equal mass at the turnover for both subpopulations. Even though the agreement with Cohen et~al.  (1998) and evidence from our Galaxy (McLaughlin \\& Pudritz 1996) suggest that this is a reasonable assumption, a direct test of this hypothesis for some early-type galaxies is not only desirable but within reach of efficient high-resolution spectrographs on $8$m-class telescopes. It is also worth bearing in mind that the turnover for the metal-rich population in the $u$ band falls near the level of $\\sim 50\\%$ completeness. This makes the use of parametric modeling (including the completeness function) a necessity. Deeper observations will be required to observe the metal-rich turnover directly.  As evidence accumulates for the GC subpopulations of undisturbed giant ellipticals being coeval (Kissler-Patig et~al. 1998; Cohen et~al.  1998; Puzia et~al. 1999; Larsen et~al. 2002) models that rely on well separated epochs of star-formation, either via major mergers or \\textit{in-situ} bursts of star formation, appear less favored as a \\textit{generic} formation mechanism, unless one is willing to push the separation of the two epochs to the point of being almost coeval.  This is not to say that major mergers of late-type galaxies are not a viable mechanism -- populations of young GCs have almost certainly been identified in ongoing mergers (e.g. Whitmore \\& Schweizer 1995) -- but their role as the primary mechanism for the formation of giant ellipticals seems unlikely. Likewise, if two \\textit{in-situ} star-forming bursts are to be the formation mechanism for bright ellipticals, there has to be enough time for the metals to diffuse and enrich the gas from which the secondary GCs form. Lacking a specific mechanism which is responsible for regulating this star formation, it is hard to constrain this idea. Given the precision in the relative ages, the most we can say is that enrichment, outflows/inflows and star formation process must have occured within a $\\sim$ few Gyr if this scenario is correct. Until a mechanism for producing precisely two bursts is identified, this scenario amounts to the \\textit{assumption} that two metallicity populations imply two periods of star formation.  Our results combined with those of Cohen et~al.  (1998) seem to show that whatever the assembly history of these galaxies, the GCs associated with M87 formed at at an early time, and within a short period. The presence of two populations differing in metallicity by a factor of $\\sim 20$, and yet with roughly the same age, suggests that the origin of the two populations may be related to differences in the local \\textit{environments} where the GC formed. According to the standard paradigm of structure formation, galaxies form via accretion and merging of small objects. Simulations predict a merger rate for massive cluster galaxies such as M87 that is highly peaked at redshifts $z \\gae 4$ (Gottl\\\"ober, Klypin \\& Kravtsov 2001). This scenario provides the necessary difference in the local environments as the GCs can form in protogalactic fragments of varying mass. Simulations of the GC metallicity distributions based on this picture have been succesful in reproducing the observations (C\\^ot\\'e, West \\& Marzke 2002).  As the bulk of star formation is expected to happen early, and within $\\sim$ a few Gyr, this picture for the formation of bright ellipticals appears consistent with the existing observations of their GC systems."
    },
    "0207/astro-ph0207077_arXiv.txt": {
        "abstract": "We present a study of the shape of the Large Magellanic Cloud disk.  We use the brightnesses of core helium-burning red clump stars identified in $V-I,I$ color-magnitude diagrams of 50 randomly selected LMC fields, observed with the CTIO 0.9-m telescope, to measure relative distances to the fields.  Random photometric errors and errors in the calibration are controlled to $\\lesssim$1\\%. Following correction for reddening measured through the color of the red clump, we solve for the inclination and position angle of the line of nodes of the tilted plane of the LMC, finding $i=35\\fdg8\\pm2\\fdg4$ and $\\theta=145\\arcdeg\\pm4\\arcdeg$.  Our solution requires that we exclude 15 fields in the southwest of the LMC which have red clump magnitudes $\\sim$0.1 magnitudes brighter than the fitted plane.  On the basis of these fields, we argue that the LMC disk is warped and twisted, containing features that extend up to 2.5 kpc out of the plane.  We argue that alternative ways of producing red clump stars brighter than expected, such as variations in age and metallicity of the stars, are unlikely to explain our observations. ",
        "introduction": "Among the many virtues of the Large Magellanic Cloud (LMC), one of the most important is that in spite of its proximity, for most purposes its stars may be considered to lie at a single distance.  As a result, the LMC is frequently used in studies of stellar evolution, is of primary importance in the calibration of the distance scale, and is the target of several projects aimed at detecting microlensing events in the Galactic halo (e.g. Alcock et al.\\ 2000, Lasserre et al.\\ 2000, Udalski et al.\\ 1999, Drake et al.\\ 2002).  The neighboring Small Magellanic Cloud (SMC) has never reached the same level of popularity, for the interesting reason that it is stretched by several kpc along the line of sight (e.g.\\ Caldwell \\& Coulson 1986, Hatzidimitriou \\& Hawkins 1989), the probable result of tidal interactions with the LMC and the Milky Way.  Probing deeper, we find that the stars of the LMC are not, after all, quite at the same distance, because the LMC disk is tilted with respect to the line of sight.  The LMC's tilt has most recently been measured by van der Marel \\& Cioni (2001), who found its value to be $\\sim$35\\arcdeg ~using asymptotic giant branch and red giant branch tip (TRGB) stars observed in the 2MASS survey (Skrutskie 1998).  The effect of the tilt is to produce a modulation of the apparent luminosity of these tracers with position angle in the LMC's disk, with an amplitude of $\\sim$0.2 magnitude peak-to-trough. Moreover, upon deprojecting the 2MASS data, van der Marel (2001) found that the LMC disk is elliptical and has a nonuniform surface density distribution, indicative of tidal forces acting on the LMC.  Clearly, the LMC is not slipping through its interactions with the Milky Way and the SMC completely unscathed.  We are thus led to ask, what further evidence for the LMC-SMC-Milky Way interaction may we find hidden in the LMC's geometry?  The question is important for chronicling the interaction history through numerical models (e.g.\\ Gardiner \\& Noguchi 1996), and may have implications for interpreting the microlensing results (e.g.\\ Zaritsky \\& Lin 1997, Zhao \\& Evans 2000).  In this paper, we report the results of our use of the apparent luminosity of ``red clump'' stars to measure the shape of the LMC disk.  The red clump, which is composed of the younger, more metal-rich analogues to the core helium-burning horizontal branch stars found in globular clusters, has been widely discussed in the literature for use as a distance indicator (e.g.\\ Girardi \\& Salaris 2001, GS01; Cole 1998).  While the red clump remains controversial as an absolute distance indicator, by using it to measure {\\em relative} distances to LMC fields we avoid the disagreement over its zero point. Red clump stars are more numerous than TRGB stars by a factor of $\\sim$100 and have tightly defined colors and magnitudes, making them easily identifiable in color-magnitude diagrams (CMDs).  In the LMC, they have a typical surface density of 3.5$\\times10^4$ stars per square degree, allowing us to search for structure with high spatial frequency.  Because they have well-defined mean colors in addition to magnitudes, we can measure reddenings with the same tracers used to measure distances, avoiding the population-dependent reddening effects discussed by Zaritsky (1999).  The errors associated with using the red clump as a distance indicator are mainly systematic ones.  We discuss their possible effect on our results in section 4. ",
        "conclusions": "We have studied the shape of the LMC by measuring the variation in luminosity of red clump stars across 50 13\\arcmin$\\times$13\\arcmin ~fields distributed over a 6\\arcdeg$\\times$6\\arcdeg area.  Our observations were taken with the CTIO 0.9-m telescope CCD camera using $V$ and $I$ filters.  We clearly detect the tilt of the LMC disk through a decrease in brightness of the red clump stars along the LMC's NE-SW axis. Our measured tilt of $35\\fdg8\\pm2\\fdg4$ is in excellent agreement with the recent measurement of van der Marel \\& Cioni (2001), {\\em as long as} we exclude 15 fields in the southwest which appear to lie out of the plane of the disk.  Including {\\em all} of the fields, we measure a lower inclination of $24\\arcdeg\\pm2\\arcdeg$.  This value is in good agreement with that measured by Caldwell \\& Coulson (1986); however, the fitted plane is of lower statistical significance.  We believe that we have detected a warp in the LMC disk which causes a region $\\sim$2 kpc wide in the southwest to lie out of the plane by $\\sim$2.5 kpc.  If it exists, this out-of-plane feature is likely a non-virialized structure.  How and when was it produced?  The location of the warp points a finger at the SMC.  According to the model by Gardiner \\& Noguchi (1996), the LMC and SMC endured a close encounter $\\sim$200 Myr ago which drew out the material that new occupies the inter-Cloud region.  It is possible that this interaction also altered the shape of the LMC disk.  However, the position of the warp also aligns it with the major axis of the LMC's elliptical disk (van der Marel 2001), the shape of which van der Marel attributes to the tidal influence of the Milky Way.  Perhaps our suggested warp is an additional response to the Milky Way's tidal force?  Weinberg (1998, 2000) studied the mutual influence of the LMC and Milky Way through numerical simulations. Although Weinberg found that the long-term effect on the LMC's structure is simply heating of the disk, transient structures might also be produced.  Further exploration of the structure of the LMC could determine whether the warp is reflected in northeastern fields more distant than those observed here.  The determination of the star formation and chemical enrichment history in the LMC's southwest would establish with certainty the contribution of age and metallicity effects to the observed red clump luminosities. In this study we avoided the LMC's inner regions, including the Bar, as these are heavily crowded.  Measurement of the LMC's shape and thickness in these regions is of strong interest for understanding the contribution of LMC self-lensing to the observed microlensing events (Zhao \\& Evans 2000).  While our results bear no consequence for the microlensing studies, they do add to the growing body of evidence suggesting that the LMC is not a simple flattened disk.  A study of the red clump in the LMC Bar, combined with the continued monitoring of the LMC for microlensing events by the SuperMACHO project (Drake et al.\\ 2002), could provide very interesting results."
    },
    "0207/astro-ph0207594_arXiv.txt": {
        "abstract": "We present evidence for approximately 400-d variations in the radial velocity of \\srco\\ (V934 Her), the suggested optical counterpart of \\srcx.  The variations are correlated with the previously reported $\\approx400$~d variations in the X-ray flux of \\srcx, which supports the association of these two objects, as well as the identification of this system as the second known X-ray binary in which a neutron star accretes from the wind of a red giant. The \\srco\\ radial velocity variations can be fit with an eccentric orbit with period $404\\pm3$~d, amplitude $K=0.75\\pm0.12\\ \\kms$ and eccentricity $e=0.26\\pm0.15$. There are also indications of variations on longer time scales $\\ga2000$~d.  We have re-examined all available ASM data following an unusually large X-ray outburst in 1997--98, and confirm that the 1-d averaged 2--10~keV X-ray flux from \\srcx\\ is modulated with a period of $400 \\pm 20$~d. The mean profile of the persistent X-ray variations was approximately sinusoidal, with an amplitude of $0.108 \\pm 0.012$~ASM~$\\cts$ (corresponding to 31\\% rms). The epoch of X-ray maximum was approximately 40~d after the time of periastron according to the eccentric orbital fit. If the 400-d oscillations from \\srco/\\srcx\\ are due to orbital motion, then the system parameters are probably close to those of the only other neutron-star symbiotic-like binary, GX~1+4. We discuss the similarities and differences between these two systems. ",
        "introduction": "The X-ray source \\srcx\\ (also known as 2A 1704+241; $l=45\\fdg15$, $b=32\\fdg99$) was first discovered in {\\it Ariel V}\\/ scans for high-latitude X-ray sources \\cite[]{2acat}, and by the $Uhuru$ (SAS A) X-ray observatory \\cite[]{4ucat}.  It has a typical flux of (1--$10)\\times10^{-11}\\ {\\rm erg\\,cm^{-2}\\,s^{-1}}$ \\cite[2--10~keV, equivalent to 0.5--5~mCrab as measured by a range of X-ray missions;][]{masetti01}, and experiences occasional outbursts. During a 100-d X-ray high state in 1997--98, \\srcx\\, was observed by the {\\it Rossi X-ray Timing Explorer}\\/ \\cite[\\xte;][]{xte96}.   This high state was the larger of two known outbursts of this system, and the X-ray flux reached a peak of 40~mCrab. The X-ray spectrum of \\srcx\\, is similar to that of other accreting neutron stars, and generally requires a blackbody with a temperature of 0.9--1.3~keV plus a hard component at higher energies for an acceptable fit. The X-ray spectral hardness and the lack of UV continuum imply that the compact object is almost certainly not a white dwarf \\cite[]{garcia83}. A black hole companion is also unlikely since the thermal X-ray component in black hole candidates is typically an order of magnitude lower in temperature than that measured in \\srcx. However, no rapid  quasi-periodic or coherent X-ray variations have been confirmed \\cite[]{masetti01}, which is rather unusual for an accreting neutron star.  Based on {\\it Einstein}\\/ position measurements, \\cite{garcia83} proposed the association of \\srcx\\ with the 8th magnitude M2 III star \\srco.  In a recent re-analysis of {\\it ROSAT}\\/ HRI data, however, \\cite{mg01} found only a 10\\% chance that the two sources are associated, but did not propose an alternative candidate.  Previous optical spectroscopy ruled out radial velocity variations with amplitudes greater than $5\\ \\kms$ on time scales of 40 min to $\\sim 1$~yr \\cite[]{garcia83}, strongly suggesting that either the orbital period is significantly longer than one year, or that the system is viewed very close to face-on (inclination angle $i\\approx0$).  \\srco\\ may be only the second symbiotic-like binary known to contain a neutron star, after the well-studied M-giant/X-ray pulsar system V2116~Oph/GX~1+4.  However, these systems are quite different in X-rays (Table \\ref{comptable}).  Furthermore, the optical spectrum of V2116~Oph exhibits extremely strong, variable emission lines, probably powered by UV photons originating from an accretion disk \\cite[]{chak97:opt}.  In contrast, \\srco\\ has much weaker line emission.  In fact, optically the spectrum is very close to that of an isolated M giant \\cite[e.g.][]{garcia83}.  Here we describe optical and X-ray observations which go some way towards resolving the uncertainties surrounding this system.  We compare the properties of \\srco/\\srcx\\ with those of V2116~Oph/GX~1+4, and discuss what may be inferred about the former in light of our results.  \\centerline{\\epsfxsize=8.5cm\\epsfbox{f1.eps}} \\figcaption{Radial velocity measurements of \\srco\\ and the corresponding $404\\pm3$~d orbital solution. The top panel shows the radial velocity measurements versus time and the estimated $1\\sigma$ uncertainties.  The bottom panel shows the measurements folded on the orbital period.  In both panels the eccentric orbital model is overplotted as a solid line (see Table \\ref{orbtbl}). \\label{orbit} } \\bigskip ",
        "conclusions": "We have found evidence for significant variation in the radial velocity measurements from \\srco, which we fit with an eccentric orbit with period $404\\pm3$~d.  We have also found evidence for a significant variation in the X-ray flux from \\srcx.  Our best estimate of the X-ray period is $400 \\pm 20$ d.  Hence, the periods measured from the radial velocity and ASM X-ray flux variations are consistent. It is compelling to suggest that these two periodicities arise from a common source. Naturally, it is also possible that the agreement between the two periodicities is merely coincidence, and that \\srco\\ and \\srcx\\ are unrelated.  Estimating the probability of such a coincidence depends upon the origin of the periodicities in each star, as well as the intrinsic distributions of these phenomena. In practice, it is not possible to make such an estimate with much confidence. However, it is clear that the likelihood of an unrelated red giant and X-ray source being at nearly the same measured position at high galactic latitude, and also both having a periodicity \\centerline{\\epsfxsize=8.5cm\\epsfbox{f2.eps}} \\figcaption{\\xte/ASM measurements of \\srcx. The top panel shows the 1~d averaged ASM measurements between 1996 January 5 and \\asmend\\ (MJD 50087--\\asmendmjd).  The horizontal lines (-- --) bracket the estimated times of phase maximum for the best period estimate (400~d; see \\S\\ref{sec:xte}).  The open triangles indicate the predicted times of periastron passage for the eccentric orbital fit ($P_{\\rm orb}=404$~d; Table \\ref{orbtbl}).  The bottom panel shows the periodogram calculated from a subset of the ASM data after the end of the outburst (1998 March 28), where the measurement error was $<0.5\\ \\cts$. The dashed vertical line shows the candidate 400~d period.  The inset shows the folded ASM profile using the time of periastron from the orbital fit as the reference phase. \\label{asm} } \\bigskip \\noindent that would be unusual for either type of source separately, is extremely small.  The detection of a 400-d modulation from both stars therefore supports the original suggestion that HD 154791 is the optical counterpart of \\srcx.   \\subsection{Cause of the Correlation} \\label{sec:causeofcor}  There are several plausible mechanisms for a correspondence between the radial velocity and the persistent X-ray flux. We consider two possibilities for the origin of the 404-d period in the \\srco\\, radial velocity curve: orbital motion, and pulsational variation of the M-giant.  Precise optical photometric monitoring could distinguish between these two possibilities, since the maximum brightness occurring at any phase besides that of maximum recessional velocity would favour orbital motion.  If the radial velocity variations are due to orbital motion, the 400-d X-ray flux modulation may arise from changes in the accretion rate onto the neutron star as it follows an elliptical orbit around the companion and moves through different nebular densities. Such modulation is common in high-mass wind-accreting X-ray binaries \\cite[e.g.][]{swr86}.  Moreover, the times of periastron passage in our best-fit orbital model are very well correlated with the times of increased X-ray flux.  For a 1.4~$M_\\odot$ neutron star and red giant mass $\\la 2.5\\ M_\\odot$, the 404~d period and low amplitude of the radial velocity variations imply that the inclination is $i\\la2^\\circ$.  Although the {\\it a priori} probability of observing such a system is $\\la7\\times10^{-4}$, such a low inclination could explain the lack of X-ray pulsations or QPOs from the source (see \\S\\ref{gxcomp}). If the 400-d variation is in fact due to orbital motion, the longer-term variations in the radial velocities may be related to irregular variations in the red giant \\cite[]{pwh01}.  Alternatively, red-giant pulsations with $P=400$~d could modulate the stellar wind, and hence the accretion rate onto the neutron star.  The IRAS fluxes of HD~154791 ($f_{12}=2.77$, $f_{25}=0.736$, $f_{60}<0.4$, $f_{100}<1.0$, where $f_{12}$ is the flux density $f_\\nu$ at 12 microns in Jy) indicate that is it not a Mira variable \\cite[]{kenyon88}.  However, non-Mira M-giants also pulsate.  Although typical fundamental pulsation periods lie in the range 20--200~d, variations with a period up to an order of magnitude larger are also seen \\cite[]{pwh01}.  If the 400-d variations are due to pulsation, one of the slower trends observed in the radial velocity variations could be related to orbital motion.  In this case, the possible range of orbital inclinations becomes much more likely.  An orbital period of 2370~d gives an inclination of 4--$6^\\circ$ while for a period of 3700~d the likely range is 5--$7^\\circ$.  Also, the low accretion rate onto the neutron star would be more understandable if the orbital period is longer than 400 d. However, a longer orbital period alone cannot account for the discrepancy in luminosity between 4U~1700+24 and GX~1+4.  Finally, the shape of the 400-d radial velocity variation argues against a red-giant pulsation.  Pulsation light curves are typically non-sinusoidal (in the sense that more time is spent with brightness below the median value than above it), and similar behavior would be expected for the radial velocity variations.  \\subsection{Comparison with GX~1+4} \\label{gxcomp}  If the 400~d periodicity in the \\srco\\, radial velocities is due to orbital motion, then the orbital periods of \\srco/\\srcx\\ and GX~1+4 may be quite similar (see Table \\ref{comptable}). This similarity makes the contrast between their other observed properties puzzling. The X-ray luminosity of GX~1+4 appears to be two to three orders of magnitude greater than that of \\srcx.  In GX~1+4, there is clear evidence for an accretion disk \\cite[]{chak97:opt}, whereas in \\srcx\\, there is no measurable UV continuum \\cite[]{garcia83}.  GX~1+4 has a rich optical emission-line spectrum, and \\srco\\, shows very little in the way of emission lines.  Finally, GX~1+4 is a 2-minute X-ray pulsar, whereas no coherent or quasi-period oscillations have been convincingly detected from \\srcx.  One explanation for the differing X-ray luminosities is a larger mass-loss rate from the late-type giant in GX~1+4.  \\cite{chak97:opt} find that the M6 giant in GX~1+4 is probably near the tip of the first-ascent red-giant branch.  \\srco\\, has a spectral type M2, and so would be expected to be losing mass at a lower rate. The relative lack of optical/UV emission lines suggests that either the ionized nebula in \\srco/\\srcx\\ is substantially smaller and/or less dense than in GX~1+4, or there are not enough UV photons to power nebular line emission in \\srcx/\\srco.  Either of these conditions can plausibly arise if the mass loss from the companion is much smaller. Moreover, if the binary separation is smaller in GX~1+4 than in \\srcx/\\srco\\ a larger fraction of the red giant wind could be accreted.  The higher X-ray luminosity in GX~1+4 could also illuminate the red giant and increase the mass loss from the red giant further, in a sense producing a feedback effect.  The observed variability of the optical spectrum of V2116~Oph is evidence for the significant role played by X-ray heating. On the other hand, in \\srcx/\\srco\\ the X-ray luminosity is perhaps low enough that very little illumination of the red giant or enhanced mass loss is expected. Even during the 1997 X-ray outburst, \\cite{tomasella97} note that the optical spectrum was unchanged.    Finally, the lack of pulsations or QPOs in \\srcx\\ remains puzzling. \\cite{garcia83} and \\cite{mg01} proposed a 900-s QPO, but this QPO was not confirmed in high-quality \\xte\\/ data \\cite[]{masetti01}.  High-precision optical photometry has also revealed a lack of pulsations (down to a limit of 0.5\\% fractional variation), or rapid variations of any kind \\cite[]{sbh01}.  This behavior is again in contrast to GX~1+4, where \\cite{jab97} find an oscillation in the optical emission with an amplitude of from 0.4 to 4.5\\%, plus stochastic variations.  Given that the optical light from GX~1+4 has a significant contribution from an accretion disk, the lack of rapid optical variations from \\srco\\, provides evidence that any contribution from a disk in \\srcx\\, is negligible in the optical regime as well as the UV regime.  The lack of X-ray pulsations is expected if the 400-d radial velocity variations are due to orbital motion.  As discussed in \\S\\ref{sec:causeofcor}, the low amplitude of the radial velocity variations requires $i\\sim 0$, in which case any magnetic hot spots on the neutron star will be continuously in view (assuming the NS spin axis is aligned with angular momentum axis of the binary).  For higher and {\\it a priori} more probable values of $i$, the non-detection of pulsations may be explained if the accretion flow is not appreciably disrupted by the magnetic field of the neutron star above the surface \\cite[e.g.][]{gl79a,gl79b}, and spreads the accreting material more or less evenly over the neutron star surface. Given the extremely low luminosity (and hence accretion rate) of this source, standard calculations of ram and $B$-field pressure balance would require that the surface magnetic field strength of the neutron star be $\\la10^6$~G, far below the canonical value for low-mass X-ray binaries or even old neutron stars in recycled ms pulsars. A more likely scenario consistent with larger $i$ is that the neutron star spin is very slow.   In conclusion, all scenarios for \\srcx/\\srco\\, have some difficulties. We have presented some options, but more observational constraints are needed to select beween them.  Additional observations to determine whether the 400-d radial velocity variations are due to orbital motion or red-giant pulsation would also enable one to better interpret the differences between \\srcx/\\srco\\, and GX~1+4."
    },
    "0207/astro-ph0207288_arXiv.txt": {
        "abstract": "In the present study we analyze the behavior of the rotational velocity, $v\\sin~i$, for a large sample of 134 spectroscopic binary systems with a giant star component of luminosity class III, along the spectral region from middle F to middle K. The distribution of $v\\sin~i$ as a function of color index (\\bv) seems to follow the same behavior as their single counterparts, with a sudden decline around G0III. Blueward of this spectral type, namely for binary systems with a giant F--type component, one sees a trend for a large spread in the rotational velocities, from a few~km~s$^{-1}$  to at least 40~km~s$^{-1}$. Along the G and K spectral regions there is a considerable number of binary systems with moderate to moderately high rotation rates. This reflects the effects of synchronization between rotation and orbital motions. These rotators have orbital periods shorter than about 250 days and circular or nearly circular orbits. Except for these synchronized systems, the large majority of binary systems with a giant component of spectral type later than G0III are composed of slow rotators. ",
        "introduction": "Rotation is one of the most important observable physical parameters in stellar astrophysics. Such a parameter can provide fundamental constraints for models of stellar evolution, as well as important information on the link between surface rotation and stellar atmospheric phenomena. The behavior of the rotational velocity for single evolved stars of luminosity class III is now well established (De Medeiros and Mayor, 1989, 1991; Gray, 1989). For this luminosity class, there is a  sudden decline in the rotational velocity around the spectral type G0III, which corresponds to (\\bv)~$\\sim$~0.70. Blueward of this spectral type, namely for F type single giants, rotational velocity scatters over a wide range of values, from about 2 to 180 km s$^{-1}$, whereas redward, namely for G and K--type single giants, stars are essentially slow rotators and rotation rates greater than 5~km~s$^{-1}$ are unusual. On the basis of  the analysis of kinematic--age relation, De Medeiros and Mayor (1991) have shown that the root cause for this discontinuity in rotation seems to be a mixing in ages associated with the rapid evolution of giant stars into the Herztsprung Gap. These observational results have been discovered from a large rotational and radial-velocity survey of about 1100 giants (De Medeiros and Mayor 1999) accomplished with the CORAVEL high-resolution spectrometer (Baranne et al. 1979) at the Haute Provence Observatory, in France, and at the European Southern Observatory in Chile.  We have now examined the complete survey for the spectroscopic binary systems containing a giant component of luminosity class III. The main challenge in the study of rotation in binary systems with evolved components is to establish the extent of the synchronization between rotational and orbital motion along the giant branch. Tidal theory (e.g., Zahn 1977) predicts that, in late type binary systems, viscous dissipation of time dependent tidal effects should produce synchronization between rotation and stellar orbital motion, as well as circularization of the orbit of the system. The most simple way to test such effects consists of the determination of precise rotational velocities for a large sample of binary systems with giant components, presenting a wide variety of values of orbital parameters. On the basis of tidal predictions, Middelkoop and Zwaan (1981) have suggested that most late type giants in close binary systems, with orbital periods shorter than about 80 days, rotate in synchronization with revolution, because most of these close binary systems have circular orbits. Such a tendency for synchronization in late type binary systems  was also observed by Giuricin et al. (1984). Mermilliod and Mayor (1992) have found that, in open clusters, binaries containing a giant have circularized orbits at orbital periods shorter than about 250 days. More recently, Boffin et al. (1993) have deduced a circularization cut off period of about 70 days from an eccentricity--period diagram for a large sample of field binary systems containing late type giants.  In the present work, we study the behavior of the rotational velocity $v\\sin~i$ as a function of the color index (\\bv), for a large sample of binary systems with a giant component of luminosity class III. We also analyze the link between $v\\sin~i$ and the orbital parameters eccentricity and orbital period. ",
        "conclusions": "Precise rotational velocities, $v\\sin~i$, are given for a large sample of 134 single lined binary systems with an evolved component of luminosity class III. For these binary systems the distribution of rotation rates as a function of color index (\\bv) presents a behavior that seems to parallel the one found for their single counterparts. Namely, there is a sudden decline in rotation around the spectral type G0III, as is the case for single giants. Binary systems located blueward of the spectral type G0III, typically those systems with a color index (\\bv) lower than about 0.70, present a large spread in the values of rotational velocity, from a few~km~s$^{-1}$ to at least 40~km~s$^{-1}$. In addition, we have found that along the G and K spectral regions there is a considerable number of moderate to moderately high rotators reflecting, clearly, the effects of synchronization between rotation and orbital motions. These stars have orbital periods shorter than about 250 days as well as circularized or nearly circularized orbits. For a given spectral type, the process of synchronization increases the observed rotational velocity of the binary systems up to about 15 times the mean rotational velocity of single giants. Except for these synchronized systems, the majority of binary stars later than spectral type G0III is essentially composed of slow rotators, following the same behavior as single giants. A few binary systems present synchronization characteristics and low rotation but probably their $v\\sin~i$ is suppressed by low inclination of their rotation axis. We have also observed that enhanced rotation is not a general property among binary systems with circularized orbit. Admittedly, in the present study the number of binary systems with F--type component is very scarce and one should be cautious with the proposed rotational discontinuity around the spectral type G0III. In addition the extent of tidal effects along the F spectral region is not yet established. In this context only the determination of rotational velocity and orbital parameters for a larger sample of binary systems with F--type component could confirm such a discontinuity on more solid basis. Finally, we would like to point out that for the large majority of stars discussed in this work the duplicity is established on very solid grounds. However for 58 stars additional measurements of radial velocity are undoubtedly necessary to establish their orbital parameters. Nevertheless, it is important to underline that the discussion of the rotational velocity for binary systems with evolved component carried out in the present work is not hampered by this fact because, in particular, the large majority of stars with no available orbital parameters are slow rotators."
    },
    "0207/astro-ph0207241_arXiv.txt": {
        "abstract": "We review all the \\bsax results relative to the search for additional \\nt components in the spectra of clusters of galaxies. In particular, our MECS data analysis of A2199 does not confirm the presence of the \\nt excess reported by Kaastra \\etal (1999). A new observation of A2256 seems to indicate quite definitely that the \\nt fluxes detected in Coma and A2256 are due to a diffuse \\nt mechanism involving the intracluster medium. We report marginal evidence ($\\sim 3\\sigma$) for a \\nt excess in A754 and A119, but the presence of point sources in the field of view of the PDS makes unlikely a diffuse interpretation. ",
        "introduction": "It is well known that X-ray measurements in the energy range 1-10 keV of thermal bremsstrahlung emission from the hot, relatively dense intracluster gas, have already contributed in an essential way to our understanding of the cluster environment. However, recent researches on \\cl have unveiled new spectral components in the intracluster medium (ICM) of some clusters, namely a cluster soft excess discovered by \\euve (Lieu \\etal 1966) and a hard X-ray (HXR) excess detected by \\bsax (Fusco-Femiano \\etal 1999) and \\rxte (Rephaeli, Gruber, \\& Blanco 1999). Observations at low and high energies can give additional insights on the physical conditions of the ICM.  Nonthermal emission was predicted at the end of seventies in \\cl showing extended radio emission, radio halos or relics (see Rephaeli 1979). In particular, the same radio synchrotron electrons can interact with the CMB photons to give inverse Compton (IC) nonthermal X-ray radiation. Attempts to detect \\nt emission from a few \\cl were performed with balloon experiments (Bazzano \\etal 1984;90), with \\heao1 (Rephaeli, Gruber \\& Rothschild 1987; Rephaeli \\& Gruber 1988), with the OSSE experiment onboard the \\textit{Compton-GRO} satellite (Rephaeli, Ulmer \\& Gruber 1994) and with \\rxte \\& \\asca (Delzer \\& Henriksen 1998), but all these experiments reported essentially flux upper limits. However, we want to remind the conclusions of the paper regarding the OSSE observation of HXR radiation in the Coma cluster by Rephaeli, Ulmer \\& Gruber in 1994: \"\\textit{..It can be definitely concluded that the detection of the IC HEX (high energy X-ray) emission necessitates an overall sensitivity a few times $10^{-6}\\pho$ in the 40-80 keV band. ..To reduce source confusion, detectors optimized specifically for HEX measurements of clusters should have $\\sim 1^{\\circ}\\times 1^{\\circ}$ fields of view. A level of internal background more than a factor of 10 lower than that of OSSE is quite realistic. Obviously, another very desirable feature of any future experiment is wide energy coverage, starting near (or below) 15-20 keV, in order to independently measure the tail of the thermal emission}\". In these conclusions it is possible to find the spectral characteristics of the Phoswich Detector System (PDS) onboard \\bsax which is able to detect hard X-ray emission in the 15-200 keV energy range. The PDS uses the rocking collimator technique for background subtraction with angle of 3.5$^{\\circ}$. The strategy is to observe the X-ray source with one collimator and to monitor the background level on both sides of the source position with the other in order to have a continuous monitoring of the source and background. The dwell time is 96 sec. The background level of the PDS is the lowest obtained so far with high-energy instruments onboard satellities ($\\sim 2\\times 10^{-4} ~\\rm {counts~s}^{-1}~\\rm keV^{-1}~$ in the 15-200 KeV energy band) thanks to the equatorial orbit of \\bsax. The background is very stable again thanks to the favorable orbit, and no modelling of the time variation of the background is required (Frontera \\etal 1997). ",
        "conclusions": "\\bsax observed a clear evidence of \\nt emission in two clusters, Coma and A2256, both showing extended radio regions. In particular, the two observations of A2256 strongly support the presence of a diffuse non thermal mechanism involving the ICM. These detections and the lack of detection in other clusters seem to indicate that the essential requirement to observe additional \\nt components at the level of the PDS sensitivity is the presence of large regions of reaccelerated electrons, with Lorenz factor $10^4$, due to the balance between radiative losses and reacceleration gains in turbulence generated by merger events that must be very recent considering the short lifetime of the electrons.  \\bsax, as it is well known, has ceased its activity at the end of April 2002. The next missions able to search for \\nt components are : \\par\\noindent $\\bullet$ \\integral. In particular, with IBIS, that has a spatial resolution of 12$'$, we have the opportunity a) to localize the source of the nonthermal X-ray emission. In the case of a point source it is possible to identify it, while in the case of a diffuse source it is possible to verify whether the \\nt emission is mainly concentrated in the cluster central region or in the external region, as predicted by the model for the Coma cluster of Brunetti \\etal (2001), or it is uniformly spread over the whole radio halo present in the cluster. b) to have a better determination of the photon spectral index. \\par\\noindent $\\bullet$ \\astroe. The Hard X-ray Detector (HXD) has a \\fov of $34'\\times 34'$ similar to that of the MECS. A positive detection of the \\nt emission already measured by \\bsax in Coma and A2256 would eliminate the ambiguity between a diffuse emission involving the intracluster gas and a point source, considering that the MECS images do not show evidence for point sources. \\par\\noindent $\\bullet$ The future missions are represented by {\\it NEXT} and \\constellation.  \\vskip5pt  These missions will be operative in the next years, but the energy range and the spectral capabilities of \\xmm/EPIC give the possibility to localize \\nt components in regions of low gas temperature as shown by the simulation regarding the radio relic of A2256 performed using the \\nt flux measured by \\bsax (see Fig. 9). This region has a gas temperature of 4 keV likely associated with the ongoing merger shown by a \\chandra observation (Sun \\etal 2001). So with \\xmm we should have the possibility, by comparing the X-ray and radio structures, to constrain the profiles of the magnetic field and of relativistic electrons."
    },
    "0207/astro-ph0207131_arXiv.txt": {
        "abstract": "Red clump giants in the Galactic bulge occupy a distinct region in the colour magnitude diagram. They have a very small spread in intrinsic luminosities and their number counts have a well defined peak. We show that these characteristics can be used to constrain the differences in the streaming motions of stars on the near side and those on the far side. We propose two methods to select two samples with one preferentially on the near side and the other on the far side. In the first method, we divide red clump giants into a bright sample and a faint one; stars in the bright sample will be on average more on the near side and vice versa. The second method relies on the fact that lensed bulge stars lie preferentially on the far side due to the enhanced lensing probability by the stars on the near side and in the disk. If the radial streaming motion is $\\approx 50\\kms$, we find the difference in the average radial velocity between the bright and faint samples can reach $\\approx 33\\kms$ while the corresponding difference is about $\\approx 10\\kms$ between the lensed stars and all observed stars. The difference in the average proper motion between the bright and faint samples is about $\\approx 1.6 \\mas\\yr^{-1}$ if there is a tangential streaming motion of 100$\\kms$; the corresponding shift between the lensed stars and all observed stars is approximately $\\approx 1\\mas\\yr^{-1}$. To observe the shifts in the radial velocity and proper motion, roughly one hundred microlensing events, and/or bright/faint red clump giants, need to be observed either spectroscopically or astrometrically. The spectroscopic observations can be performed efficiently using multi-object spectrographs already available. The proper motion signature of microlensed objects can be studied using ground-based telescopes and the Hubble Space Telescope. These observations will provide strong constraints on the Galactic bar parameters. ",
        "introduction": "Many observational groups carried/are carrying out microlensing observations (e.g., MACHO: Alcock et al. 1993; OGLE: Udalski et al. 1993; DUO: Alard \\& Guibert 1997; MOA: Bond et al. 2001; EROS: Aubourg et al. 1993). Over one thousand microlensing events in the local group have been identified (e.g., Alcock et al. 2000; Wo\\'zniak et al. 2001; Derue et al. 2001; Alard \\& Guibert 1997; Bond et al. 2001). In the coming years, $\\sim$ 1000 microlensing events are expected to be discovered by the OGLE III\\footnote{ http://www.astrouw.edu.pl/\\~\\,ogle/ogle3/ews/ews.html } and other collaborations every year, many of them in real-time.  Data collected from microlensing experiments have very diverse applications (for reviews see Paczy\\'nski 1996; Gould 1996). One of the most important applications is studying the Galactic structure. There is strong evidence that the centre of Galaxy hosts a bar (e.g., de Vaucouleurs 1964; Blitz \\& Spergel 1991; Stanek et al. 1994, 1997; Kiraga \\& Paczy\\'nski 1994; H\\\"afner et al. 2000 and references therein). However, the parameters of the bar are not well-determined, including its mass, size, and the motion of stars within it.  Data from microlensing experiments provide several ways of probing the structure of the inner Galaxy. For example, the optical depth is roughly proportional to the total mass of the bar while the event time scale distribution probes the mass function and kinematics of the bar. Another important method that was first explored by Stanek et al. (1997) uses red clump giants (RCGs). These bulge stars occupy a distinct region in the colour-magnitude diagram (see, e.g., Stanek et al. 2000 and references therein). They have very small intrinsic widths in their luminosity function, about 0.2\\,mag for RCGs in the Galactic bulge (see \\S2 in Stanek et al. 1997; Paczy\\'nski \\& Stanek 1998). The observed luminosity function has a well-defined peak and an apparent width (see Fig.~5 in Stanek et al. 1997). The apparent width depends on both the intrinsic width and the spread caused by the radial depth of the bar. Stanek et al. (1997) used this dependence to constrain the Galactic bar axial ratios and orientation.  In this paper, we use the RCGs to study the kinematics of stars in the Galactic bar. Stars on the front side and far side may have different (tangential and radial) streaming motions. We explore two effects to constrain such motions. The first effect is based on the realisation that red clump stars at the bright slope of the peak of the luminosity function must be (on average) closer to us, i.e. on the near (front) side, while the stars on the faint slope must be (on average) on the far (back) side. Therefore, one can select bright and faint samples of RCGs and examine their differences in the radial velocity and proper motion. The second effect we explore is based on the well-known fact that lensed bulge stars lie preferentially on the far side of the Galactic bar due to the enhanced lensing probability by the stars on the near side. This leads to several observable effects. For example, the lensed RCGs should be fainter compared with all observed red clump stars in the field (Stanek 1995; see also Zhao 1999a, 1999b, 2000 for other effects such as extinction and reddening for microlensing events toward the Large Magellanic Cloud). We will examine systematic differences in the proper motions and radial velocities of the lensed RCGs relative to all observed RCGs in the Galactic bar; a similar bias in the radial velocity for lensing events toward the Large Magellanic Cloud has been discussed by Zhao (1999b).  As this paper is the first feasibility study, we do not attempt to construct a detailed (or even self-consistent) model of the Galactic bulge. Our aim is to provide an order-of-magnitude estimate of the kinematic biases in the hope of motivating observations to be carried out, which may in turn provide strong constraints on the the model of the inner Galaxy and incentives for their refinement. The structure of the paper is as follows. In \\S2, we outline our model, and in \\S 3 we present our main results.  We discuss observational issues to detect these effects in \\S 4. For definiteness, we shall restrict ourselves to studying the red clump stars in Baade's window ($l=1^\\circ$, $b=-3.9^\\circ$). ",
        "conclusions": "In this paper we have studied how to use red clump giants (RCGs) to constrain the kinematics of stars in the Galactic bar. The basic idea is that stars on the near and far sides may have different streaming motions, and if we can select two samples with one preferentially on the near side and the other on the far side, then they should show differences in the radial and tangential streaming motions. We have examined two ways that we can select such samples. The first method selects a bright sample and a faint one with the peak magnitude of the number counts of RCGs as the approximate dividing line. The RCGs in the faint sample is (on average) more on the far side of the bar and vice versa. The second method relies on the fact that stars on the far side are preferentially microlensed due to the enhanced lensing probability by the stars on the near side and in the disc.  Hence the lensed stars and all observed stars should behave differently in kinematics.  We illustrate the kinematical differences in the context of two simple models of the Galactic bulge. In the first (axis-symmetric) model,  the streaming motion is purely tangential, while in the second (more realistic) bar model the streaming motion is more or less radial. In reality, the streaming motions in the Galactic bar may have both a radial component and a tangential one, so our numbers should be taken as rough order-of-magnitude estimates.  We found at the difference in the radial velocity between the bright and faint samples can reach  $\\dvr \\approx 33\\kms$ for a radial streaming motion of $50\\kms$. The difference is substantially reduced by the population of red giants that occupy the same part of the colour-magnitude diagram as the RCGs, as they have a large span in luminosity, so a brighter magnitude does not signal that a star is closer to us on average. If there is an effective way of differentiating red giants and red clump stars, then the differences in the radial velocity and proper motion can be much larger. This can be, for example, done by examining colours as RCGs occupy a slightly bluer part of the colour-magnitude diagram. Spectroscopic observations may also ultimately provide features that distinguish red giants and RCGs. For example, if the contribution of red giants can be reduced by one half, then the shift in the proper motions can be as high as $2.3\\masyr$, while that for the radial velocity can reach $50\\kms$. In comparison, we find that the shift between the lensed stars and overall population of stars is more modest, about $10\\kms$ for a radial streaming motion of $50\\kms$.  The velocity dispersion along the line of sight is about $\\sigma\\approx 100\\kms$, so in order to see this shift at the $2\\sigma$ level, from Poisson statistics, we only need to obtain the radial velocity for about $(2 \\times \\sigma/\\dvr)^2 \\approx 100$ stars even for $\\dvr=20\\kms$. The required number is moderate, but is within reach of multi-object spectrographs available on many large telescopes. For example, the FORS1 instrument on VLT has 19 slits that can be efficiently used to derive the radial velocities of lensed stars and field stars simultaneously. Other instruments such as 2dF \\footnote{http://www.aao.gov.au/2df/} which can take up to 400 spectra simultaneously should also be explored. We note that for RCGs, blending is not a severe problem as they are so bright, so their radial velocities can be measured without much difficulty. In any case, we only require an accuracy of $\\sim 10-20\\kms$ per star, since the dispersion in the radial velocity for the bulge stars is of the order of $100\\kms$.  For the shift in the proper motions, we find that a value of $1.6\\mas\\yr^{-1}$ between the bright and faint samples of RCGs for a tangential streaming motion $\\vt =100\\kms$. The shift scales roughly linearly with $\\vt$. The difference is about $1\\mas\\yr^{-1}$ between the lensed and overall population of bright stars. These shifts are not much smaller than the dispersion in the proper motion, approximately $3\\mas\\yr^{-1}$. The shifts can be detected with HST or even ground-based instruments. To do this, only a modest number of RCGs need to be monitored to measure the relative shifts in proper motions. The relative proper motion may be feasible to detect even from the ground. Soszy\\'nski et al. (2002) demonstrated that proper motions as small as $4\\mas\\yr^{-1}$ can be measured with OGLE II data. This accuracy is achieved through a combination of the new difference image analysis software (Eyer \\& Wo\\'zniak 2001) and a large number of data points for any given star, which reduces the astrometric error through Poisson statistics. The camera system of OGLE III yields much better images than that of OGLE II as the latter operates in the drift scan mode and produced somewhat worse point spread functions. Therefore it seems possible (although still to be demonstrated) that proper motions of many red clump stars, including lensed ones, can be obtained using data from microlensing surveys (e.g., OGLE III), requiring no extra observing resources.  In this paper we have only explored the shifts in the radial velocity and proper motion in Baade's window. Clearly there must be a spatial dependence of these shifts. The OGLE III collaboration is currently monitoring many fields around the Galactic centre, covering from $-7^\\circ$ to $6^\\circ$ in latitude and $-12^\\circ$ to $12^\\circ$ in longitude \\footnote{For detailed field coordinates, see http://www.astrouw.edu.pl/\\~\\,ogle/ogle3/ews/gb\\_ews.jpg}. Red clump stars and microlensing events will become available over large areas, and so it will be very interesting to detect the spatial variation of the shifts in the radial velocity and proper motion. Such observations will likely provide strong constraints on the Galactic bar parameters, and stimulate further theoretical efforts to build a better and self-consistent model of the inner Galaxy. Finally, we point out that if future astrometric missions, SIM and GAIA\\footnote{http://sim.jpl.nasa.gov; http://astro.estec.esa.nl/GAIA/}, perform as planned, they will directly measure the radial depth of the galactic bar by determining parallax's for a number of stars.  This will make our suggestion obsolete, but not sooner than in a decade.  We thank Martin Smith, David Spergel, and Hongsheng Zhao for helpful discussions. S.M. acknowledges travel support by Princeton University. This project was supported by the NSF grant AST-1206213, and the NASA grant NAG5-12212 and funds for proposal \\#09518 provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555."
    },
    "0207/astro-ph0207307_arXiv.txt": {
        "abstract": "We have observed the bipolar jet from RW\\,Aur\\,A with STIS on board the HST. After continuum subtraction, morphological and kinematic properties of this outflow can be traced to within 0\\farcs1 from the source in forbidden emission lines. The jet appears well collimated, with typical FWHMs of 20 to 30 AU in the first 2$''$ and surprisingly does not show a separate low-velocity component in contrast to earlier observations. The systemic radial outflow velocity of the blueshifted lobe is typically 50\\% larger than that of the redshifted one with a velocity difference of about 65\\,\\kms. Although such asymmetries have been seen before on larger scales, our high spatial resolution observations suggest that they are intrinsic to the ``central engine'' rather than effects of the star's immediate environment. Temporal variations of the bipolar jet's outflow velocities appear to occur on timescales of a few years. They have combined to produce a 55\\% increase in the velocity asymmetry between the two lobes over the past decade. In the red lobe estimated mass flux $\\dot{M}_j$ and  momentum flux $\\dot{P}_j$ values are around one half  and one third of those for the blue lobe, respectively. The mass outflow to mass accretion rate is 0.05, the former being measured at a distance of $0.''35$ from the source. ",
        "introduction": "\\label{intro}  Although there have been many studies of the propagation of outflows from young stars (e.g.\\ see the reviews in {\\em Protostars and Planets IV}) relatively little is known, from an observational perspective, about their generation \\citep{eisletal00}. A major problem is that the source itself is often embedded at optical/near-infrared wavelengths making high resolution observations, except in the radio band, impossible. That said, there are a number of optically visible young stars associated with outflows. Although their outflows are not as striking as those from more embedded young stars, such systems nevertheless represent the best window on the ``central engine'' that we have. With these ideas in mind, we have embarked on a Hubble Space Telescope (HST) program to examine outflows close to a small number of optically visible young stars \\citep{bmresc00} including the subject, RW\\,Aur, of this {\\em Letter}.  RW\\,Aur (HBC~80, HIP 23873) is a member of the Taurus-Auriga star forming region. It is a binary system with RW\\,Aur\\,B  located at a position angle of $258^{\\circ}$ and a projected separation of $1\\farcs50$ with respect to RW\\,Aur\\,A \\citep{lei1993}. Taurus-Auriga is usually assumed to have a distance of $\\approx$~140~pc (e.g.\\,Wichmann et al.\\,1998). In contrast the {\\em Hipparcos} distance of RW\\,Aur is $70.5\\pm 34.0\\,\\mathrm{pc}$ but, as RW\\,Aur is a close binary, and variability shifts the system's photocentre, the latter distance has been questioned (Bertout, Robinchon, \\& Arenou 1999). Here, we will assume 140~pc but one should keep in mind that this system may be closer. RW\\,Aur\\,A is one of the optically brightest T~Tauri stars in the sky with V = 10.1\\,mag \\citep{hbc1988}. The presence of strong H$\\alpha$ emission (EW[H$\\alpha$] = 84 {\\AA}, Herbig \\& Bell 1988) and high veiling at near-infrared wavelengths (Folha \\& Emerson 1999) categorise it as a classical T Tauri star (CTTS) with active accretion.  As is usual with CTTSs, accretion is accompanied by outflow: \\citet{hirth1994} first discovered an asymmetric bipolar jet (HH 229) from RW\\,Aur\\,A using long-slit spectroscopy and found that the blueshifted and redshifted jets {on extended scales} differed in absolute systemic velocity by a factor of two. Another unusual property of this bipolar outflow is that the brighter lobe (in the red [SII] doublet) is redshifted. A detailed analysis of the spectroscopic data of \\citet{hirth1994} is presented in \\citet{baccetal96} who showed that the redshifted jet from RW\\,Aur\\,A also has a low ionization fraction. Further observations by \\citet{muneis98} of the area surrounding this star revealed that its flow is much more extensive than previously thought (with a total size of at least $145''$) while \\citet{doug00} examined the bipolar jet close to RW\\,Aur\\,A using adaptive optics.  Here we report on multiple observations with the Space Telescope Imaging Spectrograph (STIS) of the bipolar jet from RW\\,Aur\\,A, which allow us to spatially resolve the jet not only along the outflow direction but transversely as well. This is the first bipolar YSO jet to be studied with STIS so close to its origin. Brief observational details are given in Section 2 and our results are presented in Section 3. ",
        "conclusions": "\\label{summ} We have obtained spectra of the bipolar jet of RW\\,Aur\\,A at an unprecedented high spatial resolution of 0\\farcs1 and for the first time studied its morphological and kinematic properties within one arcsecond from its origin. Both outflow lobes can be traced as close as 0\\farcs1 from the source in the FELs. In H$\\alpha$ the STIS spectra show two strong maxima on the star with radial velocities coinciding with those of the two outflow lobes. The jet appears well collimated very close to its origin and aymmetries in the red- and blueshifted lobe velocities arise within a region smaller than 20\\,AU from the source. This scale is almost certainly a conservative estimate as our values for the jet inclination angle to the line of sight and distance to RW\\,Aur are lower and upper limits respectively.  In contrast to previous observations we do not find a separate low velocity outflow component, which indicates this feature is variable on timescales of a few years. If this component is a disk wind, then our observations imply such winds vary on similar timescales as the higher velocity outflow and may also be episodic. Variations in the high velocity components in the RW\\,Aur\\,A jet suggest that its knots may be internal working surfaces. This is also supported by our estimates of the mass and momentum fluxes, which in both lobes are observed to decrease by at least an order of magnitude within the first 2$''$ from the source. Close to the star we find an average mass flux of about 8.5 10$^{-8}$ M$_{\\odot}$ yr$^{-1}$, which is about 5\\% of the mass accretion flux. The mass flux and momentum flux values in the red lobe are typically one half and one third of the values in the blue lobe, respectively. Finally we note that velocity variations in the blue- and redshifted lobes imply an increase in the radial velocity asymmetry of about 55\\% over the last decade."
    },
    "0207/hep-ph0207175_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\vskip -0.1cm  Prior \\footnote{The aim of this article is to provide a fairly non-technical and up to date review of the motivations for mirror matter and the evidence for its existence. For a more detailed and (hopefully) entertaining exposition see the recent book\\cite{book}.} to 1957 scientists believed that mirror reflection symmetry was respected by the interactions of the fundamental particles. Why? Perhaps because the other geometrical symmetries such as rotations and translations in space and time were observed to be good symmetries, it seemed natural that mirror reflection symmetry should be a good symmetry too. Furthermore, no experiment up until 1957 had indicated that these symmetries were anything but good symmetries of nature.  In 1956 Lee and Yang\\cite{ly6} proposed that the interactions of the fundamental particles were not mirror reflection invariant. They suggested that this could explain some known puzzles and proposed some new experiments to directly test the idea. Subsequently Madam C.S.Wu and collaborators dramatically confirmed that the interactions of the known particles were not mirror symmetric, just as Lee and Yang had suspected.  Today, it is widely believed that mirror symmetry is in fact violated in nature. God -- it is believed -- is left-handed. Actually, though, things are not so clear. What the experiments in 1957 and subsequent experiments have conclusively demonstrated is that the {\\it known} elementary particles behave in a way which is not mirror symmetric. The weak nuclear interaction is the culprit, with the asymmetry being particularly striking for the weakly interacting neutrinos. For example, today we know that neutrinos only spin with one orientation. If one was coming towards you it would be spinning like a left-handed corkscrew. Nobody has ever seen a right-handed neutrino.  The basic geometric point is illustrated in the following diagram: \\vskip 0.5cm \\centerline{\\epsfig{file=fig1a.eps,width=8.1cm}} \\vskip 0.35cm \\noindent The left-hand side of this figure represents the interactions of the known elementary particles. The forces are mirror symmetric like a perfect sphere, except for the weak interaction, which is represented as a left hand. Also shown is nature's mirror - the vertical line down the middle. Clearly, the reflection is not the same as the original, signifying the fact that the interactions of the {\\it known} particles are not mirror symmetric. If there were a right hand as well as a left hand then mirror symmetry would be unbroken. \\vskip 0.45cm \\centerline{\\epsfig{file=fig1b.eps,width=7.7cm}} \\vskip 0.5cm \\noindent However, this doesn't correspond to nature since no right-handed weak interactions are seen in experiments (this is precisely what the experiments in 1957 and subsequently have proven).  There are two remaining possibilities: We can either chop the hand off -- but this is too drastic and is therefore not shown. It corresponds to having no weak interactions at all, again in disagreement with observations. This last possibility is the most subtle and consists of adding an entire new figure with the hand on the other side. Everything is doubled even the symmetric part, which is clearly mirror symmetric as indicated in the following diagram: \\vskip 0.5cm \\centerline{\\epsfig{file=fig1c.eps,width=9.4cm}} \\vskip 0.5cm \\noindent What this figure corresponds to is a complete doubling of the number of particles. For each type of particle, such as electron, proton and photon, there is a mirror twin. Where the ordinary particles favor the left hand, the mirror particles favor the right hand. If such particles exist in nature, then mirror symmetry would be exactly conserved (we denote the mirror particles with a prime). \\vskip 0.7cm \\centerline{\\epsfig{file=fig2.eps,width=3.4cm}} \\vskip 0.5cm \\noindent As will be discussed, the mirror particles can exist without violating any known experiment. Thus, the correct statement is that the experiments in 1957 and subsequently have only shown that the interactions of the {\\it known} particles are not mirror symmetric, they have not demonstrated that mirror symmetry is broken in nature.  While many people regard the possible existence of mirror particles as highly speculative, it could also be argued that the assumption of broken mirror symmetry is equally speculative. Clearly, it is not possible to figure out which path is chosen by nature on the basis of pure thought. What really needs to be done is to understand the experimental implications of the existence of mirror particles and find out whether such things could describe our Universe.  The mirror partners have the same mass as their ordinary counterparts, which is reminisant of anti-particles. However, there is a crucial difference. Unlike anti-particles, the mirror particles interact with ordinary particles predominately by gravity only. The three non-gravitational forces act on ordinary and mirror particles completely separately [and with opposite handedness: where the ordinary particles are left-handed, the mirror particles are right-handed]. For example, while ordinary photons interact with ordinary matter (which is just the microscopic picture of the electromagnetic force), they {\\it do not} interact with mirror matter\\footnote{ Actually, as will be explained in a moment, it is possible for small {\\it new} forces to exist which connect the ordinary and mirror particles together. For the purposes of this introductory paragraph, this possibility is temporarily ignored. }. Similarly, the `mirror image' of this statement must also hold, that is, the mirror photon interacts with mirror matter but does not interact with ordinary matter.  The upshot is that we cannot see mirror photons because we are made of ordinary matter. The mirror photons would simply pass right through us without interacting at all!  The mirror symmetry does require though that the mirror photons interact with mirror electrons and mirror protons in exactly the same way in which ordinary photons interact with ordinary electrons and ordinary protons. A direct consequence of this is that a mirror atom made from mirror electrons and a mirror nucleus, composed of mirror protons and mirror neutrons can exist. In fact, mirror matter made from mirror atoms would also exist with exactly the same internal properties as ordinary matter, but would be completely invisible to us! Clearly, if there was a negligible amount of mirror matter in our solar system, we might hardly be aware of its existence at all. Thus, the {\\it apparent} left-right asymmetry of the laws of nature may be due to the preponderance of ordinary matter in our solar system rather than due to a fundamental asymmetry in the laws themselves.  While this is all very interesting, the most remarkable thing of all is that there is now a range of evidence actually supporting the mirror matter theory: \\vskip 0.15cm \\noindent 1) It predicts the existence of mirror matter in the Universe. Mirror matter would be invisible, making its presence felt by its gravitational effects. Remarkably, there is a large body of evidence for such invisible `dark' matter. There is also specific evidence that mirror stars have been observed from their gravitational effects on the bending of light from background stars.  \\vskip 0.20cm \\noindent On the quantum level, small new fundamental interactions connecting ordinary and mirror matter are possible. Various theoretical constraints suggest only a few possible types of interactions: photon-mirror photon kinetic mixing and neutrino-mirror neutrino mass mixing\\cite{flv,flv2}. Such non-gravitational forces are extremely important and open up new ways in which to test the theory: \\vskip 0.15cm \\noindent 2) Orthopositronium should have a shorter effective lifetime (in a ``vacuum\" experiment) than predicted due to the effects of photon - mirror photon kinetic mixing\\cite{gl,fg}. \\vskip 0.15cm \\noindent 3) If there are small mirror matter bodies in our solar system then this would lead to a new class of cosmic impact events when such bodies strike the Earth. Characteristics of such impacts will be the lack of ordinary fragments and other anomalous features. Such impacts will leave mirror matter embedded in the ground which could potentially be extracted and purified (assuming that the small photon-mirror photon kinetic mixing force exists, which is strong enough to oppose the Earth's feeble gravity). \\vskip 0.15cm \\noindent 4) If there is some remnant mirror hydrogen gas in our solar system, then spacecraft will experience a drag force and slow down. \\vskip 0.15cm \\noindent 5) If neutrinos have mass then oscillations between ordinary and mirror neutrinos can occur. Such effects could show up in neutrino physics experiments. \\vskip 0.1cm  At the present time there is interesting experimental/observational evidence supporting all five of these predictions. We now describe this evidence in more detail. ",
        "conclusions": "It is a known fact that almost every plausible symmetry (such as rotational invariance, translational invariance etc) are found to be exact symmetries of the particle interactions. Thus, it would be very strange if the fundamental interactions were not mirror symmetric. It is a very interesting observation that mirror symmetry requires the existence of a new form of matter called `mirror matter', otherwise there is nothing to balance the left-handed nature of the weak force. Even more interesting, is the remarkable evidence that mirror matter actually exists. But, does mirror matter really exist? I'm not sure, but I would very much like to find out. Maybe the answer lies in Jordan or maybe it is blowing in the wind\\cite{bobd}...   \\vskip 0.4cm \\noindent {\\large \\bf Acknowledgements} \\vskip 0.1cm \\noindent It is a pleasure to thank my collaborators, Sergei Gninenko, Sasha Ignatiev, Henry Lew, Saibal Mitra, Zurab Silagadze, Ray Volkas and T. L. Yoon. I would also like to acknowledge very useful correspondence from Sergei Blinnikov, Zdenek Ceplecha, Carl Feynman, Luigi Foschini, John Learned, Jesus Martinez-Friaz and Andrei Ol'khovatov.   \\vskip 0.6cm \\noindent {\\large \\bf References} \\vspace{-1.5cm}"
    },
    "0207/astro-ph0207395_arXiv.txt": {
        "abstract": "{ In most cosmological models, primordial black holes ({\\sc pbh}) should have formed in the early Universe. Their Hawking evaporation into particles could eventually lead to the formation of antideuterium nuclei. This paper is devoted to a first computation of this antideuteron flux. The production of these antinuclei is studied with a simple coalescence scheme, and their propagation in the Galaxy is treated with a well-constrained diffusion model. We compare the resulting primary flux to the secondary background, due to the spallation of protons on the interstellar matter. Antideuterons are shown to be a very sensitive probe for primordial black holes in our Galaxy. The next generation of experiments should allow investigators to significantly improve the current upper limit, nor even provide the first evidence of the existence of evaporating black holes. ",
        "introduction": "Very small black holes could have formed in the early Universe from initial density inhomogeneities (Hawking \\cite{Hawking2}), from phase transition (Hawking \\cite{Hawking3}), from collapse of cosmic strings (Hawking \\cite{Hawking4}) or as a result of a softening of the equation of state (Canuto \\cite{Canuto}). It was also shown by Choptuik (Choptuik \\cite{Choptuik}) and, more recently, studied in the framework of double inflation (Kim \\cite{Kim}), that {\\sc pbh}s could even have formed by near-critical collapse in the expanding Universe.  The interest in primordial black holes has been revived in the last years for several reasons. On the one hand, new experimental data on gamma-rays (Connaughton \\cite{Connaughton})  and cosmic rays (Barrau{\\it et al.} \\cite{barrau4}) together with the construction of neutrino detectors (Bugaev \\& Konishchev \\cite{Bugaev}), of extremely high-energy particle observatories (Barrau \\cite{Barrau3}) and of gravitational waves interferometers (Nakamura {\\it et al.} \\cite{Nakamura}) give interesting investigational means to look for indirect signatures of {\\sc pbh}s. On the other hand, primordial black holes have been used to derive interesting limits on the scalar fluctuation spectrum on very small scales, extremely far from the range accessible to CMB studies (Kim {\\it et al.} \\cite{Kim2}, Blais {\\it et al.} \\cite{Po}). It was also found that {\\sc pbh}s are a useful probe of the early Universe with a varying gravitational constant (Carr \\cite {Carr2}). Finally, significant progress has been made in the understanding of the evaporation mechanism itself, both at usual energies (Parikh \\& Wilczek \\cite{Parikh}) and in the near-Planckian tail of the spectrum (Barrau \\& Alexeyev \\cite{Barrau2}, Alexeyev {\\it et al.} \\cite{Stas}, Alexeyev {\\it et al.} \\cite{Stas2}).  For the time being there is no evidence in experimental data in favour of the existence of {\\sc pbh}s in our Universe. Only upper limits on their number density or on their explosion rate have been obtained  (Barrau {\\it et al.} \\cite{barrau4}, MacGibbon \\& Carr \\cite{MacGibbon2}). As the spectra of gamma-rays, antiprotons and positrons can be well explained without any new physics input ({\\it e.g.} {\\sc pbh}s or annihilating supersymmetric particles) there is no real hope for any detection in the forthcoming years using those cosmic-rays. The situation is very different with antideuterons which could be a powerful probe used to search for exotic objects, as the background is extremely low below a few GeV (Chardonnet {\\it et al.} \\cite{chardonnet}, Donato {\\it et al.} \\cite{fiorenza}). Such light antinuclei could be the only way to find {\\sc pbh}s or to improve the current limits. This paper is organized along the same guidelines as our previous study on {\\sc pbh} antiprotons (Barrau {\\it et al.} \\cite{barrau4}), to which the reader is referred for a full description of the source and propagation model used. The main difference is the necessity to consider a coalescence scheme for the antideuteron production. We compute the expected flux of antideuterons for a given distribution of {\\sc pbh}s in our Galaxy, propagate the resulting spectra in a refined astrophysical model whose parameters are strongly constrained and, finally, give the possible experimental detection opportunities with the next generation of experiments as a function of the uncertainties on the model. ",
        "conclusions": "As recently pointed out in Donato {\\it et al.} (\\cite{fiorenza}), antideuterons seem to be a more promising probe to look for exotic sources than antiprotons. In this preliminary study, we show that this should also be the case for {\\sc pbh}s, so that antideuterons may be the only probe to look for  such objects. They should allow a great improvement in sensitivity during the forthcoming years:  a factor 6 better than the current upper limit for AMS and a factor of 40 for GAPS.   Among the possible uncertainties mentioned in Barrau {\\it et al.} (\\cite{barrau4}), the most important one was, by far, the possible existence of a QCD halo around {\\sc pbh}s (Heckler \\cite{heckler}). The latest studies seem to show that this effect should be much weaker (Mac Gibbon {\\it et al.}, in preparation) than expected in Cline {\\it et al.} (\\cite{cline}). The results given in this work should, therefore, be reliable from this point of view.  Nevertheless, two points could make this picture a bit less exciting and deserve detailed studies. The first one is related to secondary antideuterons: the cross-sections used in this work could be slightly underestimated and some other processes could have to be taken into account (Protassov {\\it et al.}, in preparation). This could increase the background  which should be considered with the same propagation model. The second one is that the signal is extremely close to the one obtained with the annihilation of supersymmetric particles as the shape of the spectrum is mostly due to fragmentation processes. In the case of detection, it would be very difficult to distinguish between the two possible phenomena, unless collider data or indirect or direct neutralino dark matter searches have given enough information to fix the supersymmetric parameters.\\\\    {\\bf Acknowledgments}. We would like to thank K. Protassov and R. Duperray for very interesting discussions about antideuteron cross-sections and C. Renault for her great help."
    },
    "0207/astro-ph0207489_arXiv.txt": {
        "abstract": "{We have obtained \\'echelle spectroscopy of 14 Population II objects selected from those previously observed by Bonifacio \\& Molaro (1997). For one object, HD~140283, we obtained exquisite data with the High Dispersion Spectrograph on the Subaru Telescope, with $S/N$ exceeding 1000 per 0.018~\\AA\\ pixel. Li abundances have been determined by spectral synthesis from both the 6708{\\AA} resonance line and also from 6104{\\AA} subordinate feature. Firm detections of the weak line have been made in seven objects, and upper limits are reported for the remainder. Our 6708{\\AA} abundances agree with those reported by Bonifacio \\& Molaro at the 99\\%-confidence level. Abundances from the 6104{\\AA} line hint at a higher Li abundance than that determined from the resonance feature, but this evidence is mixed; the weakness of the 6104{\\AA} line and the large number of upper limits make it difficult to draw firm conclusions. NLTE-corrections increase (rather than eliminate) the size of the (potential) discrepancy, and binarity appears unlikely to affect any abundance difference. The effect of multi-dimensional atmospheres on the line abundances was also considered, although it appears that use of 3-D (LTE) models could again act to {\\em increase} the discrepancy, if one is indeed present. \\keywords {stars: abundances - stars: Population II - stars: interiors}} ",
        "introduction": "\\label{sec-intro}  Population II stars are older and more metal-poor than the Pop. I stars which make up most of the Galactic disk. They are predominantly found in the Galactic bulge, in the spherical halo around the Galaxy, and in globular clusters, although they also make up $\\sim$0.1\\% of the population of nearby stars. They formed early in the life of the Galaxy, and therefore are better indicators of early Li abundances than their younger Pop. I counterparts. This seemed to be supported by the discovery of a Li abundance plateau in these objects by Spite \\& Spite (\\cite{spite82}). The Spite plateau (Li abundance\\,=\\,2.09\\,$^{+0.19}_{-0.13}$\\,dex, Ryan et al. \\cite{ryan00} -- where A(Li) = $\\log_{10}\\frac{N({\\mathrm{Li}})}{N({\\mathrm{H}})} + 12$) includes objects in the temperature range 5600\\,K\\,$\\leq T_{\\mathrm{eff}} \\leq$\\,6400\\,K, and metallicity range [Fe/H]$< -$1.3. Since the plateau Pop. II stars are predicted to have undergone little Li depletion over the course of their lives ({$<$0.1\\,dex according to Ryan et al. 1999 --hereafter \\cite{rnb99}--, $\\leq$0.2\\, dex according to Pinsonneault et al. \\cite{p99}, but possibly as high as 0.3\\,dex according to Salaris \\& Weiss \\cite{salaris01}}) the plateau Li abundance has been suggested to be close to the primordial Li abundance, Li$_{\\mathrm{p}}$. Recent work by \\cite{rnb99} found that the plateau is ultra-thin, but slopes with metallicity ($\\frac{\\mathrm{d(A(Li))}}{\\mathrm{d[Fe/H]}}\\,=\\,+0.118\\pm0.023$), suggesting that the Galactic Li abundance increased slowly with metallicity during the time when Pop. II stars were forming. Li abundances in Pop. II stars cooler than the plateau are depleted. It is also possible that the plateau abundance is the result of stellar Li depletion from a higher value (see e.g. Pinsonneault \\cite{p97}, \\cite{p99}), though why the resulting Li abundance should then be so uniform in these objects is unclear (Th\\'eado \\& Vauclair \\cite{theado01}). Salaris \\& Weiss (2001) have developed models which include atomic diffusion (untempered by mass loss) which they compare with the plateau Li-$T_{\\rm eff}$ trend. They conclude that this mechanism leaves the stars' Li abundances depleted by $\\sim$0.3~dex relative to their initial abundances. They comment that the neglect of both mass-loss and radiative-levitation processes in their models probably leads to this being an overestimation, but more work is required to quantify this. However, the Spite plateau is probably still the best indicator available (after consideration of, and correction for, effects such as diffusion) for the primordial Li abundance, and can still be used to set limits on cosmological parameters such as $\\eta$ (the baryon-to-photon ratio) and $\\Omega_{\\mathrm{B}}$ (the universal baryon density).  There is yet another problem with assuming that the plateau abundance is the same as Li$_{\\mathrm{p}}$: the plateau abundance is derived from only one spectral line, the 6708{\\AA} resonance feature. Recent work by Kurucz (\\cite{kurucz95}), Carlsson et al. (\\cite{carlsson94}), Stuik et al. (\\cite{stuik97}), Uitenbroek (\\cite{uitenbroek98}) and Asplund et al. (\\cite{asplund99}) has suggested that this line might be unrepresentative of the actual Li abundance within the stars, due to omissions in our understanding and modelling of stellar atmospheres. Recent works on halo stars (Bonifacio \\& Molaro 1998, hereafter \\cite{bm98} -- and next subsection) and young cluster stars (Soderblom et al. \\cite{s93a}; Russell, \\cite{russell96}; Ford et al. \\cite{ford02}) have used the weak (excitation potential\\,=\\,1.8\\,eV) \\ion{Li}{i} subordinate line at 6104{\\AA} in addition to the 6708{\\AA} line (hereafter 6708 and 6104 are taken to mean `the \\ion{Li}{i} line at 6708{\\AA}' and `the \\ion{Li}{i} line at 6104{\\AA}' respectively), to attempt to determine Li abundances. This paper builds on those previous studies.  \\subsection{Multi-dimensional atmospheres: the story so far...} \\begin{itemize} \\item Kurucz (\\cite{kurucz95}) noted that one-dimensional model atmospheres will only work well if they have the same temporally- and spatially-averaged behaviour as the species they are used to predict, which might not be the case for Li. Kurucz considered the idea that conventional abundance analysis is in error, and suggested that model atmospheres could underestimate the amount of ionized Li by an order of magnitude, which would lead to Li abundance measurements also being lower than the `true' value by around 1\\,dex. Kurucz cautions us that ``{\\em since very few lines have atomic data known accurately enough to constrain the model, a match does not necessarily mean that the model is correct.}'' \\item Carlsson et al. (\\cite{carlsson94}) studied the effect of non-local thermodynamic equilibrium (NLTE) on the formation of Li lines. They used a standard flux-constant plane-parallel atmosphere to model Li line formation, but included processes ignored by assumptions of local thermodynamic equilibrium (LTE). Inclusion of NLTE processes appears to have some effect on Li lines. NLTE corrections are largest for cool ($T_{\\mathrm{eff}} \\sim $ 4500\\,K), metal-poor ([Fe/H] $\\sim$ $-$2) objects, although the magnitude of the effect varies with both temperature and metallicity. Generally, the corrections have opposite signs for 6104 and 6708, increasing the abundance derived from the subordinate line by up to 0.06\\,dex (at A(Li) = 2.2) while decreasing it in the case of the resonance feature (by around 0.004\\,dex at A(Li) = 2.2 for an object with $T_{\\mathrm{eff}}$ = 6000\\,K, $\\log\\,g$ = 4.0, and [Fe/H] = $-$2.0). \\item Stuik et al. (\\cite{stuik97}) continued the work of Carlsson et al., investigating the sensitivity of \\ion{Li}{i} and \\ion{K}{i} to activity in Pleiades stars. They stress that the variations in Li line strength in the Pleiades cannot be attributed to a spread in the stellar abundances until such time as the line-formation predictions and atmospheric models can be shown to be accurate. One test of this is the use of the potassium resonance line (which should not show any abundance spread as K is not expected to be depleted in these stars) as a proxy for Li. This line is formed in similar conditions and transitions to the Li resonance feature, and so should be affected similarly. Their results suggest that while Li and K are not sensitive to the direct effects of chromospheric activity, they can be strongly affected by temperature variations deeper in the stellar atmosphere, and by interactions between regions of different temperature stratification. Such effects are not included in the conventional one-dimensional models, and therefore might account for some of the effects seen. Pop. II stars should not be as active as young cluster objects, but there are likely to be similar (related or unrelated) omissions in our modelling of Pop. II stars. Kiselman (\\cite{kiselman97, kiselman98}) used a 3-D solar-granulation snapshot to investigate the effects of departures from LTE in line formation on 6708. He reported marked line-strength variations over the granulation pattern for NLTE and LTE. He also produced syntheses of the line under these conditions and compared them to solar observations, finding that calculations which included the detailed modelling of the line radiative transfer matched observations better than those which neglected it. \\item Uitenbroek (\\cite{uitenbroek98}) worked in a similar vein to Stuik et al. and Kiselman, considering the effect of convective surface inhomogeneities on the formation of the Li resonance line in the solar case. Uitenbroek uses 1.5- and two-dimensional NLTE calculations to investigate the effects of granulation on the lines. Again, granulation is unlikely to play a significant role in Pop. II-star Li modelling, but the fact that its omission in Pop. I models leads to a difference in abundance (although typically less that 0.1\\,dex) shows that our understanding of stellar atmospheres is not complete. \\item BM98 detected 6104 in the Pop. II star \\object{HD\\,140283}. This feature is important since it is formed deeper in the atmosphere than the resonance line, and can therefore be used to test our models, since both lines should present the same abundance. Bonifacio \\& Molaro's analysis suggested that the same abundance could be used to adequately reproduce both lines using only 1-D, homogeneous models. They took this to imply that these `simple' models were essentially correct. \\item Asplund et al. (\\cite{asplund99}) have used 3-D, time-dependent surface-convection simulations of two Pop. II objects, HD\\,140283 and \\object{HD\\,84937} to investigate the effect of multi-dimensional modelling on Li abundance in metal-poor stars. They report that, as might be expected, three-dimensional model atmospheres have a different temperature structure to one-dimensional models. The effect of their 3-D, LTE models on the stellar Li is to decrease the abundance measured from the resonance line by 0.2-0.35\\,dex, relative to that obtained using 1-D models (for models with $T_{\\mathrm{eff}}$\\,=\\,5690\\,K, $\\log\\,g$\\,=\\,3.67, [Fe/H]\\,=\\,$-$2.5 and microturbulence ($\\xi$)\\,=\\,1\\,km\\,s$^{-1}$ in the case of HD\\,140283, and $T_{\\mathrm{eff}}$\\,=\\,6330\\,K, $\\log\\,g$\\,=\\,4.04, [Fe/H]\\,=\\,$-$2.25 and $\\xi$\\,=\\,1\\,km\\,s$^{-1}$ for HD\\,84937). Note that Asplund et al. \\cite{asplund99} quote $\\log\\,g$ for $g$ in units of m\\,s$^{-1}$; for consistency with other works we use cm\\,s$^{-1}$. However, 3-D NLTE corrections almost completely cancel the 3-D LTE effect, leading to $<$\\,0.1\\,dex change (in these two cases) to the 1-D LTE abundance (Asplund et al. \\cite{asplund00}). \\end{itemize} We seek here to extend the work of \\cite{bm98}. Firstly, their sample consists of only one star, which might or might not be typical of other Pop. II objects. Secondly, they do not fit 6104 independently of the resonance feature, but synthesize it at the abundance determined from the stronger feature. 6104 is difficult to measure due to its weakness and proximity to Fe and Ca lines. At the signal-to-noise ratio ($S/N$) of the \\cite{bm98} spectrum it is likely that the 6104 Li abundance can only be determined to $\\pm$0.2\\,dex (1$\\sigma$). As this is comparable with the likely size of problems with the models and larger than the uncertainties claimed by some authors in Li$_{\\mathrm{p}}$, a more sensitive comparison of 6708 and 6104 is desirable. Because of the weakness of the line, and difficulties achieving higher $S/N$, it is sensible to measure 6104 in a number of objects to reduce the statistical errors (see Sect.~\\ref{sec-abundanceerrors} and~\\ref{sec-EWerrors}).  Following on from the work of \\cite{bm98} we have obtained a larger sample of plateau stars (and three cooler objects) and analysed them with a variety of 1-D homogeneous models, fitting both lines independently of each other to determine if the abundances from the two lines do agree. ",
        "conclusions": "\\label{sec-discuss}  \\subsection{6104 {\\em versus} 6708} \\label{sec-4v7}  Using consistent reduction and analysis techniques, we have measured 6104 in seven of our fourteen Pop.\\,II stars and obtained upper limits in the others. We have measured 6708 in all the objects.  There is evidence of a discrepancy in Li abundance between the two lines for some of our sample objects. In Fig.~\\ref{fig-4v7} which plots our 6104 abundances {\\em versus} those from 6708, it appears that the 6104 abundances of some objects (e.g. HD\\,219617 and G80-28) are higher than those from 6708. Figure 5 might also give the impression that we have uncovered a systematic discrepancy between the 6104 and 6708 abundances. However, part of this is due to the large number of upper limits, which make it impossible to draw such a conclusion. Indeed, as we describe below, there are some objects for which a large discrepancy can be ruled out. The discrepancy between 6104 and 6708 abundances is most apparent, reaching $\\sim$0.5\\,dex, at low values of A(Li) where, admittedly, there is more potential for overestimation of the 6104 equivalent width. At higher abundances the points lie closer to the one-to-one relation, but several are still $\\sim$2$\\sigma$ above it.  \\begin{figure} \\resizebox{9cm}{!}{\\includegraphics{h3440f5.eps}} \\caption{ A(Li)$_{6104}$ {\\em versus} A(Li)$_{6708}$ for sample stars. The solid line is the 1-to-1 relation. Since error bars which are linear in EW become non-linear in A(Li), especially for a line as weak a 6104, we show both 1$\\sigma$ and 2$\\sigma$ error bars for the abscissa. Upper limits are shown as inverted triangles. {\\it Upper panel}: LTE abundances. {\\it Lower panel}: NLTE abundances. } \\label{fig-4v7} \\end{figure}    A measurement of the 6104 line is more likely to be made not only in objects where the line {\\em is} stronger, but also when noise acts to make it {\\em appear} stronger. Because of this latter bias, it is not surprising that {\\it some} of the 6104 abundances exceed those of 6708. Nevertheless, it is disquieting that in three of the fourteen cases, the abundance obtained from the 6104 line exceeds that from 6708 by more than twice the formal error, {\\em i.e.} $>$2$\\sigma$. This can be seen in Fig.~\\ref{fig-specfit4}: the 6708 abundance clearly {\\em does not} fit the observations in some cases, most notably for HD\\,219617 and G\\,80-28. If the error distribution was Gaussian, and our error estimates were reliable, there would be only 2\\%\\ probability of a 6104 abundance exceeding that from 6707 by 2$\\sigma$. This would give rise to fewer than one discrepant case in fourteen. However, there are also cases where such discrepancies cannot exist: HD\\,140283 (discussed below), HD\\,108177 and HD\\,94028 all have 6104 abundances which are consistent with the 6708 abundances. Overall, this leads to a very mixed picture where the abundances from the two lines agree for some objects, some objects have A(Li)$_{6104}$\\,$>$\\,A(Li)$_{6708}$, and some objects might have A(Li)$_{6708}$\\,$>$\\,A(Li)$_{6104}$(although the latter possibility is defined by only one object in our sample; the other upper limits neither contradict nor constrain this possibility). The lack of a clear picture of what is going on adds a further complication into any interpretation of the data, in that whatever causes the elevated 6104 abundances in some stars is not present in other objects.  We have plotted the abundances for the two lines against $T_{\\mathrm{eff}}$ and [Fe/H] in Fig.~\\ref{fig-teff-met}. In the case of the  A(Li)-{\\em versus}-$T_{\\mathrm{eff}}$ plot, the plateau can be seen for both lines, with similar abundances for objects with $T_{\\mathrm{eff}} \\geq$\\,5800\\,K. The 6104 abundances appear to show a separate plateau $\\la$0.5\\,dex above the 6708 plateau. As with Fig.~\\ref{fig-4v7}, this impression is accentuated by the high number of upper limits, especially at low temperature, but concentrating only on the firm detections shows that the discrepancy is not an obvious function of effective temperature or metallicity. The Li abundances in the plateau objects all appear approximately constant with metallicity, although slopes in either direction could potentially be present. The behaviour of objects at the cool edge of the plateau is much as expected, with Li abundances from both lines decreasing with decreasing temperature, consistent with HD\\,188510 and HD\\,64090 being cool enough that some Li depletion is likely to have occurred during their lifetimes.  \\begin{figure*} \\resizebox{17cm}{!}{\\includegraphics{h3440f6.eps}} \\caption{NTLE-corrected A(Li) {\\em versus} $T_{\\mathrm{eff}}$ (upper panel) and [Fe/H] (lower panel) for sample stars. Filled symbols represent 6708 abundances and open symbols show 6104 values. Inverted triangles indicate 6104 upper limits.} \\label{fig-teff-met} \\end{figure*}  In order to investigate the abundance disparity further, more firm detections would be needed, although obtaining them is complicated by the increased difficulty in measuring such weak 6104 lines. Achieving higher $S/N$ exceeding a few hundred not only requires longer and longer exposures, but also requires highly repeatable flat-fields that correct for fringing and other defects to better than 0.1\\%.  \\subsection{HD 140283} \\label{sec-hd140}  The data for HD\\,140283 do not show a discrepancy between the 6104 and 6708 abundances. We also used our models to obtain abundances for HD\\,140283 using the EWs and model parameters quoted in \\cite{bm97} and \\cite{bm98}. In contrast to the general behaviour noted in Sect.~\\ref{sec-4v7}, we find A(Li)$_{6708}$\\,=\\,2.16$\\pm$0.01 (from EW$_{6708}$\\,=\\,47.5$\\pm$0.6\\,m{\\AA}), slightly higher than\\,A(Li)$_{6104}$\\,=\\,2.10$^{+0.07}_{-0.08}$ (from EW$_{6104}$\\,=\\,1.8$\\pm$0.3\\,m{\\AA}) but in any case in agreement within the errors. The abundance reported by \\cite{bm98} for both lines in HD\\,140283 is A(Li)\\,=\\,2.14$\\pm$0.13, which is also in agreement with our calculated abundances.  Our observation HD\\,140283 allows us to compare our results for this object directly with that obtained by BM98. The main differences between our analysis and theirs are: the significantly-higher $S/N$ of our spectrum ($\\sim$1100 per 0.018{\\AA} pixel, compared to BM98's 360 per 0.04{\\AA} pixel); the $gf$ values used in the syntheses (see Table~\\ref{table-linelist} for our values, and Sect.~2 of BM98); their use of $\\alpha$-enhanced opacities {\\em versus} our use of standard opacities (Sect.~\\ref{sec-bmcomparison}); our use of ABO broadening parameters. We have also included an \\ion{Fe}{ii} line in our synthesis of the 6104 region, although at the metallicity of the object this should make no difference. The analyses shared identical $T_{\\mathrm{eff}}$, $\\log g$, [Fe/H], $\\xi$ and $v \\sin i$ values. Despite the differences between the analyses, we have obtained abundances which are completely consistent with those reported by BM98, which is not too surprising since the differences are all in parameters which have only a small effect on abundance. Even so, it shows that there is a good degree of consistency between the two studies. Our fits to both 6708 and 6104 are remarkably close to the observations (Figs.~\\ref{fig-specfit7} and \\ref{fig-140zoom}), suggesting that our line profiles are reliable.  \\begin{figure} \\resizebox{!}{!}{\\includegraphics{h3440f7.eps}} \\caption{Observed (solid) spectrum and synthetic fit (dashed) to 6104 in HD\\,140283.} \\label{fig-140zoom} \\end{figure}  In concluding this subsection, we reemphasize that the 6104 lines are all very weak, and it is possible that we have underestimated the errors, but we discuss below the implications of a genuine discrepancy.  \\subsection{Model atmospheres} \\label{sec-disc.atmospheres}  If the possible discrepancy between the 6104 and 6708 lines in some of the objects is real, what could explain it? Errors in $T_{\\mathrm{eff}}$ of $\\pm$100\\,K lead to abundance {\\it differences} between 6708 and 6104 of $\\pm$0.03\\,dex at 6000\\,K, and $\\pm$0.05\\,dex at 5000\\,K, with 6708 abundances changing more than 6104 abundances. This has already been shown to be too small; a 1000\\,K temperature increase would be required to bring the most discrepant lines into agreement (see Sect.~\\ref{sec-abundanceerrors}), and the inclusion of other model uncertainties contributes less than 0.03\\,dex.  An alternative explanation for our result is that the model atmospheres do not adequately represent the stars we are studying, and a change in temperature gradient within the stellar atmospheres might explain our observations. Asplund et al. (\\cite{asplund99}) applied 3-D model atmospheres to two Pop.\\,II stars, HD\\,140283 and HD\\,84937, noting the predicted effect on the Li abundances relative to those from 1-D models. 3-D models alter the temperature stratification in a stellar atmosphere, with the greatest variations between 1-D and 3-D models occurring at the inner and outer boundaries (see lower panel of Fig.~\\ref{fig-cf}). The upper panel of Fig.~\\ref{fig-cf} shows the flux contribution functions ${\\mathrm{d}}(F_\\nu)$/${\\mathrm{d}} \\log_{10}(\\tau_{5000})$ computed for four wavelengths in the spectrum of HD\\,140283, using our model atmosphere. These show the range of depths over which the spectrum forms at a given wavelength (e.g. Gray 1992, Chapter 13). The greater equivalent width of 6708 and its lower excitation potential both lead to it forming on average further out than 6104.  \\begin{figure} \\resizebox{!}{!}{\\includegraphics{h3440f8.eps}} \\caption{{\\it Upper panel}: Flux contribution functions for our 1-D model of HD\\,140283, which shows the range of depths over which the spectrum forms. Shown are contributions functions for the continuum near 6708\\,\\AA\\ (thin solid curve), the Li\\,6708 line (heavy solid curve), the continuum near 6104\\,\\AA\\ (thin dashed curve), the Li\\,6104 line (heavy dashed curve). (The contribution function plotted for each spectral line is a weighted average of the contribution functions at several wavelengths within the line profile.) {\\it Lower panel}: Temperature profiles for our 1-D model (open circles), and the 1-D (solid curve) and 3-D (dashed curve) models of Asplund et al. (1999). Dotted vertical lines are to guide the eye. } \\label{fig-cf} \\end{figure}  The temperature differences between the 1-D and \\nobreak{3-D} models of Asplund et al. are fairly small at optical depths between $\\frac{1}{30} < \\tau_{5000} < \\frac{2}{3}$. However, the lower temperatures of the 3-D models beyond this range, where part of the flux at 6708 emerges, result in a stronger line forming in the 3-D models. (Switching to a 3-D model would also modify the contribution function from that shown). Asplund et al. found that, for a given EW, the 6708-derived abundance in HD\\,140283 would be 0.34\\,dex lower if derived using 3-D LTE models rather than 1-D LTE atmospheres.  Although 6104 could also be affected by the differences between 1-D and 3-D models, 6104 is so weak (EW\\,=\\,1.8\\,m\\AA) that it is barely distinguishable from the neighbouring continuum, and the magnitude of the effect is probably less for this line. Asplund (priv. comm.) confirms this expectation. In this case, use of 3-D LTE models would decrease the 6708 abundance more than the 6104 abundance, {\\em increasing} the abundance discrepancy identified in Sect.~\\ref{sec-4v7} (if it is real) between the two lines.  1-D NLTE effects, which have been corrected for in our results, have been discussed in three other studies of low-metallicity stars, namely Kurucz (\\cite{kurucz95}), \\cite{bm98}, and Molaro et al. (\\cite{molaro95}). As can be seen from our Fig.~\\ref{fig-4v7} the NLTE corrections are relatively small, affect both lines differently, and vary depending on temperature, metallicity and Li abundance. In all cases, however, application of the NLTE correction acts to {\\em increase} the 6708 and 6104 abundance difference, rather than reducing it.  Asplund et al's (1999) suspicion that overionisation of Li in 3-D models might reverse the 0.3\\,dex difference initially found for the 6708 line between the 1-D LTE and 3-D LTE models proved to be correct (Asplund 2000). However, corrections for overionisation may be expected to affect both the 6104 and 6708 lines, though not necessarily equally (due to the different depths of formation). It is clear that detailed calculations using 3-D NLTE atmospheres will be required for each programme star to fully evaluate the competing effects and their different action on the two lines.  To produce the discrepancy via temperature gradients alone, what would be required is a process which changes the temperature gradient in opposite ways for the two lines, such that it increases the temperature in the outer region where 6708 is formed, thus weakening that line and leading to an underestimate of its abundance, and decreasing the temperature in the deeper region where 6104 forms. As noted above, this is opposite to the behaviour seen in the 3-D atmospheres of Asplund et al. (\\cite{asplund99}).  \\subsection{Binarity} \\label{sec-binarity}  There are no clear systematic effects of binarity on Li abundances derived from either line. Of our seven detections, three are in objects classified as binaries by Carney (\\cite{carney83}): HD\\,84937, HD\\,219617 and HD\\,188510; both HD\\,84937 and HD\\,188510 have also been classified as definite binaries by Stryker et al. (\\cite{stryker85}). One object in our sample, BD+21\\,607, is near the significance criterion for binarity (Stryker et al. \\cite{stryker85}), and the remaining three have not been classified as binaries by either study. The upper limits do nothing to clarify the picture with respect to binarity either: LP\\,608-62, HD\\,94028, HD\\,108177 and HD\\,210891 have all been classified as either definite or probable binaries by Carney or Stryker et al., while the remainder are probably not binaries. Clearly some other solution is required to explain the abundance discrepancy between the two lines.  \\subsection{Contrast with Pop. I stars} \\label{sec-popi}  In a related paper (Ford et al. \\cite{ford02}) we found evidence that 6104 gave higher Li abundances than 6708 for some young Pleiades G/K dwarfs, whilst for others there seems to be reasonable agreement. This is very similar to the situation we find here for Pop. II stars.  For the Pleiades stars we were able to bring the abundances estimated from the two lines into agreement by introducing cool starspots in to the atmospheric models. We did this in a simple way, by modelling the atmosphere as two one-dimensional components with differing temperatures and surface areas. In young Pop. I stars there is plenty of evidence (from Doppler imaging and light-curve modulation) that such spots exist, and so it is sensible to incorporate them in the models.  Our conclusion was that plausible spot coverages and temperatures could explain the abundance discrepancies we saw in one-component models, although we had insufficient data to determine whether the stars actually did possess spots with the right properties.  We do not expect to find large, magnetically-generated spot regions associated with Pop. II stars. They are old, have spun down and their convection zones are much thinner than the Pop. I objects we considered in the Pleiades. For these reasons we do not think it appropriate to consider cool starspots as an explanation for the abundance discrepancies we see in Pop. II stars. That we still see a discrepancy (in some but not all stars) in Pop. II objects is likely to be an indication that something else is still wrong with our atmospheric models. It may also be the case that these additional problems are also present in Pop. I stars, and that the starspot interpretation is only part of the story. Indeed, even after the introduction of starspots on to the  Pop. I stars we were unable to reduce the large dispersion that is seen in their Li abundances.  \\subsection{Effects on cosmological parameters: $\\eta$ and $\\Omega_{\\mathrm{B}}$} \\label{sec-cosmology}  If we consider that the 6104-derived abundances are less likely to be affected by multi-dimensional model-atmosphere variations than those from 6708, then the plateau abundance inferred from our 6104 {\\it detections} would be A(Li)$\\sim$2.5. (We do not include possible depletion factors in the values quoted below, although our 6104 abundance is the same as that derived by Salaris \\& Weiss (\\cite{salaris01}) from their diffusion models, constrained by observations of 6708). This would lead to two possible values for the baryon-to-photon ratio ($\\eta$: values given here are $\\eta_{10}\\,=\\,10^{10}\\eta$) of either $\\sim$1.0 or $\\sim$7.0, leading to values of 3.7$\\times\\,10^{-3}h_{\\mathrm{o}}^{-2}$ or 25.5$\\times\\,10^{-3}h_{\\mathrm{o}}^{-2}$ for $\\Omega_{\\mathrm{B}}$, the Universal baryon density (where $h_{\\mathrm{o}}$\\,=\\,$\\frac{H_{\\mathrm{o}}}{100}$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and the microwave background temperature is 2.73\\,K). Values for A(Li)$_{\\mathrm{plateau}}$\\,=\\,2.09 (Ryan et al. \\cite{ryan00}) would be: $\\eta_{10}$\\,$\\sim$\\,1.2 or 6.0; $\\Omega_{\\mathrm{B}}$\\,$\\sim$\\,4.4$\\times10^{-3}h_{\\mathrm{o}}$ or 21.9$\\times10^{-3}h_{\\mathrm{o}}^{-2}$, with the $\\sim$0.4\\,dex change in A(Li) affecting $\\Omega_{\\mathrm{B}}$ by $\\sim$16\\%. In either case, this assumes that the plateau abundance is indicative of the primordial Li abundance. However, this assumption may be unsound. The work of \\cite{rnb99} found that the plateau slopes with metallicity, suggesting that the material from which the stars formed was gradually being metal enhanced over the time when Pop. II objects were being created. Despite the difficulties in measurement of weak lines, abundances from low-metallicity stars, which we assume to have formed earliest, still provide the best starting point for determinations of Li$_{\\mathrm{p}}$ and related cosmological parameters.    We have obtained data for a sample of 14 Pop. II stars with high signal-to-noise ratios, which were all consistently reduced and analysed. \\begin{itemize} \\item We measured the 6708{\\AA} \\ion{Li}{i} resonance line in all our sample objects. We compared our EWs and abundances with those of Bonifacio \\& Molaro (\\cite{bm97}), finding good agreement between the values at a 99\\% confidence level. \\item We also measured the 6104{\\AA} \\ion{Li}{i} subordinate line in seven of the stars, obtaining upper limits for the rest of the sample; \\item Abundances from the 6104{\\AA} line hinted at a systematically higher Li abundance than those from the 6708{\\AA} line, with some stars discrepant by up to $\\sim$0.5\\,dex, while in others the discrepancy is no larger than $\\sim$0.1\\,dex. This trend was very weak, and a large preponderance of upper limits prevent any firm conclusions from being drawn; \\item This difference cannot be explained by including NLTE-corrections, which actually make the discrepancy larger. Binarity does not appear likely to affect the abundance difference; \\item The 3-D model results of Asplund et al. (\\cite{asplund99}) for Pop. II stars HD\\,84937 and HD\\,140283 suggest that the use of multi-dimensional atmosphere models would most likely increase any abundance discrepancy between the lines relative to that inferred from one-dimensional models. Our results can be used to place strong constraints on the 3-D modelling of these stars, perhaps indicating areas where further improvements could be made (e.g. improved NTLE-modelling, inclusion of magnetic fields). As they stand, however, we do not believe that 3-D models can explain our results. \\end{itemize}"
    },
    "0207/astro-ph0207440_arXiv.txt": {
        "abstract": "X-ray Polarimetry is almost as old as X-ray Astronomy. Since the first discovery of X-ray sources theoretical analysis suggested that a high degree of linear polarization could be expected  due either to the, extremely non thermal, emission mechanism or to the transfer of radiation in highly asymmetric systems. The actual implementation of this subtopic was, conversely, relatively deceiving. This is mainly due to the limitation of the conventional techniques based on the Bragg diffraction at $45^{o}$, or on Thomson scattering around $90^{o}$. Acually no X-ray Polarimeter has been launched since 25 years. Nevertheless the expectations from such measurement on several astrophysical targets including High and Low Mass X-Ray Binaries, isolated neutron Stars, Galactic and Extragalactic Black Holes is extremely attractive. We developed a new technique to measure the linear polarization of X-ray sources. It is based on the visualization of photoelectron tracks in a, finely subdivided, gas filled detector (micropattern). The initial direction of the photoelectron is derived and from the angular distribution of the tracks the amount and angle of polarization is computed. This technique can find an optimal exploitation in the focus of XEUS-1. Even in a very conservative configuration (basically the already existing prototype) the photoelectric polarimeter could perform polarimetry at $\\%$ level on many AGNs. Further significant improvements can be expected from a technological development on the detector and with the use of XEUS-2 telescope. ",
        "introduction": "Historically we can group the measurements performed on Astronomical X-ray Sources into four groups:  \\begin{itemize} \\item Timing Photometry (Geiger, Proportional Counters, MCP) with Rockets, UHURU, Einstein, EXOSAT, ASCA, SAX,  XMM, Chandra.  \\item Imaging:  Pseudo-imaging (modulation collimators, coded masks), SAS-3, XTE-ASM, SAX-WFC, HETE-2  Real Imaging (grazing incidence optics + Position Sensitive Detectors: IPC, MCA, CCD) with Rockets, Einstein, EXOSAT, ROSAT, ASCA, SAX, Chandra, Newton.  \\item Spectroscopy:  Non dispersive (Proportional Counters, Si/Ge and CCD) Rockets, Einstein, EXOSAT, HEAO-3, ASCA, SAX, Chandra, Newton.  Dispersive: (Bragg, Gratings) Einstein, Chandra, Newton  \\item Polarimetry   (Bragg, Thomson/Compton) with rockets Ariel-5, OSO-8 \\end{itemize}  While in the domain of Photometry, Imaging and Spectroscopy the observing techniques have been tremendously improved, Polarimetry has only been based on the same, conventional techniques, producing important but very limited results. In fact, after OSO-8, no astronomical Polarimeter has been flown any more. ",
        "conclusions": "We conclude that with the new MICROPATTERN device, the Polarimetry of Astrophysical sources is now feasible, provided that a high throughput optics is used. This will open a new window in the sky and dramatically improve our understanding of Physics of X-ray emitting regions around NS and Black Holes. With XEUS-1 optics, and with moderate assumptions on technological developments, polarimetry to the $\\%$ level of tens of AGNs will be feasible.  Therefore we think that the inclusion of a MICROPATTERN photoelectric  polarimeter in the baseline payload for XEUS-1 should be seriously considered."
    },
    "0207/astro-ph0207506_arXiv.txt": {
        "abstract": "The BL Lac object H1426+428 was recently detected as a high energy $\\gamma$-ray source by the VERITAS collaboration \\citep{horan02}. We have reanalyzed the 2001 portion of the data used in the detection in order to examine the spectrum of H1426+428 above 250 GeV. We find that the time-averaged spectrum agrees with a power law of the shape $$ (\\frac{\\mathrm{d}F}{\\mathrm{d}E})(E) = 10^{-7.31 \\pm 0.15_{\\mathrm{stat}} \\pm 0.16_{\\mathrm{syst}}}\\cdot E^{{-3.50 \\pm 0.35_{\\mathrm{stat}}\\pm 0.05_{\\mathrm{syst}}}} \\ \\mathrm{m}^{-2}\\mathrm{s}^{-1}\\mathrm{TeV}^{-1} $$ The statistical evidence from our data for emission above 2.5 TeV is 2.6 $\\sigma$. With 95\\,\\% c.l., the integral flux of H1426+428 above 2.5 TeV is larger than 3\\% of the corresponding flux from the Crab Nebula. The spectrum is consistent with the (non-contemporaneous) measurement by \\citet{aharon02} both in shape and in normalization. Below 800 GeV, the data clearly favours a spectrum steeper than that of any other TeV Blazar observed so far indicating a difference in the processes involved either at the source or in the intervening space. ",
        "introduction": "The BL Lac object H1426+428 was discovered at optical wavelengths with a redshift of 0.129 by \\citet{remil89}. Measurements of the spectral energy distribution  of this close object are still sparse and do not sufficiently cover the expected broad two-peak structure that is familiar from other BL Lac objects (see, e.g., \\citet{donato01}). The first peak is expected at X-ray energies and is thought to represent the synchrotron emission from relativistic electrons in the source. The second peak is expected at $\\gamma$-ray energies and is explained in  so-called leptonic models as stemming from low-energy photons which are inverse-Compton scattered to $\\gamma$-ray energies by the same population of relativistic electrons that causes the synchrotron radiation. Alternative models attribute the second peak and also part of the X-ray emission to processes involving protons which are co-accelerated with the X-ray-emitting electrons or even partially produce them.  For recent reviews see, e.g., \\citet{ghisell00}, \\citet{rachen00}, \\citet{sikora00}.  A lower limit on the position of the X-ray peak of H1426+428 was placed by \\citet{costa01} at 100 keV. Currently the only published measurements above 100 keV come from Cherenkov telescopes above 280 GeV \\citep{horan02} and above 700 GeV \\citep{aharon02}.  \\citet{costa02} predict the peak of the $\\gamma$-ray emission of H1426+428 to be at several ten GeV.  The study of the high-energy spectrum of this ``new'' TeV source is especially interesting since its redshift is four times as large as that of Mkn 421 and Mkn 501, the only other extragalactic objects detected at TeV energies with good spectral information. Due to this larger distance, it is expected that signs of absorption of the $\\gamma$-radiation via interaction with the intergalactic optical and infra-red background (IIRB) will be more pronounced, possibly permitting one to infer constraints on the IIRB photon density.  In this letter, we present the results of a spectral reanalysis of H1426+428 observations made with the Whipple 10 meter $\\gamma$-ray telescope  on Mt. Hopkins, Arizona, in the first half of 2001. We use our standard method described in \\citet{mohanty98} for deriving the spectrum. Due to the weakness of the source and the special interest in the emission at the highest energies, we then examine the emission above 2.5 TeV with specially developed cuts. ",
        "conclusions": "\\begin{figure} \\plotone{f4.eps} \\caption{\\label{fig-finalspec} The differential energy spectrum of H1426+428 as measured in this analysis (filled points and upper limits) in 2001  and by HEGRA (open squares) in 1999/2000 (Aharonian et al. 2002). The 84\\% confidence level upper limits were obtained using the method by \\protect\\cite{helene83}. The solid line is the result of a power law fit to all points except the upper limit at the lowest energy.  The dashed line shows the scenario of the latter power law with an additional abrupt (super-exponential) cutoff still consistent with all data at the 95\\,\\% confidence level. The minimum cutoff energy found for this scenario is 6 TeV. Also shown is the power-law fit to the HEGRA points only (dotted) and the power-law fit to the Whipple points only (dot-dashed). } \\end{figure}  The spectral index found in this analysis agrees well with the first result ($\\alpha = 3.55\\pm0.5$) published in \\citet{horan02}. The flux normalization constant is larger than that in \\citet{horan02} but well within the statistical and systematic errors of this result which was obtained using a different analysis method and different Monte Carlo data.  The only other published detection of H1426+428 at $\\gamma$-ray energies comes from HEGRA \\citep{aharon02} and is based on a dataset of similar size. The two results are compared in the following.  From a total of 44.4 hours HEGRA obtains a total number of excess events of 199.2 after loose cuts for their spectral analysis. The significance is 4.3 $\\sigma$. We obtain a slightly higher significance  and an excess of 1540 events from 38.1 hours of data using similarly loose cuts. The relative numbers of events recorded by HEGRA and Whipple are consistent with the difference in energy threshold which is a factor of $\\approx 2.5$ . An estimation of the integral spectral index from these numbers gives a value of 2.2 which already hints that the spectrum must be steep if there was no significant change in the overall state of the source between 1999/2000 and 2001.  Direct comparison of the individual spectral points (Figure \\ref{fig-finalspec}) shows a very good agreement indicating that there really was little change in the overall state of the source between the observing periods. Given the 30\\% systematic errors of the absolute flux calibration of both our and the HEGRA measurement, this is consistent with the fact that the $\\gamma$-ray rates measured by \\citet{horan02} for the 2000 and the 2001 H1426 dataset differ only by a factor $1.5$.  To make use of the total available information, we perform a fit of a power-law to the combined points from HEGRA and our analysis. No adjustments to the overall normalization of any of the datasets was performed since the points near 790 GeV agree very well. The errors on the Whipple points included in this fit take into account both the Gaussian error of the number of excess events ($\\sqrt{N_{\\mathrm{on}}+ N_{\\mathrm{off}}}$) and the uncertainty stemming from the energy estimation for the background. The fit to all points (still excluding the upper limit at the lowest energies) results in a spectrum which is essentially identical  with what we obtain from our points only. However, the statistical errors are smaller: $$ \\frac{\\mathrm{d}F}{\\mathrm{d}E} = 10^{-7.36 \\pm 0.07_{\\mathrm{stat}}}\\cdot E^{{-3.54 \\pm 0.27_{\\mathrm{stat}}}} \\ \\mathrm{m}^{-2}\\mathrm{s}^{-1}\\mathrm{TeV}^{-1} $$ The reduced $\\chi^2$ of this fit to 11 points is 0.94 . The measurements are consistent with the assumption that the spectrum of H1426 is a continuous powerlaw between 250 GeV and 17 TeV.  In order to further quantify the evidence for high-energy emission, we introduce a super-exponential cutoff in the measured power-law. This cutoff can be described analytically ($\\exp(-(E/E_0)^2)$) and is resonably abrupt such that it can serve as a parameter describing the energy above which there is no emission from the source. By keeping the nominal values of the power-law fit for normalization and spectral index and reducing the cutoff energy $E_0$ until the $\\chi^2$ of a fit to all significant points (i.e. excluding the points shown in  Figure \\ref{fig-finalspec} as upper limits) has increased to 14.1 (95\\% confidence level for 7 degrees of freedom) we obtain $E_0 = 5.5$\\,TeV.  Independently, we verify if this value for $E_0$ is consistent with the fact that we observe from H1426 \\  $12\\pm5$\\,\\% of the Crab Nebula $\\gamma$-rate above 2.5 TeV (see section \\ref{sec-above2.5}). Taking into account the larger spill-over effects caused by the finite energy resolution of our instrument and the steep spectrum of H1426 and furthermore the 15 \\% uncertainty of our absolute energy calibration, we find that a 95\\,\\% confidence level lower limit on the integral flux of H1426 above 2.5 TeV can be put at 3\\,\\% of the Crab Nebula flux. Using our measurement of the Crab Nebula flux with the same instrument, this corresponds to $$ F_{\\mathrm{H1426}}(E > 2.5 \\mathrm{TeV}) > 1.06 \\times 10^{-9} \\mathrm{m}^{-2}\\mathrm{s}^{-1} \\ \\   (95\\,\\%\\, \\mathrm{c.l.}) $$ Varying the cutoff energy $E_0$ introduced above until the integral flux above 2.5 TeV has the value of this lower limit, we find a value of $E_0 = 6$\\,TeV which independently confirms the result from above. The spectrum with $E_0 = 6$\\,TeV is shown as a dashed line in Figure \\ref{fig-finalspec}.   The HEGRA collaboration attempts to reconstruct the intrinsic energy spectrum of the source by assuming a  model for the intergalactic IR background and calculating the expected absorption for $\\gamma$-rays coming from H1426 (redshift 0.129). They arrive at an intrinsic spectrum with index 1.9 and find that the ``upturn'' in the spectrum around a few TeV may be explained by a decrease of the slope of the $\\gamma$-ray absorption as a function of energy. Our data is not inconsistent with the absorbed spectrum derived by HEGRA. In fact, our lowest point at 370 GeV agrees very well with the hypothesis of the absorbed spectrum, but not with the extrapolation of the alternatively fitted power law (dotted line in Figure \\ref{fig-finalspec}). The latter can be excluded with a confidence level $> 99.5$\\,\\% which argues for a flattening of the spectrum at a few TeV. More data is needed to confirm this result.  Between 250 GeV and 800 GeV, our data clearly favours a spectrum which is steeper than that of any other known TeV blazar at these energies. Since H1426 is also the most distant known TeV blazar, the steepness of the spectrum may be interpreted as evidence for gamma-ray absorption in the intergalactic medium. However, the sample of TeV blazars still needs to be enlarged and the intrinsic spectrum well separated from absorption effects before one can come to firm conclusions."
    },
    "0207/astro-ph0207260_arXiv.txt": {
        "abstract": "This paper presents a new code for performing multidimensional radiation hydrodynamic (RHD) simulations on parallel computers involving anisotropic radiation fields and nonequilibrium effects.  The radiation evolution modules described here encapsulate the physics provided by the serial algorithm of~\\citet{sto92c}, but add new functionality both with regard to physics and numerics. In detailing our method, we have documented both the analytic and discrete forms of the radiation moment solution and the variable tensor Eddington factor (VTEF) closure term.  We have described three different methods for computing a short-characteristic formal solution to the transfer equation, from which our VTEF closure term is derived.  Two of these techniques include time dependence, a primary physics enhancement of the method not present in the Stone algorithm. An additional physics modification is the adoption of a matter-radiation coupling scheme which is particularly robust for nonequilibrium problems and which also reduces the operations cost of our radiation moment solution. Two key numerical components of our implementation are highlighted: the biconjugate gradient linear system solver, written for general use on massively parallel computers, and our techniques for parallelizing both the radiation moment solution and the transfer solution.  Additionally, we present a suite of test problems with a much broader scope than that covered in the Stone work; new tests include nonequilibrium Marshak waves, two dimensional ``shadow'' tests showing the one-sided illumination of an opaque cloud, and full RHD+VTEF calculations of radiating shocks.  We use the results of these tests to assess the virtues and vices of the method as currently implemented, and we identify a key area in which the method may be improved.  We conclude that radiation moment solutions closed with variable tensor Eddington factors show a dramatic qualitative improvement over results obtained with flux-limited diffusion, and further that this approach has a bright future in the context of parallel RHD simulations in astrophysics. ",
        "introduction": "This paper is a logical successor to Paper III of the~\\ztwd~series published in 1992 (see \\citet{sto92a}, \\citet{sto92b}, and \\citet{sto92c}), which describe numerical methods for carrying out radiation magnetohydrodynamic (RMHD) simulations in two dimensions. In the decade which has passed since these papers appeared, both the maximum floating point operation (FLOP) speed and available disk storage capacity have increased by three orders of magnitude: from gigaflops (GFLOP) to teraflops (TFLOP) in speed; from Gbytes to Tbytes in storage.  The increase in computing speed has risen largely from the emergence of massively parallel computer architectures as the high-performance computing paradigm.  Issues of cache memory management on modern RISC-based chips, along with the necessity of writing code for parallel execution, place demands upon application codes that were largely unknown to the bulk of the astrophysics community ten years ago.  The spectacular increase in computing power has spawned a new generation of applications featuring improvements in three major areas: higher dimensionality, higher resolution through larger (or sometimes adaptive) grids, and more realistic physics. Early universe simulations and studies of core-collapse supernovae are but two areas in which increased computing power have profoundly advanced the realism achievable in a numerical simulation; the lessons learned in the development of application codes for such problems are widely applicable to problems throughout astrophysics and engineering.  As in Paper III of the~\\ztwd~series, the focus of this paper is radiation transport and radiation hydrodynamics; in particular, we consider methods which in principle reproduce and preserve large angular anisotropies in the radiation field and which treat time dependence in the radiation field in an appropriate way.  Time dependence and angular anisotropy highlight two great shortcomings of traditional flux-limited diffusion (FLD) techniques; the ways in which our method improves upon the results of FLD form the defining theme of this paper (see Mihalas and Mihalas (1984) for a good discussion of FLD). The context of our paper is broader than that of Paper III, however, in that we have developed new algorithms for simulations on parallel computing platforms, and we identify key issues which must be addressed for the successful implementation of a parallel radiation hydrodynamics (RHD) code.  Additionally, we present a much more extensive suite of test problems than that provided in Paper III; of particular interest are the ``shadowing'' tests which, perhaps more than any other, highlight the qualitative differences between our approach and FLD.  The impetus for this project was provided by a contract, funded by the Lawrence Livermore National Laboratory, which supported two of us (Hayes and Norman) to develop radiative transfer techniques capable of treating extreme spatial and angular anisotropies in the radiation field within a medium in which both light-crossing timescales and (far longer) thermal timescales are important.  The test problem specified for benchmarking a new algorithm was the ``tophat'' (or ``crooked pipe'') test, a description of which is given by Gentile (2001).  The algorithm desired was one that could capture the aforementioned features of the problem at a fraction of the cost of more elaborate Boltzmann (e.g. S$_{n}$) or Monte Carlo methods.  We felt that a moment-based approach like that described in Paper III was an ideal candidate for treating the tophat test, and further that the original serial method could be adapted for parallel use.  The final product of this project is a new set of numerical routines for performing RHD simulations in a parallel environment.  These routines provide all of the abilities advertised for the serial routines in Paper III, and they add new functionality with regard to both physics and numerics.  In addition, these routines have been implemented within~\\zmp, the latest generation of the ZEUS code series. The initials ``MP'' refer to the ``multipurpose,'' ``multi-physics'' (HD, RHD, MHD, gravity, chemistry), and ``massively parallel'' aspects of the code design. The basic HD and MHD equations solved in~\\zmp~are identical to those documented in Paper I and Paper II of the ZEUS trilogy.  The RHD equations in~\\zmp~differ somewhat from those given in Paper III and are documented extensively in this paper. The rewriting of the ZEUS algorithm for parallel execution, with attention given to issues of cache optimization and scalability, has been documented in a refereed conference proceedings available on the World Wide Web \\citep{fied97}.  This paper is organized in the following manner:~\\S{\\ref{moment}} presents the analytic and discrete forms of the RHD moment equations solved in~\\zmp.  \\S{\\ref{transfer}} presents three different algorithms for computing the variable tensor Eddington factor (VTEF) used to close the moment equations.  Two of these algorithms include time dependence in an approximate way, in contrast to the strictly time-independent algorithm of Paper III.  \\S{\\ref{linear}} briefly describes the new linear solver we have written to solve our discrete linear systems on parallel processors, and~\\S{\\ref{parallel}} describes the main issues bearing on the implementation of the moment solution and transfer solution algorithms in a parallel environment. \\S{\\ref{tests}} provides a suite of test problems which exercise all the components of the RHD module.  The main body of the paper concludes with a summary and discussion (\\S{\\ref{discuss}}); a full listing of the linear system matrix comprising our discrete solution to the radiation moment equations is given in the Appendix. ",
        "conclusions": "\\label{discuss}  In this paper, we have presented a new code for performing RHD simulations on parallel computers.  The algorithms discussed here include and augment the functionality of the serial algorithms documented in~\\citet{sto92c}. We have documented both the analytic and discrete forms of the radiation moment solution and the Eddington tensor closure term.  We have described three different methods for computing a short-characteristic formal solution to the transfer equation, from which our VTEF closure term is derived.  Two of these techniques include time dependence, a feature not typically associated with the formal solution in astrophysics literature.  Two key numerical components of our implementation have been highlighted: the biconjugate gradient linear system solver, written for general use on massively parallel computers, and our techniques for parallelizing both the radiation moment solution and the transfer solution.  Finally, we have presented a suite of test problems which run the gamut from optically thin transport with a fixed Eddington tensor, to full RHD+VTEF calculations of radiating shocks.  This document and our code possess a number of features which are new with respect to Paper III and the serial code described therein. In the moment equations, we chose to accompany the radiation energy equation with an equation for the gas energy rather than adopting a total energy equation.  This choice allows us to employ a different matter-radiation coupling scheme with three very attractive features: (1) it is particularly robust in regimes where matter and radiation are far out of equilibrium, (2) it allows the construction of a moment solution matrix involving only one dependent field variable, and (3) it is extremely well-suited to implementation on parallel platforms. In the transfer solution, we have retained the original static method, suitable for computing Eddington tensors for static or quasistatic radiation fields, but we have added two algorithms which include time dependence in different ways.  One of these treats temporal effects by time-retarding the opacities and source functions encountered along a characteristic ray.  The second of these, which we have adopted as our default method for problems needing a time-dependent treatment, builds a discrete form of the time derivative operator directly into the characteristic solution. This ``pseudostatic'' form has an advantage over the ``time-retarded'' form in that it is not necessary to save opacities and emissivities from a previous timestep.  All three algorithms may be used to compute Eddington tensors on the full moment solution grid, or upon a subset of mesh points obtained by subsampling the moment grid at regular intervals along both coordinates.  Design for parallel execution is an entirely new feature with respect to the older code.  We have used the Message-Passing Interface (MPI) standard exclusively in constructing our code for parallel use. In the gas dynamical and radiation moment solution modules, we have distributed work among parallel processors via spatial domain decomposition.  In contrast, the transfer solution has been subdivided over the $\\mu$ angle coordinate, owing to the spatially recursive nature of the SC solution.  This work also contains a more diverse set of test problems than that provided in Paper III.  We feel that this is particularly important, because the RHD modules in Paper III have remained unused since they were written more than ten years ago, due in part to the subsequent research pursuits of their creator, and also to the fact that the publicly available version of~\\ztwd~contained an FLD module rather than the VTEF algorithm described in Paper III. As a result, the viability of this VTEF approach for multidimensional calculations has not been properly documented in the astrophysical literature.  As a first step toward correcting this shortfall, we have provided full RHD tests to which the serial algorithm was never subjected.  In addition, we have focused considerable attention on multidimensional radiation fields containing a shadowed region, consistent both the original mission of this LLNL-funded project and with anticipated astrophysics applications to come.  The test problems we have examined allow us to consider the strengths and weaknesses of the VTEF method for multi-dimensional problems. The major weaknesses of the method, as currently implemented, are the diffusivity of the discrete moment equations and limitations upon the extent to which the transport solution can track rapid time variation of the radiation field.  The diffusivity of the moment solution is a direct consequence of the difference stencil used to discretize the moment equations; the ``double divergence'' term in the radiation energy equation (see equations \\ref{rade3} and \\ref{g1term}) results in finite differences which artificially connect regions in space on either side of physical boundaries in the radiation field, such as the shadow edges in our illuminated cloud problem.  In our test case, this prevents the VTEF algorithm from reproducing such shadowed regions as fully as we would like.  Nonetheless, a noticeable shadow is maintained for very long timescales, with leakage of energy (via the R-component of the flux) into the shadowed region balanced by transmission (via the Z-component of the flux) of energy through the exit boundary.  The second problem affects the degree of consistency between the moment solution and transport solution for the radiation field. Mathematically, our radiation moment equation results from a zeroth angular moment of the time-dependent equation of transfer; using a VTEF closure for the moment equation is therefore accurate only if the radiation solutions as characterized by these two equations are mathematically consistent.  An equivalent statement is that the radiation energy density produced from the moment equation should equal that obtained through a direct integration of the specific intensity over all solid angles.  The first test problem in \\S\\ref{tests} shows that our moment solution is capable of following a sharply defined radiation front when the Eddington tensor is an idealized constant with $\\f_{11}$ = 1.  Our formal solution, however, includes time dependence in an approximate manner, and disagrees strongly with the moment solution when the radiation field shows a rapid variation in time (and therefore space).  We illustrate this behavior with comparing two streaming problems computed with the full VTEF method.  The physical characteristics of the two cases are identical and are taken from the cloud problem, albeit with the dense cloud removed in favor of a uniform medium. As with the cloud test, the illuminating source intensity increases with a prescribed e-folding time.  In the case shown in figure ~\\ref{fig17}, the e-folding time is equal to the light-crossing time along the Z-axis; in figure~\\ref{fig18} this value has been decreased by an order of magnitude.  In both graphs we have plotted the radiation temperature derived from the moment solution for the radiation energy density, and compared it to that derived from an angular integration of the specific intensity.  When the radiation source is illuminated slowly, we see close agreement at all times between the two solutions. When the radiation source is illuminated rapidly, there is a clear disagreement between the two.  In particular, the profiles computed from the transport solver begin to ``lead,'' to a significant degree, those from the moment solution as the wavefront progresses. The significance of this becomes apparent when we realize that the Eddington tensor deviates from isotropy everywhere that the specific intensity has risen from its ambient value.  Once the Eddington tensor changes from its isotropic value, the divergence terms in the flux equation become non-zero and generate flux.  This means that fluxes from the moment solution at a given point in space will become non-zero even before the radiation front has overrun that point, clearly a non-physical result.  As applied to the cloud problem, the consequences of these two shortcomings are: (1) flux in the R-direction will leak radiation into the shadowed region artificially, and (2) this behavior will occur even before the radiation front has swept over the cloud, if the radiation source is illuminated too rapidly.  As a practical matter, this test problem needed to be designed with some care; attempting to illuminate the source too quickly led to unrecoverable errors in the radiation field evolution.  While there are phenomena in astrophysics in which light-travel times may matter (e.g. the emergence of radiation fronts from optically thick surfaces), problems abound in which the radiation field may be treated as time independent, or at worst quasistatic.  In this instance, our equations show good internal consistency, and the dangers of pathological behavior are greatly reduced.  Further, we note that considerable overall improvement to the method will be obtained if our discrete solution is redesigned to eliminate (or greatly reduce) the spurious generation of fluxes at interface boundaries.  This problem is not unique to our method, and alternative schemes for tensor diffusion problems have been published in the radiation transport literature~\\citep{mor98,mor01}. With the support of the Department of Energy, we are now undertaking to implement our VTEF equations with a higher-order spatial scheme which satisfies the appropriate flux constraints.  Results of this effort will be reported in future work.  The strengths of our method, in our view, are threefold.  With regard to angular variations in the radiation intensity, the difference between the VTEF and FLD solutions is fundamental and dramatic.  While the VTEF solution shows room for quantitative improvement, the qualitative behavior is physically sensible, while that of the FLD solution is not.  This result alone compels us to pursue the VTEF approach further, with an eye toward relieving the diffusivity issues identified above.  This algorithm is also particularly suited to environments where matter and radiation are out of equilibrium, a situation often encountered in strongly dynamic environments.  The numerical scheme used to implement matter radiation coupling has the further advantage of allowing an implicit linear system solve for the radiation energy density alone, in contrast to the two-variable coupled system documented in Paper III. Finally, we note that our method is suited for large problems distributed across parallel architectures.  This statement is even truer for problems which require only a static transfer solution for Eddington tensors (owing to a far lower memory requirement), or, in the most favorable case, problems in which static Eddington tensors can be computed analytically.  In this instance, the brunt of the computational burden lies in the linear system solve for $E_{rad}$, and is thus on the same order of the cost for an analogous solution for FLD.  We therefore view this method as a topic worthy of continued study, with regard to both algorithm development and astrophysical applications."
    },
    "0207/astro-ph0207326_arXiv.txt": {
        "abstract": "We analyze STIS spectra in the 1150-1700~\\AA\\/ wavelength range obtained for six early B supergiants in the neighboring galaxy M31. Because of their likely high (nearly solar) abundance, these stars were originally chosen to be directly comparable to their Galactic counterparts, and represent a much-needed addition to our current sample of B-type supergiants, in our efforts to study the dependence of the Wind Momentum-Luminosity Relationship on spectral type and metallicity.  As a first step to determine wind momenta we fit the P-Cygni profiles of the resonance lines of N\\V, Si\\IV\\/ and C\\IV\\/ with standard methods, and derive terminal velocities for all of the STIS targets. From these lines we also derive ionic stellar wind column densities.  Our results are compared with those obtained previously in Galactic supergiants, and confirm earlier claims of `normal' wind line intensities and terminal velocities in M31.  For half of the sample we find evidence for an enhanced maximum turbulent velocity when compared to Galactic counterparts. ",
        "introduction": "The analysis of mass-loss and stellar winds in early-type supergiants, important {\\em per se} because of the physical insight it provides about the atmospheres of massive stars, has gained momentum in recent years with the realization that it can also provide the basis for the determination of stellar distances, through the dependence of the wind momentum on luminosity. With the work of \\citet{puls96} and \\citet{kudritzki99}, the Wind Momentum-Luminosity Relationship (WLR) has been established for OBA supergiants in our own Galaxy and in the Magellanic Clouds. The latter paper showed the impact of the differing stellar atmospheric parameters between spectral types on the parameterization of the WLR. The empirical verification of the predicted dependence of the global properties of stellar winds, and consequently of the WLR, on metallicity has however just begun, because of the considerable observational efforts required (\\citealt{mccarthy95,mccarthy97}). It is nevertheless reassuring that preliminary results on the massive stellar winds in galaxies even beyond the Local Group agree with the present calibration of the WLR (\\citealt{bresolin01,bresolin02}).  As a next step in the understanding and calibration of the WLR we are carrying out several programs, including the study of OB supergiants of known distances in the Galaxy (\\citealt{herrero01}), M33 (\\citealt{urbaneja02}) and M31 (this work). This allows us to investigate stellar winds at different metallicities and spectral types, by removing the uncertainty in the distances which afflicts most of the Galactic work (\\citealt{kudritzki99}).  In this paper we present our first results on early B (B0-B3) supergiants in the neighboring galaxy M31, obtained from Space Telescope Imaging Spectrograph (STIS) ultraviolet spectra. Given the nearly solar abundance of the observed fields, as inferred from \\hii\\/ region studies (\\citealt{blair82}), these stars were originally chosen to be directly comparable to their Galactic counterparts.  A comparison of the WLR obtained for two A-type supergiants in M31 with similar objects in the Galaxy has been presented in the work of \\citet{mccarthy97}, and by \\citet{kudritzki99}. However, for the purpose of calibrating the WLR, early B supergiants are currently under-represented in the available samples of stars with accurate distances. Moreover, as shown by \\citet{kudritzki99}, changes in the ionization stages of the metals driving the wind around the late O and early B types modify the WLR considerably, and must be thoroughly tracked by gathering more data for stars in this spectral range.  For A-type supergiants the wind parameters, including the mass-loss rate and the terminal velocity $v_\\infty$, can be derived from fits of the H$\\alpha$ line to hydrodynamical models of the expanding atmosphere. However, for OB supergiants, profile fits to the UV P-Cygni resonance lines are required to obtain accurate terminal velocities (\\citealt{kud00}), hence the necessity for HST spectroscopy.  Individual OB supergiants in M31 have been observed previously in the vacuum ultraviolet with the IUE satellite (\\citealt{bianchi91}, and references therein) and with the Faint Object Spectrograph (FOS) on the Hubble Space Telescope (HST; \\citealt{hutchings92}, \\citealt{bianchi94}). Some of this early work suffered from low spectral resolution and sensitivity, and suggested that the stellar winds of OB supergiants in M31 could be one order of magnitude weaker than in the Galaxy. This result was challenged by \\citet{herrero94}, who concluded from a study of H$\\alpha$ spectra in M31 AB supergiants that their winds should be comparable to those of similar Galactic stars. A reanalysis of FOS data by \\citet{bianchi94} and \\citet{haser95b} established that mass-loss rates in the studied M31 OB supergiants are comparable to those found in some of their Galactic counterparts.  The targets for the present investigation are six early B supergiants chosen among the brightest from the surveys of massive stars in M31 by \\citet[CCD {\\em UBV} imaging]{massey86}, \\citet{humphreys90} and \\citet[optical spectroscopy]{massey95}.  The nomenclature adopted here for the parent OB associations (see Table~1) is the one introduced by \\citet{vandenbergh64}. One of the supergiants lies in Baade's Field IV (\\citealt{baade63}), within the boundaries of van den Bergh's association OB\\,184, one of the outermost associations in the disk of M31, $96\\arcmin$ (22~Kpc) southwest of the nucleus. This star is indicated as IV-B59 in Table~1, and is one of the three early-type supergiants included in the study of \\citet{humphreys79}. Of the remaining stars, two are found in OB\\,78 (=\\,NGC~206), and one in OB\\,48, both among the richest OB associations in M31, located in the most actively star forming annular region, between 9 and 15~Kpc, identified by \\citet{vandenbergh64}. The last two objects are part of OB\\,8 and OB\\,10, at the smallest galactocentric distance ($\\sim$\\,6~Kpc) in the sample considered here.  High resolution optical spectra of all the STIS targets have been obtained by us with the William Herschel Telescope (WHT) on La Palma and with the Keck~I telescope.  The analysis of the WHT spectra are presented in a companion paper (\\citealt{trundle02}), deriving stellar parameters and photospheric abundances. A future analysis of the complete data set will be achieved in a unified, consistent approach to further yield mass-loss rates and a full test of the WLR for the M31 galaxy. We also note that previous optical spectroscopic work on individual B- and A-type M31 supergiants has been carried out by \\citet[the star OB\\,78-478 is in common with our present sample]{herrero94}, \\citet{mccarthy97}, \\citet{venn00} and \\citet[star OB\\,10-64 in common with current work]{smartt01}.  In this paper we concentrate on the measurement of the terminal velocities from the analysis of the main resonance lines in the UV spectrum. We describe the STIS observations of the six B supergiants in M31 in \\S~2, and the method used to derive the terminal velocities in \\S~3. We comment on the individual stars in \\S~4, and we briefly discuss our results in \\S~5. ",
        "conclusions": "In our previous work on the UV spectra of B supergiants in the Galactic Cyg OB2 association (\\citealt{herrero01}) we were able to use some of the stellar parameters derived from optical spectra, together with the known Galactic WLR (\\citealt{puls96}) and an assumed set of abundances, to estimate mass-loss rates and mean ionization fractions.  Given the uncertainties in abundances and ionization fractions for the current sample of B supergiants in M31, we defer a more complete analysis to a future paper, where we will analyze WHT and Keck optical spectra in order to measure the stellar parameters, and in addition the mass-loss rates from fits to the H$\\alpha$ line profiles. When combined with the terminal velocities obtained in this work, we will be able to derive a WLR for early B supergiants in M31. In the following we will make some qualitative comparisons with Galactic supergiants, and draw some preliminary conclusions.  The abundance gradient in M31 has been characterized by relatively few studies of \\hii\\/ region emission lines (e.g.,~\\citealt{dennefeld81}; \\citealt{blair82}; \\citealt{galarza99}). The large angular extent of the galaxy and the general faintness of the \\hii\\/ regions has precluded more extensive and complete spectroscopic surveys of the gaseous abundances. Moreover, the low excitation of the nebulae implies that the auroral lines necessary to determine the electron temperature remain undetected, with the consequence that empirical strong line methods must be used to estimate the abundances. A recent reanalysis of the available \\hii\\/ region abundances, based on the empirical abundance calibration of \\citet{mcgaugh91}, has been presented by \\citet{smartt01}. The resulting oxygen abundance gradient would imply a range between $12+\\log (O/H)=8.9$ (at the radial distance of OB\\,8-17 and OB\\,10-64) and 8.6 (IV-B59) for the current sample of B-type supergiants, i.e., between a roughly solar abundance and a 0.3 dex lower value.  On the other hand, the scant data available from direct stellar abundance determinations (\\citealt{smartt01}, \\citealt{venn00}) is consistent with a flat gradient between 5 and 20 Kpc, with a value close to that found in the Galactic solar neighborhood ($12+\\log (O/H)=8.7\\pm0.1$, from the mean NLTE value of \\citealt{rolleston00}). The two methods (nebular vs.~stellar) of abundance determination might not be in disagreement, given the large scatter in the derived nebular $O/H$ ratio at a given radial position, and the small number of stellar data available. However, the result obtained for OB\\,10-64 by \\citet[$12+\\log (O/H)=8.7$]{smartt01} is well below the value obtained from \\hii\\/ regions at a similar galactocentric distance.  A detailed abundance study of the B supergiants observed with HST is beyond the scope of the present paper, and is addressed in the companion paper by \\citet{trundle02}. Therefore here we directly compare the UV spectra of B-type supergiants in M31 with stars of similar spectral type in the Galaxy.  Any large discrepancies would immediately reveal differences in the stellar abundances and/or the wind properties. In Fig.~\\ref{comp3} the spectra of two B0\\,Ia supergiants (OB\\,10-64 and OB\\,48-358) are shown superimposed on the IUE spectrum of the Galactic O9.7 Iab star HD~149038. The latter was retrieved from the IUE archive at STScI and rectified in the same way as done for the M31 stars.  There is an overall excellent agreement between the three spectra. In particular, we note a good match in the spectral regions most affected by metal (mostly iron) lines, between 1410 and 1500~\\AA, and at wavelengths redwards of C\\IV\\w1550. The Si\\IV\\w1400 and C\\IV\\w1550 wind line profiles are also very similar in the M31 and Galactic supergiants, with nearly the same terminal velocities and peak intensities. A second example is illustrated in Fig.~\\ref{comp5}, where the UV spectra of IV-B59 and $\\kappa$~Cru (B3\\,Ia) are compared. In this case a slight underabundance is suggested in the spectrum of the M31 star, which would be in agreement with its large galactocentric distance (22~Kpc). The WHT spectra of IV-B59 are not of high enough quality for the accurate measurement of metal lines, hence an abundance analysis of this star is not presented in \\citet{trundle02}. The \\hii\\/ region BA\\,500 (\\citealt{baadearp64}) lies only $10\\arcsec$ east of IV-B59, and its emission line fluxes were measured by \\citet{dennefeld81}. Applying to their measurements the analytical expression given by \\citet{kobulnicky99} for the determination of empirical abundances, based on the strength of the [O\\,II]\\w3727 and [O\\,III]\\ww4959,5007 lines, and on the \\citet{mcgaugh91} photoionization models, we obtain $12+\\log (O/H)=8.53$, which is approximately 2/3 of the oxygen abundance in the solar neighborhood. However, the F5 supergiant A-207, also located in Baade's Field IV, has been found to have a roughly solar abundance by \\citet{venn00}. More detailed abundance studies of individual BA supergiants in M31 are clearly needed to clarify the issue of the abundance gradient in this galaxy.  Similar qualitative comparisons have been shown by \\citet[see also \\citealt{bianchi94}, \\citealt{haser95b}]{bianchi96}, who compared the UV spectral appearance of M31 and M33 early B supergiants with Galactic and LMC ones. The effects of the lower metal abundance were clearly seen in the M33 stars (see also \\citealt{urbaneja02}), with weaker absorption in the Si\\IV\\/ and C\\IV\\/ lines, and in the metal lines in general, in a similar fashion to what is observed in the LMC stars. On the other hand, in the case of M31 their results are comparable to those shown in Fig.~\\ref{comp3}-\\ref{comp5}.  Our data (see Table~2), together with previous reliable determinations of terminal velocities in M31 supergiants (\\citealt{bianchi94}; \\citealt{haser95b}; \\citealt{bianchi96}), support the idea that the winds in M31 stars are comparable to those observed in the Galaxy. This is also supported by the H$\\alpha$ spectra studied by \\citet{herrero94} and by the differential analysis of the optical spectra by \\citet{trundle02}, which indicates that four of the stars have very similar abundances to those derived in B-type supergiants in our Galaxy within 1-2 Kpc of the Sun's position. The oxygen abundances derived are listed for reference in column 3 of Table~2.  The absolute value for OB\\,8-17 appears significantly lower than the others at 8.4 dex.  However, a differential abundance analysis with two solar neighbourhood supergiants indicates that OB\\,8-17 has a very similar oxygen abundance compared to these Milky Way stars. This is also supported by a comparison of the UV spectrum of this object with that of its Galactic counterpart, HD~148688.  Although in one instance the measured $v_\\infty$ is well below the mean Galactic value for the given spectral type, all available data lie within the range observed for Galactic stars, and are therefore consistent with the expected scatter in M31. This is summarized in Fig.~\\ref{vinf}, where the Galactic data have been taken from the compilations of terminal velocities of O and B supergiants by \\citet{haser95} and \\citet{howarth97}. A similar conclusion was reached by \\citet{bianchi96}, although their terminal velocities are in general less accurate then the ones measured with STIS, because of the lower resolution of some of their data.  Our next step will be to add the information from the high-resolution optical spectra, in order to derive wind momenta for the supergiants studied here. This will allow us to provide better constraints on the WLR for early B supergiants of solar abundance."
    },
    "0207/astro-ph0207110_arXiv.txt": {
        "abstract": "Gas and stars in spiral galaxies are modelled with the {\\sc dual} code, using hydrodynamic and N-body techniques. The simulations reveal morphological differences mirroring the dual morphologies seen in $B$ and $K'$ band observations of many spiral galaxies. In particular, the gaseous images are more flocculent with lower pitch angles than the stellar images, and the stellar arm-interarm contrast correlates with the degree of morphological decoupling. ",
        "introduction": "There is much more to spiral galaxies than is apparent from simple observations in optical light. Optical observations are dominated by young stars and gas which may constitute only 5\\% of a galaxy's mass. The majority of the baryonic mass of a galaxy is only revealed in infrared light. Despite the fact that young and old stars orbit in the same potential, observations in the optical and near infrared can reveal radical differences in morphology. Block \\& Puerari (1999) show that $B$ and $K'$ band images of the same galaxy can be completely decoupled and that $K'$ band images are mostly one or two armed whereas $B$ band images are frequently multi armed. Grosbol \\& Patsis (1998) find that grand design spirals are common in $K'$ band images, but in the $B$ band, most galaxies are flocculent with tighter spiral pitch angles of upto $7\\deg$.  The {\\sc dual} code, combining hydrodynamic and N-body techniques, has been used to recreate these dual morphologies, and ask if predictions can be made about intrinsic characteristics of spiral galaxies from their morphologies. High density regions of hydrodynamic simulations are assumed to reproduce the morphology of high luminosity regions in $B$ band images, and N-body methods to trace the underlying mass distribution of the galaxy seen in $K'$ band images. ",
        "conclusions": "Simple simulations have reproduced some of the differences in spiral galaxy morphology in the optical and infrared. These differences can be explained by the dynamics of stars and gas, without invoking interactions, star formation or dust obscuration. In particular, gas images are more flocculent with lower pitch angles than stellar images. Stellar arm-interarm contrast and $Q$ correlate inversely with the degree of decoupling. The effect of gas self-gravity is not known, but it is needed to model the full range of galaxies self-consistently."
    },
    "0207/astro-ph0207332_arXiv.txt": {
        "abstract": "The possibility to unambiguously determine the equation-of-state of the cosmic dark energy with existing and future supernovae data is investigated. We consider four evolution laws for this equation-of-state corresponding to four quintessential models, i.e. i) a cosmological constant, ii) a general barotropic fluid, iii) a perfect fluid with a linear equation-of-state and iv) a more physical model based on a pseudo-Nambu-Goldstone boson field. We explicitly show the degeneracies present not only within each model but also between the different models\\,: they are caused by the multi-integral relation between the equation-of-state of dark energy and the luminosity distance. Present supernova observations are analysed using a standard $\\chi^2$ method and the minimal $\\chi^2$ values obtained for each model are compared. We confirm the difficulty to discriminate between these models using present SNeIa data only. By means of simulations, we then show that future SNAP observations will not remove all the degeneracies. For example, wrong estimations of $\\Omega_m$ with a good value of $\\chi^2_{min}$ could be found if the right cosmological model is not used to fit the data. We finally give some probabilities to obtain unambiguous results, free from degeneracies. In particular, the probability to confuse a cosmological constant with a true barotropic fluid with an equation-of-state different from $-1$ is shown to be $95\\,\\%$ at a $2\\,\\sigma$ level. ",
        "introduction": "Present supernovae data strongly support cosmological models containing a perfect fluid with a negative pressure (Riess et~al. 1998, Perlmutter et al. 1999). The oldest and most studied candidate for this fluid is the cosmological constant, which acts like a perfect fluid whose equation-of-state (hereafter, EOS) is $w \\equiv p\\,/\\,\\rho = -1$, where $p$ is the fluid pressure and $\\rho$ its density. But the vacuum energy density associated with the cosmological constant is 60-120 orders of magnitude smaller than its natural value derived from quantum field theories. This discrepancy is known as the cosmological constant problem (Abbott 1988, Weinberg 1989, Carroll et~al. 1992, Sahni \\& Starobinsky 2000) and has led theorists to find alternative dark energy candidates with a present negative pressure. The simplest models are based on a generalisation of the cosmological constant, e.g. a barotropic fluid ($w(z) = w_0 =$ constant $< 1$\\,; Gonz\\'alez-D\\'\\i az 2000, Di Pietro \\& Demaret 2001) or a homogeneous fluid for which $w$ is a linear function of $z$ ($w(z) = w_0 + w_1\\,z$ with $w_0$ and $w_1$ constant\\,; Goliath et~al. 2001, Maor et al. 2001, 2002).  These models are interesting because they are described by simple field equations. However, they suffer from a lack of physical justification. Other models with a more general EOS and stronger physical interpretation have been proposed (Peebles \\& Ratra 1988, Ratra \\& Peebles 1988, Wetterich 1988, Ferreira \\& Joyce 1998, Steinhardt et~al. 1999). The most popular ones are based on a dynamical quintessential component represented by a minimally coupled scalar field evolving in a potential. Mathematically, such a quintessence fluid can be described by an EOS $w$ function of the redshift $z$ ($-1 \\leq w(z) \\leq 1$). Among these quintessence models, we shall consider the one which assumes the existence of an ultra-light pseudo-Nambu Goldstone boson (PNGB) field relaxing to its vacuum state (Frieman \\& Waga 1998, Waga \\& Frieman 2000, Ng \\& Wiltshire 2001a). From the quantum viewpoint, the PNGB models are the simplest way to introduce a naturally ultra-light scalar field able to reproduce the cosmological observations (Frieman et~al. 1992, 1995). Moreover PNGB models provide an interesting theoretical framework to any spontaneous symmetry breaking which could justify the neutrino mass found in the Mikheyev-Smirnov-Wolfenstein solution to the solar neutrino problem (Wolfenstein 1979, Mikheyev \\& Smirnov 1985).  In summary, many theoretical models describing a cosmological fluid with negative pressure have been proposed\\,: some are designed to be mathematically simple while others rely on a more physical justification. The first aim of this paper is to put in evidence the strong degeneracies existing between the luminosity distance predicted by four quintessence models\\,: a cosmological constant, a general barotropic fluid, a fluid with a linear equation-of-state and a more physical model based on a pseudo-Nambu-Goldstone boson field. Second, we shall determine the constraints on those models coming from the present observations of type Ia supernovae (hereafter, SNeIa) and confirm the need for more data to discriminate between them. This may be fulfilled by the proposed SNAP satellite (SuperNova/Acceleration Probe\\,; see SNAP URL). The main objective of this instrument is to detect a very large number of supernovae up to a redshift of $1.7$, in order to yield a more precise determination of the cosmological parameters and therefore to provide information on the nature of dark energy. The third objective of this paper is to analyse, by means of simulated data, how well future SNAP observations alone will be able to break the degeneracies that we put in evidence.  Several authors have already studied the feasibility of SNAP to determine the properties of the dark energy. Depending on the method used for reconstructing the cosmological model, their conclusions are quite different\\,: some authors are optimistic regarding the possible determination of the EOS of the dark energy using SNeIa (Huterer \\& Turner 1999, Nakaruma \\& Chiba 1999, Saini et~al 2000, Chiba \\& Nakaruma 2000, Weller \\& Albrecht 2001, 2002) while others are more cautious (Barger \\& Marfatia 2001, Astier 2001, Maor et~al. 2001, 2002, Gerke \\& Efstathiou 2002). Most of these discrepancies can be traced back to the differences in initial assumptions, prior knowledge, ... used for the reconstruction of the cosmological model (Goliath 2001). Usually, optimistic conclusions result from strong assumptions, such as an accurate knowledge of $\\Omega_m$.  The structure of the paper is as follows. In Section~2, we present the field equations describing the four cosmological models we have considered. In Section 3, we explicitly display the degeneracies between the luminosity distances predicted by these four models. The constraints brought by present SNeIa data on the parameters of each model are presented in Section~4. In Section~5, we explore the ability of SNAP to break the degeneracies and so to discriminate among these models. Finally, we summarize our work and give some conclusions in Section~6. ",
        "conclusions": "\\label{summary} Making use of four quintessence models, of which three are mathematically simple and one is coming directly from particle physics, we have first shown that the various cosmological models may predict apparent magnitudes for objects at given redshifts which differ from each other by less than 0.04 mag till $z = 2$. These magnitude differences are thus small compared to the intrinsic spread of the SNeIa maximum luminosity (0.15 mag). This indicates that an unambiguous discrimination between the cosmological models will be difficult to reach, with cosmological tests only based on luminosity distances.  Second, we have fitted these models to the present supernovae data and found equally good fits with the four models. Finally, we have explored the discriminatory power of future SNeIa data to constrain the dark energy properties. We have simulated a large number of SNAP data with each particular model and have re-analysed them with the four cosmological models. We have then compared the positions and the values of the $\\chi^2_{\\rm min}$ obtained using wrong models with those obtained with the fiducial model. This led us to confirm the more skeptical conclusions already made in some previous works on the subject (Barger and Marfatia 2001, Astier 2001, Maor et~al. 2001, 2002)\\,: some degeneracies between the curvature, the matter density and the equation-of-state of dark energy will be difficult to break at the $3\\,\\sigma$ level and even at the $2\\,\\sigma$ and $1\\,\\sigma$ level. We also found that different estimates of a basic parameter like $\\Omega_m$ can be obtained, depending on the model used for the data processing. Moreover, whatever the true model is, the presently admitted FLRW model with $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$ will not be rejected on the basis of the future SNAP observations alone.  In order to compare results expected from SNAP with the ones which could be obtained by other supernovae surveys, we have also generated data samples from models C and D for another redshift distribution\\,: 2050 SNeIa between $z=0$ and $z=0.5$ and 50 between $z=2$ and $z=2.5$ (with $\\sigma_i$ = 0.15 mag). This distribution reflects future observations to be made with the 4m Liquid Mirror Telescope (see ILMT URL) and with the Next Generation Space Telescope (see NGST URL). No significant difference has been found with the results presented here.  Many recent papers discussed the strong constraints that could be expected from future SNeIa data on cosmological parameters. However, the most important question is\\,: how can we be sure that we consider the right model~? Indeed how useful are strong constraints on cosmological parameters if they do not describe the model chosen by Mother Nature~? The degeneracy problem presented in this paper is expected to affect every cosmological test using luminosity distance. Therefore it will be very difficult for SNAP alone to obtain any strong and safe constraints on the cosmological parameters. However, we expect that the combination of the CMB data with the future SNAP ones could help to break some of the luminosity distance degeneracies. This fact is already known in the framework of a single cosmological model (see e.g. Huterer \\& Turner 2001, Frieman et al. 2002, Bean \\& Melchiorri 2002)."
    },
    "0207/astro-ph0207618_arXiv.txt": {
        "abstract": "{The {\\bllac} Object {\\source}, at a red-shift of $z$$=$$0.129$, has been monitored by the {\\cat} telescope from February 1998 to June 2000. The accumulation of 26 hours of observations shows a {\\gr} signal of 321 events above $250\\:\\mathrm{GeV}$ at 5.2 standard deviations, determined using data analysis cuts adapted to a weak, steep-spectrum source. The source emission has an average flux of ${\\Phi}_\\mathrm{diff}({400\\:\\mathrm{GeV}})=6.73\\pm1.27^\\mathrm{stat}\\pm1.45^\\mathrm{syst}\\times10^{-11}\\:\\mathrm{cm^{-2}\\:s^{-1}TeV^{-1}}$, and a very steep spectrum, with a differential spectral index of $\\gamma$$=$$-3.60$$ \\pm $$0.57$ which can be refined to $\\gamma$$=$$-3.66$$ \\pm $$0.41$ using a higher flux data subset. If, as expected from its broad-band properties, the Very High Energy emission is hard at the source, these observations support a strong absorption effect of {\\grs} by the Intergalactic Infrared field. ",
        "introduction": "{\\bllac} objects, together with Flat Spectrum Radio Quasars (FSRQs), constitute the extreme class of Active Galactic Nuclei known as blazars. They exhibit a highly variable non-thermal spectral energy distribution (SED) which is interpreted as the Doppler boosted emission of a jet pointing towards the Earth.  While at least 66 blazars are reported in the third {\\egret} catalogue (\\cite{Hartman99}), in the TeV range only two nearby {\\bllac} objects, {\\mfr} ($z$$=$$0.031$) and {\\mfv} ($z$$=$$0.034$), have been firmly established so far. This can be explained partially by the small field of view of ground-based telescopes (a few degrees), although, as compared to satellite-borne instruments, they benefit from a much larger sensitivity. The attenuation of \\grs\\ through pair-production with Intergalactic Infrared field (IIR) is another limiting factor for viewing distant sources in the Very High Energy ({\\vhe}) domain. However, the detection of absorption features in blazar spectra can provide interesting constraints on the poorly-measured $0.5-20\\:\\mu\\mathrm{m}$ IR band.    In this context, a survey of nearby {\\bllacs} has been carried out with the {\\cat} telescope since 1997. Among them, {\\source} ($\\alpha_{J2000}$$=$$\\mathrm{14^h28^m32.7^s}$, $\\delta_{J2000}$$=$$+42\\dg40'21''$) at a red-shift of $z=0.129$, occupies a peculiar position: the characteristics of its radio-to-\\xr\\ SED are very similar to those of the two TeV {\\bllacs}, especially its hard \\xr\\ emission which flags the presence of ultra-relativistic electrons. On the other hand, the relatively high red-shift of \\source\\ can imply significant absorption by the IIR photons.    VHE \\gr\\ detections of  {\\source} have recently been reported by Whipple and {\\hegra} during 1999-2001 (\\cite{Horan02}, \\cite{Aharonian02}). In this letter, we report on its detection based on observations made by the \\cat\\ telescope from February 1998 to June 2000. In Sect. 5 we will also give a first estimation of its spectrum. ",
        "conclusions": "Observations of \\source\\ from 1998 to 2000 by \\cat\\ show a signal at $5.2\\sigma$ with an average event rate of order of $\\sim 0.2$ Crab. This is the most distant detected \\bllac\\ in the \\vhe\\ range.  For the three observing seasons studied here the source average emission seems stable, within the measurements' accuracy, while there is some evidence for time variability with transient emission rates $> 0.5$ Crab.  Given its broad-band properties, \\source\\ was one of the most promising \\vhe\\ candidates, despite its red-shift of $>\\:0.1$  (\\cite{Costamante99}). Its detection supports the unifying scheme by \\cite{Ghisellini98}, where the lower luminosity blazars with high frequency peaked synchrotron emission (HBLs) are efficient accelerators to very high energies. It is remarkable that \\xr\\ measurements (\\cite{Costamante01}), contemporaneous to \\cat\\ observations in February 1999 reported a peak emission energy of $\\sim 100\\:\\mathrm{keV}$, comparable to that of {\\mfv} during its strong activity in 1997.  If, as suggested by these data, the same acceleration/cooling mechanisms are at play within the two sources under similar conditions, one would expect a hard \\vhe\\ differential spectrum with $\\gamma \\sim 2.0$ at the source (see \\cite{Djannati99}). The very steep spectrum observed here, $\\gamma=3.66\\pm0.41$, supports a strong absorption of {\\grs} by the diffuse Intergalactic Infrared field, even though the effect is compatible with several --- among many --- estimates and models of its density in the $0.5-20\\:\\mu\\mathrm{m}$ band.  Any definite conclusion on this subject and any measurement of the IIR density will require the detection of more sources at different red-shifts."
    },
    "0207/astro-ph0207042_arXiv.txt": {
        "abstract": "The opacity due to photodissociation of $^{24}$MgH is investigated in the atmospheres of cool stars. The lowest two electronic transitions $A~^2\\Pi\\leftarrow X~^2\\Sigma^+$ and $B'~^2\\Sigma^+\\leftarrow X~^2\\Sigma^+$ are considered where the cross sections for the latter were published previously (Weck, Stancil, \\& Kirby 2002a) while the former are presented in this work. Model atmospheres calculated with the {\\tt PHOENIX} code are used to investigate the effect of the photodissociation opacity on spectra of cool stars. The $A~^2\\Pi\\leftarrow X~^2\\Sigma^+$ photodissociation cross sections are obtained using a combination of {\\it ab initio} and experimentally derived potential curves and dipole transition moments. Partial cross sections have been evaluated over the accessible wavelength range $\\lambda\\lambda~1770-4560$~\\AA~for all rotational transitions from the vibrational levels $v''=0-11$. Assuming a Boltzmann distribution of the rovibrational levels of the $X~^2\\Sigma^+$ state, LTE photodissociation cross sections are presented for temperatures between 1000 and 5000~K. Shape resonances, arising from rotational predissociation of quasi-bound levels of the $A~^2\\Pi$ state near threshold, characterize the LTE photodissociation cross sections. A sum rule is proposed as a check on the accuracy of the photodissociation calculations. ",
        "introduction": "The lack of accurate and complete molecular line and continuum opacity data has been a serious limitation to developing atmospheric models of cool stars (M and later), solar system planets, and Extrasolar Giant Planets (EGPs). Sophisticated modeling programs, such as {\\tt PHOENIX} \\citep{hau99}, require high quality opacity data in order to produce synthetic spectra and predict physical parameters (e.g., surface chemical composition, effective temperature, etc). Typically, atmosphere models include molecular bands with hundreds of millions of spectral lines, mostly derived from molecular band or Hamiltonian models. Moreover, atmospheric models do not consider the effect of molecular photodissociation processes, which may play a role in the opacity at visible and UV wavelengths. An exception is the work of \\citet{kur87} who included LTE photodissociation calculations of CH and OH in solar and cool stellar atmosphere models. Further, it has been suggested by \\citet{sho94,sho96} that a missing opacity source between 2500 and 4000 \\AA~ might be due to MgH, based on the ground rovibrational photodissociation cross section of \\cite{kir79} which peaks at 2920 \\AA~, but has a threshold at 3100 \\AA. However, the LTE $B'~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ MgH calculations of \\citet*{wec02a} find that the cross sections have significant amplitude to nearly 4500 \\AA.  In the present work, we have performed extensive calculations of rovibrationally-resolved photodissociation cross sections of $^{24}$MgH through the $A~^2\\Pi \\leftarrow X~^2\\Sigma^+$ transition, using the most accurate available molecular data. Calculations were performed for the full range of 313 rovibrational levels $(v'',J'')$ in the ground electronic state. Assuming a Boltzmann distribution of the rovibrational levels of the $X~^2\\Sigma^+$ state, LTE photodissociation cross sections are also presented for temperatures between 1000 and 5000~K. These photodissociation cross sections as well as those of \\citet{wec02a} for the $B'~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ transition have been included in cool stellar atmosphere calculations to test if MgH photodissociation is the missing opacity postulated by \\citet{sho94,sho96}. In this paper we investigate the importance of the bound--free opacity in both very cool dwarfs and the objects discussed in \\citet{sho94,sho96} (K type giants and solar type dwarfs).  Atomic units are used throughout to discuss the molecular calculations unless otherwise stated. ",
        "conclusions": "Photodissociation cross sections have been calculated for the $A~^2\\Pi\\leftarrow X~^2\\Sigma^+$ transition of $^{24}$MgH for all rotational transitions from the vibrational levels $v''=0-11$ and over the accessible wavelength range $\\lambda\\lambda~1770-4560$~\\AA. As predicted earlier by \\cite{kir79}, using the Franck-Condon picture, photodissociation cross sections through the $A~^2\\Pi$ are several orders of magnitude smaller than through the $B'~^2\\Sigma^+\\leftarrow X~^2\\Sigma^+$ pathway. However, interesting features such as a large number of shape resonances, arising from rotational predissociation near thresholds, appear in the LTE photodissociation cross sections calculated for temperatures between 1000 and 5000~K using a Boltzmann distribution of the rovibrational levels of the $X~^2\\Sigma^+$ state. Inclusion of the $^{24}$MgH $A\\leftarrow X$ and $B'\\leftarrow X$ photodissociation cross sections in atmosphere models of cool stars results in only minor changes of the computed spectra. Though it does not explain the missing opacity found by Short \\& Lester (1994, 1996), the MgH bound-free opacities will have to be considered for the analysis of high resolution spectroscopy at high S/N. Further, as MgH is a trace molecule, the opacity due to photodissociation may be more significant for molecules with larger abundances such as H$_2$O."
    },
    "0207/astro-ph0207568_arXiv.txt": {
        "abstract": "We compute the structure function scaling of the integrated intensity images of two \\jco\\ maps of Taurus and Perseus. The scaling exponents of the structure functions follow the velocity scaling of super--sonic turbulence, suggesting that turbulence plays an important role in the fragmentation of cold interstellar clouds. The data also allows to verify the validity of the two basic assumptions of the hierarchical symmetry model, originally proposed for the derivation of the velocity structure function scaling. This shows that the same hierarchical symmetry holds for the projected density field of cold interstellar clouds. ",
        "introduction": "Due to the complexity of the Navier--Stokes equations, mathematical work on turbulence is often inspired by experimental and observational measurements. Since geophysical and laboratory flows are predominantly incompressible, turbulence studies have been limited almost entirely to incompressible flows (or to infinitely compressible ones, described by the Burgers equation). Little attention has been paid to highly--compressible, or super--sonic turbulence.  The cold interstellar medium (ISM) of galaxies such as the Milky Way is both highly turbulent and highly super--sonic, within a range of scales from hundreds of parsecs to approximately a tenth of a parsec. Large observational surveys of the cold ISM of our galaxy and of dust within it have become available in the last few years and new surveys of unprecedented sensitivity and resolution will be obtained in the near future by FIR satellites such as SIRTF and Herschel.  The cold ISM provides a good laboratory for super--sonic turbulence and results from new observational surveys should motivate new mathematical work. Furthermore, the interpretation of the astronomical data requires a knowledge of basic properties of super--sonic turbulence.  Previous works have tried to investigate the properties of ISM turbulence by estimating the second order structure function or the power spectrum of the velocity field in dark clouds \\citep{Kleiner+Dickman87,Hobson92,Miesch+Bally94,% Brunt+Heyer2002calibration,Brunt+Heyer2002results}, as sampled by molecular emission lines, with a number of different methods. The centroid velocity at each map position is used as an estimate of the local radial velocity. This centroid velocity results from a convolution of density, velocity and excitation temperature along the line of sight and is not easily related to the three dimensional velocity structure. Other works have instead studied the column density distribution of dark clouds, estimated from the integrated intensity of molecular emission lines or from dust thermal emission. Fractal dimensions have been estimated \\citep{Beech87,Bazell+Desert88,Scalo90,Dickman+90,% Falgarone+91,Zimmermann+92,Henriksen91,Hetem+Lepine93,Vogelaar+Wakker94,% Elmegreen+Falgarone96}. Multifractal \\citep{Chappell+Scalo01} and wavelet \\citep{Langer+93} analysis have also been proposed as a way of characterizing the projected density structure of molecular clouds.  The multifractal analysis applied by \\cite{Chappell+Scalo01} to dust continuum images is related to the scaling of the moments of the projected density. In the present work we study the scaling of integrated intensity differences (structure functions) of the J=1--0 transition of $^{13}$CO from the Taurus and the Perseus molecular cloud complexes. If we had computed the moments of intensity, instead of intensity differences, then or analysis would be equivalent to the multifractal analysis by \\cite{Chappell+Scalo01}.  We have already shown in previous works that super--sonic turbulence in a roughly isothermal gas, such as the cold ISM, generates a complex density field with density contrasts of several orders of magnitude and with statistical properties consistent with observational data from molecular clouds \\citep{Padoan+98cat,Padoan+99per,Padoan+Nordlund99mhd}. We usually refer to this process as {\\it turbulent fragmentation}. Turbulent fragmentation of star forming clouds is unavoidable, since super--sonic turbulent motions are ubiquitously observed. The results of this work provide further evidence of the importance of turbulent fragmentation in star forming clouds.  In \\S~2 we compute the relative scaling of the structure functions and in \\S~3 we discuss our results. The hierarchical structure model is briefly presented in \\S~4 and its validity for our data is verified with the so called $\\beta$ and $\\gamma$ tests. We draw our conclusions in \\S~5. ",
        "conclusions": "We have computed the structure functions of the integrated intensity images of two \\jco\\ maps of Taurus and Perseus. The structure functions scale as power laws within the range of scales 0.3--3~pc. The scaling exponents have been computed up to the 20th order and are statistically significant at least up to the 15th order in Taurus and 8th order in Perseus. They are found to follow very closely the velocity scaling of super--sonic turbulence proposed by \\cite{Boldyrev2002}, within 5\\% and 10\\% accuracy for Taurus and Perseus respectively.  We have verified that the projected density field (or integrated intensity) of the Taurus and Perseus molecular cloud complexes can be described by a hierarchical model as the one proposed by \\cite{She+Leveque94} for the velocity structure functions of incompressible turbulence. We have done so by testing the validity of the two basic assumptions of the hierarchical model for our data. The validity of the assumptions of the hierarchical model means that the integrated intensity images we have analyzed (an approximate estimate of the projected density of the Taurus and Perseus molecular cloud complexes) are the result of a multiplicative process with Log--Poisson statistics \\citep{Dubrulle94}.  The complete derivation of the relation between the structure functions of velocity and projected density is the subject of future work. However, the close similarity of the structure functions of projected density in Taurus and Perseus with that of the velocity field of turbulence provides additional evidence that super--sonic turbulence is the major factor controlling the density field in the range of densities and scales sampled by the maps we have analyzed.  It is well established that super--sonic turbulence plays an important role in the dynamics of the cold ISM \\citep{Larson81,Padoan+98cat,Padoan+99per,Padoan+2001cores,Padoan+Nordlund99mhd} The statistical properties of this ISM turbulence need to be discovered and understood in order to elaborate a statistical theory of star formation \\citep{Padoan+Nordlund2002imf} We have shown in the present paper that this can be achieved with existing observational data, by studying the projected density field of molecular clouds. This type of study will be greatly improved by far infrared imaging of the dust thermal emission from turbulent ISM clouds, obtained by future satellite missions such as SIRTF."
    },
    "0207/gr-qc0207005_arXiv.txt": {
        "abstract": "In this paper we look at the gravitational spin--spin interaction between macroscopic astronomical bodies. In particular, we calculate their post--Newtonian orbital effects of order $\\mathcal{O}(c^{-2})$ on the trajectory of a spinning particle with proper angular momentum ${\\bf s}$ moving in the external gravitomagnetic field generated by a central spinning mass with proper angular momentum ${\\bf J}$. It turns out that, at order $\\mathcal{O}(e)$ in the orbiter's eccentricity, the eccentricity the pericenter and the mean anomaly rates of the moving particle are affected by long--term harmonic effects. If, on one hand, they are undetectable in the Solar System, on the other, maybe that in an astrophysical context like that of the binary millisecond pulsars there will be some hopes of measuring them in the future. ",
        "introduction": "In recent years the topic of the general relativistic post--Newtonian effects of  order $\\mathcal{O}(c^{-2})$ on various features of the motion of spinning particles freely orbiting a central astronomical body has received great attention both theoretically and experimentally. More precisely, when a spinning particle moves in an external gravitational field one has to describe both its spin precession and the influence of the spin on its trajectory$^1$.  The geodetic, or De Sitter, precession$^2$ refers to the coupling of the static, gravitoelectric part of the gravitational field due to the Schwarzschild metric generated by a central, non--rotating object to the spin of a particle freely orbiting around it. It has been measured for the Earth--Moon orbit, thought of as a giant gyroscope, in the gravitational field of the Sun with 1$\\%$ accuracy$^3$.  It should be measured for four superconducting gyroscopes in the gravitational field of the Earth by the important GP-B mission$^4$ at a claimed relative accuracy level of $2\\times 10^{-5}$. Finally, it might be possible to detect it also in some binary pulsar systems$^5$.  The Lense--Thirring drag of the inertial frames$^6$ is an effect due to the stationary, gravitomagnetic part of the gravitational field of a rotating central mass with proper angular momentum ${\\bf J}$ on the geodesic path of a freely falling test particle, i.e. considered to be not spinning itself\\footnote{Such effect could be thought of as a spin--orbit interaction between the spin of the central object and the orbital angular momentum of the test body. For some spin--orbit effects induced by the rotation of the Sun on the orbital angular momentum of the Earth--Moon system see reference$^7$.}. In 1998 the first evidence of the Lense--Thirring effect in the gravitational field of the Earth has been reported$^8$ with a claimed accuracy of almost $20\\%$. It is based on the analysis of the laser--ranging data of the LAGEOS and LAGEOS II geodetic satellites. The launch of the proposed LAGEOS--like LARES satellite$^9$ could allow to measure such effect with an accuracy probably better than $1\\%$. Another interesting gravitomagnetic effect of order $\\mathcal{O}(c^{-2})$ on the orbit of a test particle has been recently derived in reference$^{10}$; it is due to the temporal variability of the Earth's angular momentum. Unfortunately, it is too small to be detected with Satellite Laser Ranging.  The spin of the central object affects also the spin of  a particle freely orbiting it in a way discovered by Schiff$^{11}$ in 1959. The detection of this subtle precessional effect, in addition to the geodetic precession, is one of the most important goals of the GP-B mission$^4$; the claimed accuracy amounts to $1\\%$.  In this paper we are interested in looking for some orbital effects due to the spin--spin gravitational coupling on the geodesic path of a spinning extended particle with mass $m$ and proper angular momentum ${\\bf s}$ freely orbiting around a central body of mass $M$ and proper angular momentum ${\\bf J}$. The gravitational spin--spin coupling in the quantum mechanical domain has been treated in references$^{12,\\ 13}$.  The paper is organized as follows. In Section 2 we derive the gravitational Stern--Gerlach force with some simplifying assumptions. In Section 3 we work out the long--term orbital effects of such interaction on the Keplerian orbital elements of the moving particle. In Section 4  we look at the Sun--Mercury system and to PSR B1259-63 and PSR B1913+16 in order to see if the predicted effect could be measured. Section 5 is devoted to the conclusions. ",
        "conclusions": "In this paper, in the linearized approximation of gravitoelectromagnetism and at order $\\mathcal{O}(c^{-2})$,  we have calculated the influence of the spin ${\\bf s }$ of a particle on its geodesic orbital motion in an external gravitomagnetic field generated by a rotating body with spin ${\\bf J }$. It has been assumed that the orbit lies in the equatorial plane of the central object and that the spins are both vertical to it. The orbital--averaged, long--term effects on the Keplerian orbital elements of the orbiter have been calculated by neglecting all terms of order $\\mathcal{O}(e^{2})$. It turns out that the eccentricity, the pericenter and the mean anomaly are affected by such spin--spin gravitomagnetic interaction by means of harmonic perturbations with half the period of the pericenter. In order to see if they are measurable, we have examined two possible astronomical scenarios: the Sun--Mercury system and a pair of binary millisecond pulsar systems. It turns out that they are far too small. Only for the pulsar PSR1913+16 the predicted periastron rate is not too far from the present experimental sensitivity. The possible discovery of new, highly eccentric and close binary pulsar systems, together with notable improvements of the accuracy in measuring the periastron rates could give some hopes to detect such tiny effects in future."
    },
    "0207/astro-ph0207274_arXiv.txt": {
        "abstract": "{ A 50 x 50 arc sec region near the solar disc center containing a bright point (BP) was observed with the {\\it SUMER}- spectrograph of the {\\it SOHO} observatory. The data consist of two hours observation of four far-UV emission lines formed between 2 10$^4$ -- 6 10$^5$ K, with 2 arc sec spatial, 2.8 min temporal  and 4 km/s  spectral resolution. A striking feature was the strong microflaring of the major persistent BP (with size 8 x 8 arcsec) and the appearance of several short lived transients. The microflaring  of  each individual 2 x 2 arc sec  pixel inside the main BP was coherent, indicating strong interaction of the possible sub arc sec building blocks (magnetic flux tubes). Using  the emission measure at 10$^5$ K  as an indicator of the loop foot point area and magnetic filling factor, we suggest 10 per cent filling factor for the BP observed. This is similar to that on the average  surface of a medium-active solar type  star. ",
        "introduction": "The nature of solar coronal bright points (BP) has been an enigma since their discovery in late 1960's (e.g. Vaiana et al., 1970). Their correspondence with small bipolar magnetic regions was found by combining ground-based magnetic field measurements with simultaneous space-borne X-ray imaging observations (Krieger et al., 1971; Golub et al., 1977). The bright points are clearly seen in Ca K spectroheliograms as bright ``knots'' (e.g. Golub \\& Pasachoff, 1997). The daily number of BP's found on the Sun varies between several hundreds up to a few thousand (Golub et al., 1974). Zhang et al. (2001) found their density around 800 BP's for the entire solar surface at any moment.  Golub et al. (1974) found that the diameters of the bright points in X-rays are around $10^{4}$~km (10 - 20 arc sec) and their lifetimes  range from 2 hours to 2 days (see also Zhang et al. 2001). The physical studies have indicated the temperatures to be fairly low, $T \\approx 2 \\times 10^6 $~K, and the electron densities  $n_e \\approx 5 \\times 10^9$~cm$^{-3}$ (e.g. Golub \\& Pasachoff, 1997), although  cooler BP's exist (Habbal, 1990). Assuming that almost all BP's represent new magnetic flux emerging at the solar surface, their overall contribution to the solar magnetic flux would exceed that of the active regions (Golub and Pasachoff, 1997).  BP's often show irregular intensity and spatial shape variations on time scales of minutes which can be called as microflares (Shimojo \\& Shibata, 1999; Zhang et al. 2001). In particular, Shimojo and Shibata show that the frequency distribution of these microflares, as a function of the peak intensity, show a power law with index 1.7$\\pm{0.4}$ which is consistent with that of ordinary (stronger) flares. However, the total  coronal heating can not energetically be due to these microflares, they  are too few. Brown et al. (2001), using {\\it TRACE}, studied for the first time BP's for their entire lifetime with a cadence of 2 min and spatial resolution of 0.5 arcsec, using hot FeXII and FeXI lines. In particular, they suggested that BP's are made up of a complex system of dense loops.  In the present paper we report observations of one particular bright point observed with the {\\it SUMER}-spectrometer onboard {\\it SOHO}  using four ultraviolet lines formed between 2 10$^4$ - 6 10$^5$ K. ",
        "conclusions": "The main bright point (A in Fig. 2)  had dimensions 8 x 8 arcsec between 2 10$^4$ - 6 10$^5$ K. The non-thermal velocities  (40 - 45 km/s) were somewhat smaller but close to the sonic ones. The most striking feature of the present observation, lasting 120 min, was the strong microflaring  and  the appearance of  short lived transients. One of these   (BP-B in Fig.2) showed clear mass upflows with velocities 20 km/s.   No flows were detected elsewhere in the observed region.  Inside the 8 x 8 arcsec area of the BP-A no clear structure could be resolved within the  2 arcsec spatial resolution. All spatial pixels varied  coherently with strong correlation between individual pixels (Fig. 4). It is   possible that the BP-A consisted of a dense system of small loops  triggering each others, so that the whole BP looked like a single variable object. A much better spatial resolution (like in {\\it TRACE}, Brown et al. 2001) is needed to resolve the issue and to find the BP building blocks.  Comparing  the sun and solar type active stars, Vilhu (1987, 1994) suggested a correlation between the CIV 1550 line intensity (10$^5$ K) and the magnetic filling factor. Lines formed at 10$^5$ K arise close to the loop foot points and are good indicators of their surface area. In the quiet solar network the magnetic filling factor is at one per cent level (one per cent  of the surface covered by equipartition fields) while in the most active (and consequently rapidly rotating) stars the filling factor approaches 100 per cent when the CIV line surface flux  is  100 times that of the quiet sun. This suggests that the filling factor of BP-A was around 10 per cent (see Fig. 5). In this picture  BP-A was similar to the average surface of a moderately active star (like solar type stars in the Hyades cluster, rotating with periods around 10 days or less).  Further, some bright points (like the present one) may well be essentially like over dense quiet sun regions. However, much better spatial resolution will be needed to make any definite conclusion.  We also investigated the underlying magnetic field structure of the region during our observation (from the SOI-archive). The MDI-instrument onboard {\\it SOHO} observed the region  between UT 15:16 - 15:36 and 16:45 - 17:05. In both MDI-observations (around the BP-A of Fig.2)  two bright positive flux regions (10  and 20 arcsec apart) and two negative flux fragments (30 arcsec apart) are present and look rather persistent. According to Brown et al. (2001) one third of the bright points lie over emerging regions of magnetic flux while the remaining two thirds lie above cancelling magnetic features. It seems that the bright point studied in the present paper belongs to this latter category."
    },
    "0207/astro-ph0207512_arXiv.txt": {
        "abstract": "{ We present the complete results of the planet search program carried out at the ESO Coud\\'e Echelle Spectrometer (CES) on La Silla, using the Long Camera from Nov.\\,1992 to April 1998. The CES survey has monitored $37$ late-type (F8V -- M5V) stars in the southern hemisphere for variations in their differential radial velocities ($RV$) in order to detect Doppler reflex motions caused by planetary companions. This led to the discovery of the first extrasolar planet in an Earth-like orbit around the young (ZAMS) and active G0V star $\\iota$ Horologii (K\\\"urster et al. \\cite{martin00}). Here we present the $RV$ results for all survey stars and perform a statistical examination of the whole data-set. Each star is tested for $RV$ variability, $RV$ trends (linear and non-linear) and significant periodic signals. $\\beta$~Hyi and $\\epsilon$~Ind are identified as long-term, low-amplitude $RV$ variables. Furthermore, for $30$ CES survey stars we determine quantitative upper mass-limits for giant planets based on our long-term $RV$ results. We find that the CES Long Camera survey would have detected short-period (``51 Peg-type'') planets around $all$ $30$ stars but no planets with $m\\sin i < 1~{\\rm M}_{\\rm Jup}$ at orbital separations larger than 2 AU. Finally, we demonstrate that the CES planet search can be continued without applying velocity corrections to the $RV$ results coming from the currently installed Very Long Camera at the CES. ",
        "introduction": "The exciting discoveries of giant planets orbiting solar-type stars by precise Doppler searches have caused a shift in our paradigm of the structure and formation of planetary systems. Although we now know that planets have also formed around stars other than the Sun, their orbital characteristics turned out to be quite exotic (for an overview see e.g. Marcy, Cochran \\& Mayor~\\cite{proto}).  Extrasolar giant planets were detected in very close-by orbits around their host stars with periods on the order of a few days, while orbital eccentricities at longer periods appear to be distributed quite uniformly. To date no Solar System analogue has been detected which is primarily due to the insufficient time baseline and long-term $RV$ precision of present Doppler surveys. However, the detection of Jovian-mass companions with $P>10$ yrs will become possible in the near future.  One of these long-term $RV$ surveys is the planet search program at the Coud\\'e Echelle Spectrometer (CES) at ESO La Silla, which was begun in Nov. 1992 using the 1.4 m CAT telescope. The highlight of this program so far was the discovery of an extrasolar giant planet in an Earth-like orbit around the young (ZAMS) and modestly active G0V star $\\iota$ Horologii (K\\\"urster et al.~\\cite{martin00}).  It is important to set such discoveries into the context of the complete results obtained by planet search programs. The pioneering study by Walker et al. (\\cite{walker}) first presented long-term (12 years) $RV$ results for a sample of 21 stars and discussed the implications of their non-detections on the occurrence of Jovian-type planets around solar-type stars. In an even earlier work, Murdoch et al. (\\cite{murdoch}) presented an analysis of their $RV$ measurements for 29 stars over 2.5 years, finding no brown dwarf companions within 10 AU in their sample. Since then $RV$ measurement precision (e.g. Butler et al.~\\cite{butler96}) and the size of target samples has increased dramatically. Extrasolar giant planets, which can be detected by present Doppler searches, exist around $\\approx 3 - 5\\%$ of the observed solar-type stars. Another study of the long-term $RV$ behaviour of a sample of stars was presented by Cumming et al. (\\cite{cumming}). These authors examined 11 years of $RV$ data collected by the Lick survey for 76 F-, G-, and K-type stars and derived companion limits for these stars.   With this work we present all $RV$ measurements of the CES survey and a complete analysis thereof over the time period of November 1992 to April 1998. During that time observations were performed with the same telescope and spectrograph configuration and thus form a homogenous data set. After April 1998 the CES instrument underwent major modifications and the results based on data collected after that point of time will be presented in an upcoming paper.  The structure of this paper is the following: Sect. 2 gives an overview of the CES planet search program, Sect. 3 presents the complete $RV$ results of the CES targets (appendix A displays the $RV$ measurements graphically for each star), Sect. 4 is a statistical examination of the CES $RV$ data where we perform tests to identify variable stars, the presence of linear and non-linear trends and periodic signals, in section 5 we set quantitative upper mass-limits for orbiting planets based on the $RV$ results (appendix B shows the derived limits for each survey star) and finally Sects. 6 and 7 contain the discussion and summary. ",
        "conclusions": "The planet around $\\iota$ Hor remains the single clear detection of an extrasolar planet by the CES Long Camera survey. This corresponds to a detection rate of $\\approx$ 3\\%, a value similar to other precise Doppler searches. The discovery of $\\iota$~Hor b demonstrated for the first time the feasibility of the $RV$ technique in planet detection in the case of young and thus moderately active stars.  Seven stars (19\\%) of the CES sample show minor signs of variability (minor in the sense that they pass one but fail at other tests): $\\zeta$~Tuc, HR~506, $\\zeta^{2}$~Ret, $\\epsilon$~Eri, $\\phi^{2}$~Pav and HR~8883. While for $\\epsilon$~Eri this is an indication for the presence of the highly eccentric RV signature of the planet, the cause of variability for the other stars remains unknown since no convincing periodicity or trends were found. For HR~8883 the $RV$ variations can be explained by the high instrinsic activity of this star (high X-ray luminosity).  Low amplitude linear $RV$ trends were found for the following 5 targets (14\\%): $\\beta$~Hyi, $\\alpha$~For, HR~6416, $\\epsilon$~Ind and HR~8501. For the known binaries $\\alpha$~For, HR~6416 and HR~8501 these trends agree well with the expected acceleration by the stellar secondary. The large scatter around the linear trend of $\\alpha$~For is probably also due to high stellar activity (again a high X-ray luminosity).  $\\beta$~Hyi and $\\epsilon$~Ind are identified as candidates for having long-period and probably stellar companions. But could these trends also be caused by planets? If we assume that the minimum orbital period would be 4 times the monitoring time span (i.e. $P\\approx20$ years) and that the $RV$ semi-amplitude is of the order of the $RV$ shift over the 5 years ($\\beta$~Hyi: $38~{\\rm m\\,s}^{-1}$, $\\epsilon$~Ind: $21~{\\rm m\\,s}^{-1}$) and $e=0$, the observed trends could be caused in the case of $\\beta$~Hyi by a planet with $m\\sin i \\approx 4~{\\rm M}_{\\rm Jup}$ at $a\\approx 7.6$ AU and for $\\epsilon$~Ind by an $m\\sin i \\approx 1.6~{\\rm M}_{\\rm Jup}$ companion at $a\\approx 6.5$ AU. Such planetary systems with a distant giant planet would resemble our Solar System more closely than the extrasolar planetary systems found so far. For $e>0$ orbits the period can even be much shorter than 20 years and we therefore conclude that although the linearity of the $RV$ trends points towards distant and previously unknown stellar companions, both stars constitute prime targets for follow-up observations by the CES planet search program.  $\\phi^{2}$ Pav has been earlier announced by our team as a possible candidate for having a planetary companion with an orbital period of about 43 days and $m\\sin i = 0.7~{\\rm M}_{\\rm Jup}$ (K\\\"urster et al.~\\cite{martin99}). This signal was found with a low confidence level and based on a preliminary analysis of a subset of the Long Camera data, using an early version of the $Radial$ code (Cochran \\& Hatzes~\\cite{radial}) to obtain the $RV$ measurements. The analysis of the complete data set of $\\phi^{2}$ Pav using the $Austral$ software did not confirm the presence of this companion. The total rms scatter over the entire 5 1/2 years is $35.4~{\\rm m\\,s}^{-1}$, slightly larger than the mean internal error of $31.3~{\\rm m\\,s}^{-1}$ for this star. No apparent Keplerian signal is present. This is consistent with results coming from the Anglo-Australian planet search (Butler et al.~\\cite{butler01}), who collected 7 measurements of $\\phi^{2}$ Pav over the course of 1 year, which reveal a total rms scatter of only $5~{\\rm m\\,s}^{-1}$ (the Anglo-Australian planet search uses the UCLES echelle spectrometer which covers a much larger spectral region and the entire $1000~{\\rm \\AA}$ of the I$_2$-reference spectrum, hence the higher precision of their results). However, we have identified in our much longer and higher sampled data a periodic signal of $\\approx 7$ days, again with low confidence (the FAP of this signal is still higher than $0.001$ but it appears in both the unbinned original $RV$ data as well as in the nightly averaged results). If the periodic signal is indeed real what could produce such an $RV$ signature? $\\phi^{2}$~Pav belongs to the $\\zeta$~Ret stellar kinematic group, a group of metal deficient stars with an age of $\\approx 5$ Gyr (del Peloso et al.~\\cite{peloso}). The iron abundance was determined as $[$Fe/H$] = -0.37$ by Porto de Mello \\& da Silva (\\cite{gustavo}) and as $[$Fe/H$] = -0.44$ by Edvardsson et al. (\\cite{edvardsson}). This low metallicity can account for the observed large $RV$ scatter, since fewer and shallower absorption lines in the small CES bandpass degrade our measurement precision. In fact $\\phi^{2}$ Pav has the second largest internal $RV$ error of the F-type stars in the CES sample. Based on H$\\alpha$ emission, $\\phi^{2}$~Pav appears to be slightly more active than the Sun (del Peloso et al.~\\cite{peloso}) and the star is already evolving into the subgiant phase (Porto de Mello \\& da Silva~\\cite{gustavo}). With this higher level of activity we suspect that the $P=7$ day $RV$ variation is in fact the stellar rotation period and that our $RV$ measurements are affected by cool spots in the photosphere of $\\phi^{2}$ Pav. These spots would modulate the $RV$ measurements with a typical timescale of $P_{\\rm Rot}$. Since these spots appear and disappear on short timescales compared to the monitoring duration and the overall activity level might change over 5 1/2 years, the amplitude as well as the phase of this modulation varies with time. Such a signal is therefore difficult to detect significantly, which is exactly what we observe here. The expected size of the subsurface convection zone for a low-metallicity F-type star is smaller than for a star of solar metallicity. Even with $P_{\\rm Rot}=7$ days such a star would not display a much larger activity level than $\\phi^{2}$~Pav due to the inefficiency of the dynamo. From the $v\\sin i = 6.7~{\\rm kms}^{-1}$ and $R_{*}=1.86~R_{\\odot}$ (Porto de Mello priv. comm.) and the $P_{\\rm Rot}$ value of $7$ days we derive a viewing angle of $\\approx 30^{\\circ}$. The continued monitoring of $\\phi^{2}$ Pav will demonstrate whether the $P=7$ day is robust and can be recovered with a higher confidence level.  Roughly 50\\% of the targets (18 stars) of the CES Long Camera survey show absolutely no sign of variability or trends in their $RV$ data. Within the given $RV$ precision of the CES Long Camera survey the following stars were found to be $RV$-constant: HR 209, HR 448, HR 753, $\\zeta^{1}$Ret, $\\delta$ Eri, HR 2667, HR 4523, HR 4979, HR 6998, HR 7373, HR 7703, HR 8323 and GJ 433.  In the cases of the binaries $\\kappa$~For, HR~2400, HR~3677 and $\\alpha$~Cen A \\& B (see Endl et al.~\\cite{michl01}) no sign of significant periodic signals were found in the $RV$ residuals after subtraction of the binary orbit. Interestingly, $\\kappa$~For does not show any excess scatter although based on its $L_{X}$-flux and H$\\alpha$-emission (Porto de Mello priv. comm.) it is an active star. Still, the residuals after subtraction of the binary orbit are consistent with our measurement errors.  The CES Long Camera survey is in $all$ cases sensitive to short-period (``51 Peg''-type) planets with orbital separations of $a < 0.15$ AU. This result confirms the general bias of precise Doppler searches towards short-period companions. For 22 stars of the CES Long Camera survey these mass-limits reach down into the sub-Saturn mass regime at $a \\approx 0.04$ AU.  For most stars the region where planets with $m \\sin i < 1~{\\rm M}_{\\rm Jup}$ could have been detected is confined to orbital separations of less than 1 AU. Beyond 2 AUs no planets with $m \\sin i \\approx 1~{\\rm M}_{\\rm Jup}$ were found to be detectable around any star of the survey. Subsequently, in order to detect a Solar System analogue the time baseline and (if possible) the $RV$ precision of the CES planet search has to be increased. Within the limitations of our numerical simulations ($e=0$, $P^{2/3}$ sampling) we can rule out the presence of giant planets within 3 AU of the CES survey stars according to the limits presented here (with the exceptions of HR 209, HR 8883 and periods inside the non-detectability windows).  Spectral leakage is the main cause for the windows of non-detectability. Even for well observed stars like e.g. $\\tau$~Cet or $\\beta$~Hyi these windows exist close to the seasonal one year period. This demonstrates how difficult the detection of $RV$ signals with a one year periodicity is.  The average detection threshold \\={F}$_{\\rm CES}$ for the examined 30 Long Camera survey stars is $2.75$, meaning that on average detectable planetary signals have $K$ amplitudes which exceed the noise level by a factor of $2.75$.  \\subsection{Outlook}  With the decommissioning of the Long Camera in April 1998, phase I of the CES planet search program came to an end. All results based on this homogeneous set of observations are included in this work or were already presented earlier.  Although the CES was modified quite substantially after that, with the installation of the Very Long Camera (VLC) yielding a higher resolving power of $R\\approx220,000$ and an optical fibre-link to the 3.6 m telescope being the most significant changes, the CES planet search program was continued using the same I$_2$-cell for self-calibration. This ensures the capability to merge the $RV$ results from phase I with the newer phase II data set without the need to compensate for velocity zero-point drifts as demonstrated for HR 5568 in Fig.~\\ref{gj570a_vlc}. The displayed $RV$ results now cover almost 3 years for this star (compared to 1 year of the Long Camera survey). For the intermediate time when the 1.4 m CAT and the 3.6 m telescope were used in combination with the VLC and the 2K CCD (which meant a reduced bandwidth of $\\approx 18~{\\rm \\AA}$ due to the higher spectral dispersion) we observe a small $RV$ offset of $\\approx 25~{\\rm m\\,s}^{-1}$. This offset can be explained by the difference in spectral regions which were analysed to obtain the $RV$s. However, after the VLC was equipped with a longer 4K CCD the spectral bandwidth was increased to $36.5~{\\rm \\AA}$. To assure that the $RV$ results are based on the same spectral regions we analyse both the Long Camera and the VLC data using a stellar template spectrum obtained with the most current instrumental setup (i.e. VLC \\& 4K CCD). A comparison of the $RV$ results derived with the current CES and Long Camera results does no longer show any velocity offset (see Fig.~\\ref{gj570a_vlc}). This demonstrates that the I$_2$-cell technique successfully compensates even for major instrumental setup changes. This guarantees a high long-term $RV$ precision and allows a smooth continuation of the CES planet search program. \\begin{figure} \\centering{ \\vbox{\\psfig{figure=MS2400f24.eps,width=9.0cm,height=6.0cm,angle=270}} \\par } \\caption[]{$RV$ monitoring of HR 5668 during the refurbishment of the CES. A comparison of the Long Camera results (full diamonds) with data collected with the new VLC and the 2K CCD (boxes and circles) show a slight offset. This offset disappears with the installation of the 4K CCD (triangles) which increased and equalized the spectral bandwidth (see text for details). The total rms scatter over the 3 years is $15.5~{\\rm m\\,s}^{-1}$, and $7.5~{\\rm m\\,s}^{-1}$ without the intermediate data (between the vertical dotted lines). } \\label{gj570a_vlc} \\end{figure}  The Very Long Camera at the CES is promising to increase the $RV$ precision of the CES planet search due to several reasons: the resolving power is doubled with respect to the Long Camera while the spectral bandwidth is not reduced by a large amount ($36.5~{\\rm \\AA}$ instead of $48.5$), and the $S/N$-ratio of spectra is higher due to the usage of image-slicers and the larger aperture of the 3.6 m telescope. With the successful merging of the new Very Long Camera data with the Long Camera survey and an expected better $RV$ precision the CES planet search might become sensitive to Solar System analogues in the near future."
    },
    "0207/astro-ph0207454_arXiv.txt": {
        "abstract": "{The three-point correlation function of cosmic shear, the weak distortion of the images of distant galaxies by the gravitational field of the inhomogeneous matter distribution in the Universe, is studied here. Previous work on three-point statistics of cosmic shear has mainly concentrated on the convergence, or on aperture measures of the shear. However, as has become clear recently for the two-point statistics of cosmic shear, the basic quantity that should be used is the correlation function: first, it is much easier to measure from observational data, since it is immune against complicated geometries of data fields (which contain gaps and holes, e.g. due to masking); second, all other (linear) two-point statistics can be expressed as integrals over the correlation function. The situation is the same for the three-point statistics. However, in contrast to the two-point correlation function, the invariants (with respect to rotations) of the shear three-point correlation function have not been employed yet. Here we consider the transformation properties of the shear three-point correlation function under rotations. We show that there are four complex linear combinations of components of the three-point correlation function, which we shall call `natural components', since they are multiplied just by a phase factor for arbitrary rotations, but do not mix. In particular, their moduli are invariant under rotations and thus (non-linear) invariants of the three-point correlation function.  In terms of these natural components, the invariance of the statistical properties of the shear field under parity transformations are easily obtained. Our results do not apply only to cosmic shear, but also to other quantities with the same mathematical properties -- that of a polar. For example, practically every relation derived here applies also to the polarization of the cosmic microwave background radiation. ",
        "introduction": "The weak gravitational lensing effect by the large-scale matter distribution of the Universe, called cosmic shear, has long been recognized as a unique tool to study the statistical properties of the cosmological (dark) matter distribution, without referring to luminous tracers of this distribution (Blandford et al. 1991; Miralda-Escude 1991, Kaiser 1992, 1998, Jain \\& Seljak 1997, Bernardeau et al.\\ 1997, Schneider et al.\\ 1998, van Waerbeke et al. 1999; Bartelmann \\& Schneider 1999; Jain et al.\\ 2000, White \\& Hu 2000; see Mellier 1999 and Bartelmann \\& Schneider 2001 for recent reviews).  Owing to the smallness of the effect, its actual measurement has only fairly recently been achieved, nearly simultaneously by several groups (Bacon et al. 2000; Kaiser et al. 2000; van Waerbeke et al. 2000; Wittman et al. 2000). This breakthrough became possible due to the usage of wide-field optical cameras and the development of special-purpose image analysis software specifically designed to measure the shape of very faint galaxies and to correct their ellipticity for effects of PSF smearing and anisotropy. By now, several additional cosmic shear measurements have been reported (Maoli et al. 2001; van Waerbeke et al. 2001; Rhodes et al. 2001; Bacon et al.\\ 2002; Refregier et al.\\ 2002; H\\\"ammerle et al. 2002; Hoekstra et al.\\ 2002), both from the ground and from HST imaging, partly with appreciably larger sky area than the original discovery papers.  In all of these papers, the cosmic shear signal detected was one related to the two-point correlation function of the shear, or some function of it, such as the shear dispersion or the aperture mass. To measure higher-order statistical properties of the cosmic shear, the quantity of data must be larger than for the second-order measures. It has been pointed out by a number of authors (e.g., Bernardeau et al.\\ 1997; Jain \\& Seljak 1997; Schneider et al.\\ 1998; van Waerbeke et al.\\ 1999; Hamana et al.\\ 2002) that the third-order statistics (e.g. the skewness) contains very valuable cosmological information, such as the density parameter $\\Omega_{\\rm m}$. In particular, the near-degeneracy between $\\sigma_8$ and $\\Omega_{\\rm m}$ in two-point cosmic shear statistics (see, e.g., van Waerbeke et al.\\ 2002) can be broken if the three-point statistics is employed.  Encouragingly, Bernardeau et al.\\ (2002a) have reported the detection of a third-order statistical signal in their cosmic shear survey.  Apart from the larger difficulty to obtain a measurement of the third-order statistics of the cosmic shear, there is also the problem of an appropriate statistical estimator for the third-order shear. Whereas for the second-order, the statistically independent shear measures are known, we are not in this position for the third-order shear statistics. We shall briefly summarize the situation for the two-point statistics, and explain why the three-point shear statistics is substantially more complicated in Sect.\\ 2 below.  In Sect.\\ts 3 we shall then define the components of the shear three-point correlation function (3PCF) and study their transformation behavior under spatial rotations.  From that, we shall then find in Sect.\\ts 4 the natural components of the shear 3PCF, which can be considered analogous to the natural components of the two-point correlation function of the shear. These components are `natural' in the sense that they have the simplest behavior under rotation transformations: each of the four complex natural components is just multiplied with a phase factor when an arbitrary rotation is applied, which in particular means that the moduli of these natural components are invariants under rotations.  We shall discuss the importance of these natural components in Sect.\\ts 5, where we also outline the perspectives of future work that can be based on the use of these natural components.  In an appendix, we shall consider the projection of the shear onto several particular reference points, defined by the various centers of a triangle.   We want to point out that all the relations derived in this papers are not confined only to cosmic shear. In fact, this paper investigates the three-point correlation function of a polar -- a polar is a two-component quantity which transforms under a rotation of the coordinate frame by a phase factor ${\\rm e}^{2{\\rm i}\\vp}$. Alternatively, a polar can be viewed as the trace-free part of a symmetric $2\\times 2$ matrix. In the case of cosmic shear, this matrix is the Hessian of the deflection potential. Another polar of great cosmological importance is the polarization of the cosmic microwave background. Hence, all the results presented below do equally well apply to the three-point correlation function of the CMB polarization. ",
        "conclusions": "The natural components of the shear 3PCF provide a generalization of the corresponding natural components $\\xi_\\pm$ of the two-point correlation function. Given that upcoming cosmic shear surveys will be substantially larger than the current ones, it is obvious that the 3PCF will be measurable in the future with a similar accuracy as the two-point correlation function in current surveys; in fact, a first significant measurement of the 3PCF has been reported in Bernardeau et al.\\ (2002a). Therefore, it is worth to explore the dependence of the 3PCF on various parameters, as has been done for the two-point function. In particular, the detailed study of the interrelations between $\\xi_\\pm$ and the underlying power spectrum, as presented in Crittenden et al.\\ (2002) and Schneider et al.\\ (2002a), shall be generalized to the 3PCF. Of course, this will be substantially more difficult from a technical point of view, and will be deferred to future work. Nevertheless, we can outline a few aspects of what can be expected from such work.  In close analogy to the two-point correlation function, one can expect that the natural components can be more easily calculated from the bispectrum of the underlying mass distribution than the individual components of the 3PCF. As is the case for $\\xi_\\pm$, one can expect that $\\Gamma^{(0)}$ probes the bispectrum in a different way than the $\\Gamma^{(k)}$, $1\\le k\\le 3$. In particular, the different natural components will have a different dependence on cosmological parameters.  The natural components of the 3PCF are not independent of each other; provided that the shear indeed is due to a surface mass density field, there should be integral relations which interrelate them. Again one should note the analogy with the two-point correlation function, where $\\xi_+$ can be obtained as an integral over $\\xi_-$ and vice versa (Crittenden et al.\\ 2002; Schneider et al.\\ 2002a). The interrelations between the components of the 3PCF provide a redundancy which can be profitably combined to reduce the noise in real measurements.  If the shear is not solely due to a surface mass distribution, the shear field may contain a B-mode contribution (e.g., from intrinsic alignments of the galaxies from which the shear is measured). In the case of the two-point correlation function, the presence of a B-mode can be probed from integral relations between the two correlation functions $\\xi_\\pm$ (Crittenden et al.\\ 2002; Schneider et al.\\ 2002a); it is expected that similarly in the case of the 3PCF, the integral relations between the natural components will be modified in the presence of a B-mode.  All linear three-point statistics of the shear can be expressed in terms of the shear 3PCF. This is obvious from the fact that the bispectrum of the surface mass density can be expressed in terms of the 3PCF; on the other hand, all linear three-point statistics are linearly related to the bispectrum, and can therefore be expressed directly in terms of the 3PCF. In particular, the third-order aperture mass statistics (Schneider et al.\\ 1998; van Waerbeke et al.\\ 1999) can be expressed as an integral over the 3PCF. It remains to be seen whether the third-order aperture mass statistics is as useful for a separation of the shear field into E- and B-modes as it is the case for the two-point statistics (Crittenden et al.\\ 2002; Schneider et al.\\ 2002a; for applications in cosmic shear surveys, see e.g.\\ Pen et al.\\ 2002; Hoekstra et al.\\ 2002).  Concerning practical measurements of the 3PCF, the procedure is straightforward: from a given catalog of galaxy images with position vectors $\\vc \\theta_i$ and measured ellipticities $\\eps_i$, triplets are selected. For each such triplet, the sides of the corresponding triangle can be calculated, and for practical reasons the largest side be called $x_3$. The three points are then labeled $\\vc X_l$, $1\\le l\\le 3$ in the (unique) way such that the orientation of the points is as described at the beginning of Sect.\\ts 3, and the longest side connects $\\vc X_1$ and $\\vc X_2$. Since the connecting vectors $\\vc x_l$ need to be calculated anyway, it is simplest to project the ellipticities along the sides of the triangle, or, with opposite sign, towards the orthocenter, without having to calculate any trigonometric function. The eight ($=2\\times 2\\times 2$) triple products of the projected ellipticities are calculated and summed up in eight three-dimensional bins. Those are conveniently labeled by a scale $x_3$, and two shape parameters, $q_1=x_1/x_3$, $q_2=x_2/x_3$, so that the grid of bins runs as $0\\le q_{1,2}\\le 1$, $0<x_3<\\infty$. After summing up over all triplets of points, for each of the eight components the sums in the bins are divided by the number of triplets contributing, which then yields an estimate of the 3PCF. Those can then be combined into the natural components for further analysis.  For a quantitative analysis of the three-point statistics of cosmic shear in relation to predictions of cosmological models, one can either use integrated properties of the shear 3PCF -- such as the aperture mass or the integral of $\\Gamma^{(\\rm B)}$ over an elliptical region as done in Bernardeau et al.\\ (2002b), or consider the (more noisy, but much more numerous) estimates of the natural components of the 3PCF directly. The first method yields one-dimensional functions of the three-point shear statistics depending on a scale parameter, and are thus very convenient for graphical displays. In contrast, using the full shear 3PCF employs multi-dimensional data which is difficult to display, but contains {\\it all} the information about third-order statistics from the data. Hence, even though the signal-to-noise in each bin of the 3PCF can be small, its overall information content cannot be smaller than that of any of the integrated quantities, and should therefore be employed in extracting cosmological information from the measurements. To obtain a reliable figure-of-merit for the comparison of the observational results with model predictions, one needs to know the covariance matrix of the shear 3PCF -- this, however, is a 64-components quantity depending on 6 arguments, and will be very difficult to obtain analytically (see Schneider et al.\\ 2002b for the difficulties of obtaining the covariance of the two-point correlation functions). Corresponding covariances of integrated quantities may be slightly easier obtainable. On the other hand, it is quite likely that figures-of-merit will have to rely heavily on future ray-tracing simulations through N-body-generated cosmological mass distributions, as in, e.g., van Waerbeke et al.\\ (1999), Jain et al.\\ (2000) and Bernardeau et al.\\ (2002a).  After this paper was finished, two recent preprints were posted on the Web which discussed very similar issues (Zaldarriaga \\& Scoccimarro 2002; Takada \\& Jain 2002b). In the latter one, the components $\\gamma_{\\mu\\nu\\lambda}$ of the correlation function were calculated from ray-tracing simulations; in particular they verified that all of the eight components of the 3PCF are non-zero in general and thus contain cosmological information."
    },
    "0207/astro-ph0207348_arXiv.txt": {
        "abstract": "We analyze spectra of 44 emission line stars detected in a low resolution optical spectroscopic survey of the galactic open cluster NGC~6871. The survey of 1217 stars is complete to $V = 14.9$ and includes stars with $V < 16.5$ between the zero age main sequence (ZAMS) and the $10^7$ yr pre-main sequence (PMS) isochrone. Of the 44 emission line stars in this survey, 28 show obvious emission in $H\\alpha$ and 16 have weak $H\\alpha$ absorption (compared to $H\\gamma$). We use the reddening to separate foreground and background stars from the cluster members; the position in the HR diagram (HRD)  or the presence of forbidden emission lines ([NII], [SII]) then yields the evolutionary status of the emission line stars. A comparison of the $H\\alpha$ spectral index distribution in NGC~6871 with the distribution of young stars in the Taurus-Auriga molecular cloud indicates that the late type PMS stars in our sample may be weak line T Tauri stars (wTTs). Many of these stars show [SII] ($36\\%$) and [NII] ($45\\%$) emission. ",
        "introduction": "Establishing mass-and age-dependent color-luminosity relations is important for understanding the physical processes leading from molecular clouds to main sequence stars. As in main sequence and post main sequence stars, observational HR-diagrams for PMS stars allow comparisons between observations and predictions of evolutionary models. Because isochrone fitting allows a more reliable measurement of distances and ages for star clusters than for single stars, the detection of PMS stars as members of star clusters offers an advantage over the observation of individual stars in more extended star forming regions. This advantage leads to accurate colors and luminosities for the possible PMS members of open clusters with well defined constraints on age and mass.  The aim of this paper is to describe 37 new emission line objects along the line of sight to the galactic open cluster NGC~6871. We identify these stars from a program to measure accurate spectral types, to derive the reddening, to separate cluster members from foreground and background stars, and to construct the Hertzsprung-Russell diagram (HRD) for cluster members. We will describe these results in a future paper.  With an age of roughly $10^7$ yr, NGC~6871 is at an important stage in the early evolution of an open cluster. Younger clusters with ages of $\\sim 10^6$ yr are often embedded in a bright nebulosity or a dark molecular cloud. Most of the young stars in these clusters have circumstellar disks \\citep{hais01}. Older star clusters with ages of $10^8$ yr have little nebulosity or molecular gas. Few stars in these cluster have disks \\citep{habi01,span01}. In NGC~6871, the lack of significant nebulosity or molecular gas allows optical spectroscopic observations to derive spectral types and luminosities of cluster members, and to measure IR excesses from circumstellar disks. Optical data provide a reliable measurement of the IMF for comparison with other young clusters \\citep[e.g.,][]{sles02}. Recent work suggests that circumstellar disks disappear on timescales of 5-10 Myr \\citep{hais01}. Observations of IR excess emission in NGC6871 thus yield better constraints on disk timescale.  NGC~6871 has been relatively well studied photometrically. The cluster is part of the Cyg OB3 association \\citep{garm92}, one of the largest nearby star formation complexes. The most recent UBV CCD photometry published by \\citet{mass95} has a limiting magnitude $V \\simeq16.5$. Several age estimates exist for the cluster members. From isochrone fitting, \\citet{mass95} derived ages of 2 - 5 Myr for the stars with $M > 25 M_\\odot$ (spectral type B0 and earlier). From uvby photometry for stars with spectral type earlier than B5 ($M > 4M_\\odot$), \\citet{reim89} estimated an age of 10 - 20 Myr. The difference between the two ages suggests either a large age spread among the cluster members, contamination from nearby stars in the local spiral arm \\citep{sles02}, or large uncertainties in age estimates derived from optical photometry. The presence of O stars in the cluster sets an upper limit for ages of high mass stars. According to \\citet{mass95}, the highest mass stars in the cluster have $M = 40~M_{\\odot}$ and main sequence lifetimes of 4-5 Myr \\citep{schal92}. However, \\citet{mass95} note that the cluster contains evolved 15 $M_\\odot$ stars with main sequence lifetimes of 11 Myr.  One of the goals of our program is to test the origin of the age difference by deriving accurate spectral types for cluster members. According to the WEBDA catalog of J-C. Mermilliod \\citep{merm95}, 12 stars in the cluster have reliable spectral types \\citep[e.g.,][]{rom51, hilt56,wal71,mass95}. The SIMBAD database lists another 151 stars with spectral classifications in the field of NGC~6871. Only 58 of these stars are included in the photometric database of \\citet{mass95} from which we selected our spectroscopic candidates. The discrepancy between the numbers of stars with spectral types in different catalogues can be explained by the incompletness of the WEBDA and the uncertain spatial bounderies of the catalogue of \\citet{mass95}. Some stars with spectral type must lie outside the CCD field of \\citet{mass95}. Our sample of 1217 stars will yield a good comparison between ages derived from photometric and spectroscopic techniques. This large sample also provides a good measurement of the frequency of $H\\alpha$ and other emission lines as a function of spectral types. Emission lines disappear when disk accretion ends; our data thus provide an estimate of the time scale for the end of disk accretion for intermediate and low mass stars.  We adopt an automatic approach and measure spectral types from spectral indices derived from digital spectra. This method allows robust measurement of spectral types and can be rapidly applied to large numbers of spectra. The accuracy of the method is comparable to visual estimates. We also use spectral indices to identify emission line objects. We compare the strength of the H$\\alpha$ and H$\\gamma$ indices and identify an object as emission line star if the H$\\alpha$ absorption is weak relative to H$\\gamma$. This method yields an automatic identification of emission line stars in cases where the emission is not obviously detectable by visual examination.  In this first paper of the series we report the discovery of 37 new emission line stars and analyze their spectra. Eight stars with accurate spectral types are known emission line stars (we detect seven of these stars). Here, we extend this sample to 45 stars and begin to separate cluster members from foreground and background stars. ",
        "conclusions": "We have begun a spectroscopic survey of the galactic open cluster NGC~6871. Our spectra cover $3700-7400$\\AA~at $\\simeq 6$\\AA~resolution. The survey is complete to $V=14.9$ and includes stars with $V < 16.5$ between the ZAMS and $10^7$ yr PMS isochrone.  Using $H\\alpha$ and $H\\gamma$ spectral indices, we identify 44 emission line stars. The spectral types of these stars range from O9 - K8; one star may have an M5 spectral type. Eight stars have [NII] emission; six stars have [SII] emission. A few stars may have Li I $\\lambda6708$ absorption; higher resolution spectra with good signal-to-noise would test these detections.  We use the reddening $E_{B-V}$ to divide the emission line sample into foreground stars, background stars, and cluster members. Nine stars have very low reddening and are probably foreground stars. Eighteen stars are much more highly reddened than the cluster stars. However, most of these lie above the cluster main sequence on the HRD, suggesting that these stars are cluster stars with extra (circumstellar) reddening. We thus conclude that 24 emission line stars are cluster members.  The 9 foreground stars fall into two groups. Four stars have essentially no reddening and are therefore nearby stars with $d \\lesssim 300$ pc. These might be part of a nearby association, similar to the TW Hydrae association \\citep{reza89,greg92,zuck01} or the $\\eta$ Chamaeleontis cluster  \\citep{mam99}. The other five stars also appear to lie at the same distance and might form another association or group along the line of sight to NGC~6871.  The emission line stars in the cluster consist of two groups. Stars with spectral type $\\lesssim$ A0 lie on the main sequence and have ages of $\\rm t < 1.4x10^7$ yr \\citep{siess00}. Some of these stars may be classical Be stars. Near infrared data would allow a test of the evolutionary status of these stars. \\citet{lada92} found that PMS Be stars and classical Be stars occupy well-defined, different regions of the J-H, H-K color-color diagram. A few F-type emission line stars may belong to this group. Several stars with spectral types A - G lie above the cluster main sequence and may be PMS stars. The distribution of $H\\alpha$ indices for these stars coincides with the peak observed for wTTs in the Taurus-Auriga molecular cloud. If these are PMS stars, they are wTTs not cTTs. Some of these stars show [NII], [SII] emission and may have LiI absorption, which provides some evidence for their PMS status.  The discovery of [SII] and [NII] in apparent wTTs in NGC~6871 is a new and unexpected result of our survey. Previous observations of wTTs in other regions rarely detect [NII] or [SII] emission \\citep{keny98,harti95}. We detect [NII] emission in $45\\%$ and [SII] emission in $36\\%$ of our candidate wTTs; none of these show [OI] emission features in their spectra. In many PMS stars, [OI], [NII], and [SII] emission lines are associated with disk accretion phenomena \\citep[e.g. jets and disk-wind interactions;][]{hart89,harti95,rei01}. As a class, wTTs show very little evidence for a wind; the lack of an infrared excess indicates that they have already lost their disk. Thus the origin of [NII] and [SII] emission in our candidate wTTs is uncertain: If this emission is associated with these stars, the emission must be excited by some mechanism other than accretion. If the emission is not associated with these stars, a faint HII region in NGC 6871 is a plausible alternative source of forbidden line emission. Our test for nebular background emission in Section 4 appears to rule out this source of emission, but high resolution spectra would provide a better test."
    },
    "0207/astro-ph0207381_arXiv.txt": {
        "abstract": "{ We present medium resolution (R$\\sim$1500) ISO-SWS 2.4--45$\\mu$m spectra of a sample of 29 galaxies with active nuclei. This data set is rich in fine structure emission lines tracing the narrow line regions and \\mbox{(circum-)}nuclear star formation regions, and it provides a coherent spectroscopic reference for future extragalactic studies in the mid-infrared. We use the data set to briefly discuss the physical conditions in the narrow line regions (density, temperature, excitation, line profiles) and to test for possible differences between AGN sub-types. Our main focus is on new tools for determining the properties of dusty galaxies and on the AGN-starburst connection. We present mid-IR line ratio diagrams which can be used to identify composite (starburst + AGN) sources and to distinguish between emission excited by active nuclei and emission from (circum-nuclear) star forming regions. For instance, line ratios of high to low excitation lines like [O\\,IV]25.9$\\mu$m/[Ne\\,II]12.8$\\mu$m, that have been used to probe for AGNs in dusty objects, can be examined in more detail and with better statistics now. In addition, we present two-dimensional diagnostic diagrams that are fully analogous to classical optical diagnostic diagrams, but better suited for objects with high extinction. Finally, we discuss correlations of mid-infrared line fluxes to the mid- and far-infrared continuum. We compare these relations to similar relations in starburst galaxies in order to examine the contribution of AGNs to the bolometric luminosities of their host galaxies. The spectra are available in electronic form from the authors. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207662_arXiv.txt": {
        "abstract": "We have used optical spectroscopy to investigate the active galaxy populations in a sample of 20 nearby Abell clusters. The targets were identified on the basis of 1.4 GHz radio emission, which identifies them as either AGN or galaxies forming stars at rates comparable to or greater than that of the Milky Way. The spectra were used to characterize the galaxies via their emission and absorption features.  The spectroscopy results reveal a significant population of star forming galaxies with large amounts of nuclear dust extinction. This extinction eliminates bluer emission lines such as [OII] from the spectra of these galaxies, meaning their star formation could easily be overlooked in studies which focus on such features. Around $20\\%$ of the cluster star forming galaxies have spectra of this type. The radial distributions of active galaxies in clusters show a strong segregation between star forming galaxies and AGN, with star forming galaxies broadly distributed and AGN preferentially in the cluster cores.  The radial distribution of the dusty star forming galaxies is more centrally-concentrated than the star forming galaxies in general, which argues that they are a consequence of some cluster environmental effect. Furthermore, we note that such galaxies may be identified using their $4000\\mbox{\\AA}$ break strengths.  We find that discrepancies in reported radio luminosity functions for AGN are likely the result of classification differences. There exists a large population of cluster galaxies whose radio fluxes, far-infrared fluxes, and optical magnitudes suggest their radio emission may be powered by stars, yet their spectra lack emission lines. Understanding the nature of these galaxies is critical to assessing the importance of AGN in the radio luminosity function at low luminosities. We also find that regardless of this population, the crossover point where the radio luminosity function is comprised equally of star forming galaxies and AGN occurs at lower luminosities in clusters than the field. This is likely a simple consequence of the reduction in star formation in cluster galaxies and the morphological mix in clusters compared to the field. ",
        "introduction": "The role of environment in the evolution of galaxies has been a central question in extragalactic astronomy. One of the stronger pieces of observational evidence for galaxy evolution in the environment of rich clusters of galaxies is the Butcher-Oemler effect \\citep{butc1978,butc1984}. Subsequent photometric work to examine this effect and its causes has confirmed its existence with larger samples \\citep{rako1995,lubi1996,marg2000,marg2001} and revealed more evidence for evolution in cluster populations. \\citet{dres1997} used HST imaging of galaxies in distant clusters to confirm that the morphology-density relationship still holds, but with important differences. Specifically, the fraction of S0 galaxies is several times lower in distant clusters, arguing that while elliptical galaxies formed in the distant past the S0 galaxies seem to have formed after cluster virialization.  Spectroscopic studies motivated by the Butcher-Oemler effect have also produced evidence for evolution in clusters \\citep[e.g.,][]{dres1983,couc1987,dres1992,dres1999}. \\citet{dres1983} uncovered an unusual population of galaxies in distant clusters, dubbed ``E+A'' galaxies, and noted that their frequency in clusters increased at higher redshift. While these galaxies are more common in clusters than the field at high redshift, their presence in the field (even locally) implies that they can not be explained solely by environmental effects exclusive to cluster environments \\citep{zabl1996}. The spectroscopic catalog of \\citet{dres1999} noted the presence of an additional important population of galaxies, those which exhibited strong Balmer absorption with slight emission of the star formation indicator, [OII] $\\lambda3727$. In general, the frequency of these galaxies mimics that of the E+A galaxies, being more common at higher redshift. \\citet{pogg1999} applied evolutionary models to investigate these galaxies and argued that they were most likely dusty starbursts, and drew parallels to galaxies in the nearby universe \\citep{liu1995,pogg2000}. Further models confirm that these dusty starbursts can be the precursors to the post-starburst galaxies \\citep{shio2001}.  Ironically, the exciting evidence for evolution in clusters has diverted attention from nearby clusters. Since clusters in the nearby universe seem to exhibit less activity (in the sense of harboring galaxies with strong current or recent star formation), they are often overlooked in favor of their predecessors at intermediate redshift \\citep[although see][]{cald1993,cald1997,rose2001}. This is unfortunate, since nearby clusters can more easily be investigated in greater detail. Not only can such investigations provide accurate benchmarks for studies of clusters at higher redshift, but they can influence future high redshift studies by uncovering trends not easily apparent given the practical limits of present-day telescopes and detectors.  In this paper, a large sample of active galaxies drawn from nearby clusters is used to assess galaxy evolution in the cluster environment. The galaxies are radio selected, with 1.4 GHz luminosities consistent with that of the Milky Way or greater. Radio emission is an excellent indicator of activity, in the form of either star formation or active galactic nuclei \\citep{cond1992}. The sample is primarily that identified in \\citet{mill2001}, augmented by comparable radio observations of two additional clusters. Cluster membership of the radio galaxies has been achieved via optical spectroscopy, producing a net sample of 411 confirmed members in 20 Abell clusters (plus an additional 12 potential members for which velocities were not obtained). Much of the data used in this paper correspond to quality long-slit spectra of a large subset of the confirmed cluster members. Classification of the galaxies is performed via this spectroscopy, including line ratio diagnostics to characterize the activity (star formation, AGN, or both) and assess the importance of dust extinction. The distributions of these various classes of galaxies are then compared to elucidate the role of environment on cluster galaxy evolution.  Specifically, this paper examines the sample collectively to explore several issues of galaxy evolution in the cluster environment. First, what types of active galaxies are seen in nearby clusters, and in what numbers? In part, this amounts to understanding the relative importance of star formation and AGN in nearby clusters. Within the star forming galaxies, the spectra also provide a means to assess the importance of dust extinction. Second, what can the distributions of these active galaxies tell us about evolutionary mechanisms in clusters? Because the presence of radio emission implies recent activity, the locations of the active galaxies are tied to their local environment. Consequently, the distributions will elucidate the importance of various evolutionary models. Lastly, what can the sample tell us about the composition of the radio luminosity function and its application to studies of star formation at higher redshift? While these issues are addressed using the sample collectively, in future papers we will make cluster-by-cluster comparisons in order to shed light on issues of specific cluster parameters and their effects on member galaxies.  The paper is organized as follows: In Section \\ref{sec:data}, the data are summarized. This includes an overview of the sample and observations, as well as the procedures used to classify the cluster galaxies on the basis of their optical spectra. The different classes of radio galaxies are then investigated in Section \\ref{sec:analysis}. This includes an assessment of the importance of dust extinction and an investigation of how the different classes of radio galaxies are distributed in the clusters. Section \\ref{sec:rlf} presents the cluster radio luminosity function, broken down by activity class. This breakdown also leads to a deeper exploration of the assigned activity classes, in particular the nature of those galaxies with absorption-line spectra.  The implications of all these results are discussed in Section \\ref{sec:discuss}, followed by a brief summary of the conclusions in Section \\ref{sec:conclude}. Unless otherwise noted, we have assumed $H_o = 75$ km s$^{-1}$ Mpc$^{-1}$ and $q_o = 0.1$ throughout. ",
        "conclusions": "\\label{sec:conclude}  We have used a comprehensive sample of active galaxies in 20 nearby Abell clusters to assess galaxy evolution in the cluster environment. Our multiwavelength data have revealed a number of interesting results:  \\begin{itemize}  \\item{Classification of galaxies using only bluer spectral features can mask a population of cluster star forming galaxies. These galaxies exhibit dust extinctions great enough to greatly reduce or even remove the presence of [OII] from their spectra. These galaxies represent somewhere near $20\\%$ of all star forming galaxies in nearby clusters.}  \\item{The distributions of radio galaxies in nearby clusters appear to be segregated by activity. Normal star forming galaxies are broadly distributed in clusters, extending well past the classical Abell radius whereas AGN are centrally concentrated. Galaxies which are actively forming stars but show unusually strong dust extinction appear to have a different distribution. Their core radius is intermediate to those of the star forming galaxies and the AGN, and they seem to peak at $\\sim0.5$ Mpc while avoiding the very cores of clusters. This argues for such galaxies being the consequence of cluster environmental effects.}  \\item{Spatially, the dust extinction in the dusty star forming galaxies is nuclear. Their Balmer decrements are large when measured for nuclear apertures (2\\arcsec), yet normal when measured off the nucleus. Consequently, they appear to be the result of nuclear starbursts excited by some aspect of the cluster environment.}  \\item{In higher redshift clusters for which H$\\alpha$ is beyond the range of the optical spectrum, candidates for these dusty star forming galaxies may be identified on the basis of their $D_{4000}$ values. The dusty star forming galaxies have signficantly lower $D_{4000}$ than non-star forming galaxies whose shorter wavelength spectra are characteristic of old stellar populations.}  \\item{There is a population of galaxies with spectra dominated by old stellar populations whose radio luminosities place them in the realm of either star forming galaxies or low luminosity AGN. These galaxies are frequently weak FIR detections, and ratios of their radio, optical, and FIR fluxes do not unambiguously identify them as star forming or AGN. It is likely that the assumed classification of these galaxies is responsible for differences in the RLFs reported by different authors.}  \\item{The radio luminosity at which AGN and star formation contribute equally to the overall RLF is greater in clusters than it is in large volume-limited samples. This effect can be traced to the lower fractions of star forming galaxies in clusters relative to the field.}  \\end{itemize}  Even though these results have been obtained using a large spectroscopic database, additional spectroscopy would prove very useful. The fairly large fraction of galaxies classified on the basis of radio and FIR fluxes means that more precise estimates of the frequency of various spectral types (including dusty star forming galaxies, AGN, etc.) can not be made.  While these results have been obtained from the collective sample, additional information can be obtained by assessing each cluster individually in comparison to the rest of the sample. Specifically, are these evolutionary results more common in some clusters than others? If so, are they a function of factors such as richness or dynamical state? It is possible that an understanding of these issues will help to understand any underlying causes for discrepancies in evolutionary studies at higher redshift, if these causes can be related to selection effects in the different higher redshift studies."
    },
    "0207/hep-ph0207211_arXiv.txt": {
        "abstract": "The simplest explanation for early time acceleration (inflation) and the late time acceleration indicated by recent data is that they have a common origin. We investigate another generic cosmological implication of this possiblity, that the baryon asymmetry of the universe may be generated in such models. We identify several novel features of baryogenesis in such a universe, in which a rolling scalar field is always part of the cosmological energy budget. We also propose a concrete mechanism by which the baryon asymmetry of the universe may be generated in this context. We analyze the generic properties of and constraints on these cosmologies, and then demonstrate explicitly how a complete cosmology may develop in some specific classes of models. ",
        "introduction": "\\label{intro} Rolling scalar fields are a mainstay of modern cosmology. This is perhaps best-illustrated by the inflationary paradigm~\\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}, in which most implementations involve a scalar field rolling towards the minimum of its potential in such a way that the potential energy of the field is the dominant component of the energy density of the universe. There are, however, many other cosmological instances in which scalar fields are invoked.  During the last few years a new consistent picture of the energy budget of the universe has emerged. Large scale structure studies show that matter (both luminous and dark) contributes a fraction of about 0.3 of the critical density, while the position of the first acoustic peak of the cosmic microwave background power spectrum indicates that the total energy density is consistent with criticality. The discrepancy between these two measurements may be reconciled by invoking a negative pressure component which is termed {\\it dark energy}. While there are a number of different observational tools to study dark energy -- number counts of galaxies~\\cite{Newman:1999cg} and galaxy clusters~\\cite{Haiman:2000bw} for example -- the most direct evidence to date comes from the light-curve measurements of intermediate redshift type IA supernovae~\\cite{Riess:1998cb,Perlmutter:1998np}. Consistency between these observations and others such as weak gravitational lensing~\\cite{Huterer:2001yu} and large scale structure surveys~\\cite{Hu:1998tk} implies that the dark energy $X$ satisfy $\\Omega_X \\sim 0.7$ and that the equation of state be~\\cite{Perlmutter:1999jt,Wang:1999fa} \\begin{equation} w_X \\equiv \\frac{p_X}{\\rho_X}\\leq -0.6 \\ , \\end{equation} leading to the acceleration of the universe.  It is of course possible that this mystery component is a cosmological constant $\\Lambda$, for which $w_{\\Lambda}=-1$. However, understanding the nature of such an unnaturally small $\\Lambda$ is at least as difficult as undestanding one that is zero. Alternatively, it has been suggested~\\cite{Wetterich:fm}-\\cite{Caldwell:1997ii} that if the cosmological constant itself is zero, the dark energy component could be due to the dynamics of a rolling scalar field, in a form of late-universe inflation that has become known as {\\it quintessence}. Although there are a number of fine-tuning problems associated with this idea, it does provide a way to ensure the late-time acceleration of the universe, albeit at the expense of introducing a second (after the inflaton) cosmologically relevant rolling scalar field. While not addressing the cosmological constant problem itself, and suffering from fine-tuning, quintessence itself has the advantage of avoiding a future horizon in space-time, and hence makes consistency with what is known about perturbative string theory more likely.  It is natural to wonder whether the inflaton and the quintessence field might be one and the same~\\cite{Spokoiny:1993kt}. In fact, specific models for this have been proposed~\\cite{Spokoiny:1993kt}-\\cite{Dimopoulos:2001ix}. Clearly such models are attractive because we need only postulate a single rolling scalar, but may be problematic either theoretically or phenomenologically.  In this paper we investigate how we may further limit the proliferation of rolling scalar fields required in modern cosmology by studying how the scalar field responsible for late-time acceleration of the universe might also solve another outstanding cosmological puzzle. Specifically we will be interested in the role that such a field may play in the generation of the baryon asymmetry of the universe. The spectacular success of primordial nucleosynthesis requires that there exist an asymmetry between baryons and antibaryons in the universe at temperatures lower than an MeV. This is quantified by the requirement \\begin{equation} \\label{BAUobserved} 4\\times 10^{-10}\\leq \\eta \\equiv \\frac{n_B}{s} \\leq 7\\times 10^{-10} \\ , \\end{equation} where $n_B\\equiv n_b - n_{\\bar b}$, with $n_{b({\\bar b})}$ the number density of (anti)baryons and $s$ is the entropy density. To generate such an asymmetry, the underlying particle physics theory must satisfy three necessary conditions -- the Sakharov conditions~\\cite{Sakharov:dj}. These are baryon number $B$ violation, the violation of the discrete symmetries $C$ and $CP$ and a departure from thermal equilibrium, this last condition resulting from an application of the $CPT$ theorem. In this paper we are interested in how these conditions may be met within the context of dark energy models.  The relationship between early-time acceleration -- inflation -- and baryogenesis has been explored in some detail (for example see~\\cite{Affleck:1984fy}-\\cite{Nanopoulos:2001yu}). Here we investigate the opposite regime, that the quintessence field may be associated with the generation of the baryon asymmetry. Naturally, it would be particularly efficient if a single scalar field could be responsible for three fundamental phenomena in cosmology -- inflation, baryogenesis and dark energy, and indeed we will show that baryogenesis occurs quite generically in models in which a single scalar is responsible for the two periods of cosmic acceleration.  The outline of this paper is as follows. In section~\\ref{quintinf} we will review some details about quintessence and explain how inflation and quintessence may be unified by generalizing the quintessential inflation model of Peebles and Vilenkin~\\cite{Peebles:1998qn}. In section~\\ref{qbg} we will describe how quintessence and quintessential inflation may naturally yield a baryon asymmetry without the introduction of any new fields into the theory. We term this model {\\it quintessential baryogenesis}, borrowing the phrasing from Peebles and Vilenkin. This turns out to depend to some extent on the details of quintessential inflation. In section~\\ref{sec:constraints} we will discuss experimental and astrophysical constraints on our models and comment on how we may test the physics involved. We offer our comments and conclusions in the final section of the paper. ",
        "conclusions": "\\label{conclusions} Modern particle cosmology concerns the search for dynamical explanations for the initial conditions required by the standard FRW cosmology. Particular attention has been paid recently to one of these initial condition problems, that of the size of the vacuum energy contribution to the total energy density of the universe. The root of this issue is that vacuum energy, or something approximating it, can lead to the acceleration of the universe.  The best fit cosmology to all current observational data is one in which the universe undergoes two separate epochs of acceleration. The first of these, inflation, is the only clear way to seed adiabatic, scale-free perturbations in the cosmic microwave background radiation. The second epoch, that of dark-energy domination, is required to simultaneously understand the power spectrum of the CMB and the expansion of the universe at intermediate redshifts as revealed by type IA supernovae. These requirements have led cosmologists to introduce a new scalar field to account for the newly required late-time acceleration.  In this paper we have extended the approach of Spokoiny~\\cite{Spokoiny:1993kt} and of Peebles and Vilenkin~\\cite{Peebles:1998qn} in exploring the extent to which the dynamics of a single scalar field can be responsible for setting multiple parameters required by the standard cosmology. Our contribution has been to generalize the mechanisms by which inflation and dark energy domination may be due to a single scalar and to introduce the idea that this same rolling scalar might be responsible for generating the baryon asymmetry of the universe. The mechanism that we propose, quintessential baryogenesis, is an application of the spontaneous baryogenesis model of Cohen and Kaplan\\cite{Cohen:1988kt} to the quintessential inflation case.  It seems a particularly powerful idea to us that a single rolling scalar field might be responsible for a number of the fundamental initial conditions required to make the standard cosmology work. In the case of the baryon asymmetry, this allows us to associate the existence of an asymmetry with the spontaneous breaking of $CPT$ and the direction of the rolling of the scalar field. \\begin{figure} \\epsfig{file=graphic.eps, height=3.3in, width=6.4in} \\caption{The evolution of the universe in this model. As the scalar field rolls down its potential the universe goes through a succession of phases, beginning with inflation, generating the baryon asymmetry along the way, and ending with dark energy domination.} \\label{fig:graphic} \\end{figure}  The evolution of the universe we envisage may be summarized as follows (see figure~\\ref{fig:graphic}). At the earliest times in the universe, inflation occurs due to the potential energy dominance of the field $\\phi$ which begins rolling at very large and negative values. Inflation ends when the kinetic energy of the scalar field becomes important and the slow-roll conditions are violated. Since our potential does not have a minimum at finite $\\phi$, unlike typical inflationary models, conventional reheating does not occur. Instead, matter is created gravitationally due to the mismatch of vacuum states between the approximately de-Sitter state of inflation and that of the kination era. Since kinetic energy density redshifts more rapidly that radiation energy density, the universe eventually becomes radiation-dominated. At this stage the rolling scalar has negligible effect on the expansion rate of the universe. However, the direction of rolling spontaneously violates $CPT$. If $\\phi$ couples to other fields, as we expect it to generically, then the expectation value of the baryon number operator in this background in thermal equilibrium is nonzero. Thus a baryon excess is generated. After the electroweak phase transition baryon number violation is no longer effective in the universe and the baryon number existing at that time is frozen in. The scalar field continues to evolve and in the late universe, after matter-domination has begun, its potential energy can once again become dominant leading to a new period of dark-energy domination.  The couplings required to make quintessential baryogenesis effective may be generated in a number of different ways, for example by gravitational effects coupling the inflaton/dark energy sector to visible sector fields. We have considered the current experimental constraints on the necessary operator and have found that there exist considerable regions of parameter space in which our mechanism is consistent. Further, it is possible that a restricted region of this parameter space may be accessible to future experiments.  We have left a number of questions unanswered and will return to them in future work. Perhaps the most pressing issue is one that plagues rolling scalar models of dark energy in general, namely the question of technical naturalness of the potentials involved, and their stability to quantum corrections. We have omitted any discussion of this here, while laying out the general features of the model, but these issues must be addressed to put our mechanism on firmer ground. For example, it may be most natural to identify the field $\\phi$ with a pseudo-Goldstone boson~\\cite{Dolgov:1996qq,Frieman:1995pm}, since its coupling to the current $J^{\\mu}$ is derivative. However, this is a general issue for quintessence models, and is not specific to our baryogenesis mechanism. We have therefore chosen to concern ourselves with this issue separately.  Taken at face value, current observations imply that our universe is entering an accelerating phase that may be governed by the rolling of a scalar field. If we are to understand the physics of such a field then it is important that we investigate other ways in which it may impact cosmology and particle physics. In particular, if, as we have suggested here, the field is responsible for the generation of the baryon asymmetry, then the result will be a more economical and attractive cosmology."
    },
    "0207/astro-ph0207094_arXiv.txt": {
        "abstract": "After almost 2.5 years of actively accreting, the neutron star X-ray transient and eclipsing binary MXB 1659--29 returned to quiescence in 2001 September.  We report on a {\\it Chandra} observation of this source taken a little over a month after this transition. The source was detected at an unabsorbed 0.5--10 keV flux of only $(2.7 - 3.6) \\times10^{-13}$ \\funit, which implies a 0.5--10 keV X-ray luminosity of approximately $(3.2 - 4.3) \\times10^{33}\\,\\,\\, (d/10\\,\\,\\,{\\rm kpc})^2$ \\Lunit, with $d$ the distance to the source in kpc. Its spectrum had a thermal shape and could be well fitted by either a blackbody with a temperature $kT$ of $\\sim0.3$ keV or a neutron star atmosphere model with a $kT$ of $\\sim0.1$ keV.  The luminosity and spectral shape of MXB 1659--29 are very similar to those observed of the other neutron star X-ray transients when they are in their quiescent state. The source was variable during our observation, exhibiting a complete eclipse of the inner part of the system by the companion star. Dipping behavior was observed before the eclipse, likely due to obscuration by an extended feature in the outer part of a residual accretion disk.  We discuss our observation in the context of the cooling neutron star model proposed to explain the quiescent properties of neutron star X-ray transients. ",
        "introduction": "}  During outburst episodes, neutron star X-ray transients can be detected at luminosities of $\\sim10^{36-38}$ \\Lunit~(e.g., Chen, Shrader, \\& Livio 1997).  During those outbursts, the transients are very similar to the persistent sources with respect to their X-ray properties. The high X-ray luminosity is very likely due to the accretion of matter onto the neutron star.  These transients are characterized by their bright outbursts, but most of the time they spend in a quiescent state in which they are orders of magnitude dimmer at all wavelengths. Fortunately, using sensitive imaging instruments, we are still able to detect them at X-ray luminosities of $\\sim 10^{32-34}$ \\Lunit~(e.g., van Paradijs et al. 1987; Asai et al. 1996, 1998). The high sensitivity camera's aboard {\\it Chandra} and {\\it XMM-Newton} are well suited to detect quiescent systems and obtain good X-ray spectra for the brightest systems (see, e.g., Daigne et al. 2002; in 't Zand et al. 2001; Rutledge et al. 2001a, 2001b; Wijnands et al. 2001b, 2002b).  To explain the low quiescent X-ray properties, several models have been developed. For example, the X-rays could be due to the residual accretion of matter onto the neutron star or magnetospheric boundary, or the pulsar emission mechanism might be active (see, e.g., Stella et al. 1994; Corbet 1996; Campana et al. 1998b; Menou et al. 1999; Campana \\& Stella 2000; Menou \\& McClintock 2001). Currently the most successful model is that in which the X-rays are due to the thermal emission from the neutron star surface, which will be referred to as 'the cooling neutron star model'.  \\subsection{The cooling neutron star model}  In the cooling neutron star model (e.g., van Paradijs et al. 1987; Campana et al. 1998b; Brown, Bildsten, \\& Rutledge 1998 and references therein) the emitted radiation below a few keV is thermal emission originating from the neutron star surface.  Brown et al. (1998) argued that the neutron star core is heated by the nuclear reactions occurring deep in the crust when the star is accreting and this heat is released as thermal emission during quiescence. If the quiescent emission is dominated by the thermal emission of the cooling neutron star, then the quiescent luminosity should depend on the time averaged (over $10^{4-5}$ years) accretion luminosity of the system (Campana et al. 1998b; Brown et al. 1998).  Thus, the quiescent luminosities of the detected systems can directly be compared with the predicted ones obtained from estimates of the long term accretion history of the sources.  The neutron star cooling model also gives clear predictions for the spectral shape of the quiescent X-ray spectrum, which should be thermal.  Although a simple blackbody model can be fitted to the data, the obtained radii of the emitting regions (of the order of only a few kilometers) are considerably lower than the predicted radii of neutron stars (Shapiro \\& Teukolsky 1983). To circumvent this discrepancy, it has been proposed that quiescent neutron star systems do not emit a true blackbody spectrum but a modified one (Brown et al. 1998). When using blackbody models to fit such modified spectra, the effective temperatures will be overestimated and the emitting areas underestimated. By fitting more realistic models, such as the so-called 'neutron star atmosphere models' (the non-magnetic models are appropriate for quiescent neutron star systems; e.g., Zavlin, Pavlov, \\& Shibanov 1996 ), to the X-ray data, emitting radii were obtained which are consistent with the expected radii of neutron stars (Rutledge et al. 1999, 2000).   The cooling neutron star model cannot fully explain all characteristics of the quiescent emission. For example, the power-law shaped spectral component which dominates the quiescent spectra above a few keV in several systems\\footnote{Note that not in all detected quiescent systems this power-law component could be detected and that the flux ratio of the power-law component with the thermal component varies considerably between sources.} (e.g., Asai et al. 1996, 1998; Campana et al. 1998a) cannot be explained by the cooling models. It is conceivable that this component might be described by one or more of the alternative models discussed above (in particular the residual accretion model). However, the observational results on this component and our understanding of its nature are very limited.   \\subsection{The quasi-persistent X-ray transients}  Recently, a sub-group of neutron star X-ray transients has received extra attention because of their potential to test the cooling neutron star model and to determine some of the physical properties of the neutron star crust and core. These particular transients do not have traditional outbursts which only last weeks to at most a few months, but instead they stay active for several years to over a decade (and maybe even longer). These systems have been called long-duration transients or quasi-persistent sources (e.g., Wijnands et al. 2001b, Wijnands 2002). The long outburst behavior of those sources might be related to the extended episodes (several months to several years) of low-level activity seen in other transients usually after they have exhibited bright outbursts (e.g., in 4U 1630--47, Aql X-1, 4U 1608--52, or SAX J1808.4--3658; Kuulkers et al. 1997; Bradt et al. 2000; Wijnands et al. 2001c; Wachter et al. 2002), although those episodes are generally less luminous ($<10^{36}$ \\Lunit) than the outbursts of the quasi-persistent sources ($10^{36} - 10^{37}$ \\Lunit).   In 'ordinary' (i.e., short-duration) transients, the accretion of matter will have only a very minor effect on the thermal state of the crust, but for these quasi-persistent sources the prolonged accretion episodes can heat the crust to high temperatures, considerably higher than that of the neutron star core (see Rutledge et al. 2002). When those systems become quiescent again, it might take years to decades for the crust to return to thermal equilibrium with the core and the initial quiescent properties of those systems might be dominated by the crust emission and not by the state of the core (as is the case in ordinary transients). Monitoring observations of those systems in quiescence might even allow one to follow the cooling of the crust from which the heat conductivity of the crust can be determined (Rutledge et al. 2002).  Recently, one of the quasi-persistent systems (KS 1731--260) suddenly turned off after having actively accreted for over 12.5 years. A {\\it Chandra} observation taken a few months after this transition showed the source at a 0.5--10 keV luminosity of $\\sim10^{33}$ \\Lunit~(Wijnands et al. 2001b). An {\\it XMM-Newton} observation of this system performed about half a year after the {\\it Chandra} observation, showed that the system had declined by a factor of $\\sim$3 in luminosity (Wijnands et al. 2002b).  If the quiescent emission from this system was dominated by the state of the crust, the decrease in luminosity within half a year strongly indicates that the crust must have a high heat conductivity (Wijnands et al. 2002b; using the cooling curves calculated for this system by Rutledge et al. 2002). In this scenario, the core temperature is expected to be lower than the crust temperature and the luminosity measured with {\\it XMM-Newton} can be used as an upper limit on the core luminosity. The quiescent luminosity of KS 1731--260 is much lower than expected from its long term accretion history, and can only be explained in terms of the standard cooling model if this system is dormant for at least several thousand of years between outbursts (assuming all outbursts of this system are very similar to the last one, which might not be a valid assumption; Wijnands et al. 2001b; Rutledge et al. 2002). Alternatively, enhanced cooling processes might be active in the core, rapidly cooling it (e.g., Wijnands et al. 2001b, 2002b).  \\subsection {MXB 1659--29}  \\begin{figure}[t] \\begin{center} \\begin{tabular}{c} \\psfig{figure=f1.eps,width=7cm} \\end{tabular} \\figcaption{The {\\it RXTE}/ASM light curve of MXB 1659--29 clearly showing the 1999--2001 outburst. The time of our {\\it Chandra} observation is indicated by the solid line. Note that our {\\it Chandra} observation was performed at times that the {\\it RXTE}/ASM (and also the {\\it RXTE}/PCA) could not detect the source anymore (see text). \\label{fig:asm} } \\end{center} \\end{figure}   In 2001 September, the opportunity arose to use another quasi-persistent system to test the cooling neutron star model: MXB 1659--29.  This source is an X-ray transient and was discovered in 1976 by Lewin, Hoffman, \\& Doty (1976) during type-I X-ray bursts, which clearly demonstrates that the compact object in this system is a neutron star. The source was detected several times between 1976 October and 1978 September with {\\it SAS3} and {\\it HEAO} (Lewin et al. 1978; Share et al. 1978; Griffiths et al. 1978; Cominsky, Ossman, \\& Lewin 1983; Cominsky \\& Wood 1984, 1989) and irregular X-ray variability was found in this system (Lewin 1979; Cominsky et al. 1983). Cominsky \\& Wood (1984, 1989) reported on the discovery of eclipses every $\\sim$7.1 hours, which can be identified with the orbital period of the system.  MXB 1659--29 is one of only several non-pulsating neutron star low-mass X-ray binaries for which total eclipses have been observed (The other confirmed eclipsing neutron star systems are EXO 0748--676, GRS 1747--312, and AX J1745.6--2901, although several other systems [e.g., 4U 2129+47] show partial eclipses; Parmar et al. 1986; In 't Zand et al. 2000; Maeda et al. 1996; McClintock et al. 1982). During later observations, using a variety of satellites (e.g., {\\it Hakucho}, {\\it EXOSAT}, {\\it ROSAT}), the source could not be detected anymore (Cominsky et al. 1983; Verbunt 2001). The pointed {\\it ROSAT} observations in the early 1990's failed to detected the source with a 0.5--10 keV upper limit on the unabsorbed flux of $(1 - 2) \\times 10^{-14}$ \\funit~(Verbunt 2001; Wijnands 2002; Oosterbroek et al. 2001).  The source remained dormant until 1999 April, when in 't Zand et al. (1999) reported it to be active again in observations obtained with the {\\it BeppoSAX} Wide Field Camera. Wachter, Smale, \\& Bailyn (2000) and Oosterbroek et al. (2001) obtained an updated ephemeris for the orbital period (using data obtained with the {\\it Rossi X-ray Timing Explorer} [{\\it RXTE}] and {\\it BeppoSAX}). Studies of its X-ray spectrum were performed using {\\it BeppoSAX} (Oosterbroek et al. 2001) and {\\it XMM-Newton} (Sidoli et al. 2001), and, using {\\it RXTE} data, Wijnands, Strohmayer, \\& Franco (2001a) found $\\sim$567 Hz oscillations during X-ray bursts. Those oscillations are likely related to the neutron star spin frequency. The source remained bright for almost 2.5 years before it became dormant again in 2001 September (Wijnands et al. 2002a).  Because of its long outburst duration, MXB 1659--29 may be classified as a quasi-persistent source. We had a Cycle 3 {\\it Chandra} TOO proposal approved to obtain a quiescent observation of the next quasi-persistent source that could turn off, within a month after the transition. As part of this proposal, MXB 1659--29 was observed on 2001 October 15 using {\\it Chandra} for $\\sim19$ ksec.  Here we report on this observation. ",
        "conclusions": "}  We have presented a {\\it Chandra} observation of MXB 1659--29 performed $\\sim$5 weeks after the last clear detection of the source with the {\\it RXTE}/PCA (indicating that at that time the source was still actively accreting; Wijnands et al. 2002a). During our {\\it Chandra} observation, we detected the source at a luminosity of $(3.2 - 4.3) \\times10^{33} (d/10\\,\\,\\, {\\rm kpc})^2$ \\Lunit, and its spectrum could be well described by a thermal component (either a blackbody model or a neutron star atmosphere model).  The obtained luminosity and the shape of the X-ray spectrum of MXB 1659--29 resemble those obtained for other quiescent neutron star systems, strongly suggesting that the source was quiescent during our {\\it Chandra} observation.  \\subsection{The cooling neutron star model}  As argued by Rutledge et al. (2002), for systems which are actively accreting for long periods, the crust might have been heated to very high temperatures (probably considerably higher than that of the core) and the quiescent properties might be dominated by the thermal state of the crust rather than that of the core. Although the duration of the accretion episode of MXB 1659--29 is considerably shorter (factor of $\\sim$5) than that of KS 1731--260 or X 1732--304 (both of which had accretion episodes of more than a decade; Wijnands et al. 2001b; Wijnands et al. 2002b; see also Guainazzi, Parmar, \\& Oosterbroek 1999 for X 1732--304), here we assume that the 2.5 year accretion episode of MXB 1659--29 has had a considerable effect on the state of the crust, similar to what has been argued for KS 1731--260 (Rutledge et al. 2002; although likely less extreme), .  This assumption is also supported by the fact that in the early 1990's, the source could not be detected in quiescence using a {\\it ROSAT} observation and only a 0.5--10 keV upper limit of $(1 - 2) \\times 10^{32} (d/10\\,\\,\\, {\\rm kpc})^2$ \\Lunit~ could be obtained (Verbunt 2001; Wijnands 2002; Oosterbroek et al. 2001) which is about an order of magnitude lower than the luminosity we have detected during our {\\it Chandra} observation. Therefore, we can conclude that the {\\it Chandra} quiescent luminosity is not the rock bottom quiescent luminosity of this system and the quiescent properties for MXB 1659--29 during the {\\it Chandra} observation were dominated by that of the crust and not by the core. As an upper limit on the flux due to the cooling neutron star core, we will assume the upper limit provided by {\\it ROSAT} and this upper limit will be used to test the cooling neutron star model.  In order to test the model, the time averaged accretion rate has to be estimated. The last outburst was fully covered with the {\\it RXTE}/ASM instrument (Fig.~\\ref{fig:asm}) and lasted for $\\sim2.5$ years. The 2--10 keV luminosity during this state as obtained with {\\it BeppoSAX} and {\\it XMM-Newton} was about $6\\times 10^{-10}$ \\funit~(Oosterbroek et al. 2001; Sidoli et al. 2001). Using the spectral model and the spectral parameters given by Oosterbroek et al. (2001), we inferred (by simulating the spectrum in XSPEC) that the bolometric luminosity can be at least a factor of two larger. However, the luminosity itself was variable (by a factor of a few; Fig.~\\ref{fig:asm}) during the outburst, and large uncertainties might be present in the 'typical' outburst bolometric luminosity. However, in the rest of the discussion we assume that bolometric outburst flux was typically $(5 -10) \\times 10^{-10}$ \\funit~during the last outburst.  The 1999--2001 outburst could be an atypical one for MXB 1659--29 and previous outbursts might have been less bright and/or less long (i.e., more like those of the ordinary transients). We have searched the literature for reports on detections of this source in the past and we found that the source was conclusively detected in X-rays in 1976 October, 1977 June, July\\footnote{During 1977 June and July, the observations were not very sensitive to persistent X-ray emission and none was detected (e.g., Cominsky \\& Wood 1984, 1989). But X-ray bursts were observed from the source, indicating that also during those observations, the source was accreting, albeit at a low level.}, and September, and 1978 March and September using {\\it SAS3} and {\\it HEAO} (Lewin et al. 1976; Lewin et al. 1978; Share et al. 1978; Griffiths et al. 1978; Lewin 1979; Cominsky et al. 1983; Cominsky \\& Wood 1984, 1989) and in the optical on 1978 June 1 and 1979 June 27 -- July 2 (Doxsey et al. 1979; Cominsky et al. 1983). No information is available during the periods in-between those observations. Although it cannot be excluded that at those occasions, the source was in quiescence, we consider it unlikely that the source would only be active during times when it was observed with an X-ray or optical instrument and dormant when no instrument looked at the source. Therefore, it is likely that during the complete period from 1976 October until early July 1979 the source was actively accreting for over 2.5 years, especially because the recent outburst had a similar duration. If true, then this would constitute the first indication that different outbursts of quasi-persistent sources may have similar durations and that the long duration of those outbursts might be a common property of those sources.  The first reported non-detection of the source was on 1979 July 17--25 (Cominsky et al. 1983) in optical (V$>$22--23) and with {\\it Hakucho} (no X-ray upper limits were provided). Prior to 1976, the source might have also been detected during the period 1971 to 1973 using {\\it Uhuru} (classified as 4U 1704--30; Forman et al. 1978), although this identification with MXB 1659--29 is not certain and we will assume that they are two different objects (if 4U 1704--30 can be identified with MXB 1659--29 then the source might have exhibited an extra outburst during that period or it might have been active for a period of $\\sim$7 years).  The exact fluxes during the observations in the period October 1976 to early July 1979 are in the range 1 to 6 $\\times 10^{-10}$ \\funit, but they have large uncertainties because the exact energy range was not always quoted (if quoted it was 1--10 keV or 2--10 keV), it was unclear if the fluxes were absorbed or unabsorbed, and the assumed spectral shape was not always similar (often assumed to be Crab like) and quite different than observed with {\\it BeppoSAX} and {\\it XMM-Newton} during the 1999--2001 outburst (Oosterbroek et al. 2001; Sidoli et al. 2001). However, we will assume that all fluxes are for the 2--10 keV range and unabsorbed, and a bolometric flux of about twice the quoted values (as inferred above for the {\\it BeppoSAX} results on MXB 1659--29). Therefore, during the outburst in the late 1970's, the source was actively accreting for a period of at least $\\sim$2.5 years at a bolometric flux level of 2 to 12 $\\times 10^{10}$ \\funit. Although quite uncertain, this is remarkably similar to the values of the 1999--2001 outburst and for simplicity we assume that the typical outburst duration is 2.5 years and that the bolometric fluxes during outburst is 5--10 $\\times 10^{-10}$ \\funit.  Using the Brown et al. (1998) model (assuming standard cooling processes), the predicted quiescent flux $ F_{\\rm q}$ for this source would then be (Wijnands et al. 2001b; see also Rutledge et al. 2002) $ F_{\\rm q} \\approx {t_{\\rm o} \\over t_{\\rm o} + t_{\\rm q}} \\times {\\langle F_{\\rm o} \\rangle \\over 135}$, with $\\langle F_{\\rm o} \\rangle$ the average flux during outburst (5--10 $\\times 10^{-10}$ \\funit), $t_{\\rm o}$ the average time the source is in outburst (2.5 years), and $t_{\\rm q}$ the average time the source is in quiescence ($\\sim$21 year). This results in a predicted quiescent flux of 4--8 $\\times 10^{-13}$ \\funit. Remarkably, this value is very similar to the quiescent flux observed during our {\\it Chandra} observations. However, as explained above, based on the {\\it ROSAT} non-detection of the source, the core flux is likely at least an order of magnitude lower than this, which would make the predicted core flux considerably higher than that truly originating from the core. Similar to KS 1731--260 (Wijnands et al. 2001b, 2002b) this low core flux (and thus temperature) might be due to enhanced core cooling instead of the assumed standard core cooling in the Brown et al. (1998) model.  Despite the fact that the last two outburst episodes are likely of similar duration, it cannot be excluded that they are not typical for the source and that most of the time MXB 1659--29 exhibits short duration outbursts. If the typical outburst duration of MXB 1659--29 is not 2.5 years but instead 0.25 years (3 months) or shorter, then the predicted quiescent flux will be consistent (within the uncertainties of the model and assumptions) with the {\\it ROSAT} upper limit. The higher {\\it Chandra} quiescent luminosity is again due to the state of the crust which should be considerably heated during the long accretion episode.  \\subsection{Crust cooling}  If during the {\\it Chandra} observation the X-ray emission was dominated by thermal emission from the crust, then further quiescent observations of MXB 1659--29 will enable studies of the cooling of the crust in this system. The {\\it ROSAT} flux upper limit suggests that the crust flux should eventually decrease to at least this level. When the crust will be thermally relaxed with the core, no significant further decrease of the quiescent flux is expected and from this bottom flux level the state of the core can be inferred from which the cooling models can be better constrained. For KS 1731--260 it had already been found that its quiescent luminosity decreased by a factor of 3 within half a year time (likely due to a temperature decrease), indicating a highly conductive crust (Wijnands et al. 2002b). It would be of interest to determine if such a rapid cooling will also be observed for MXB 1659--29 or if the neutron star crust in this system has a significantly lower conductivity.  In the latter case, the system should be at the {\\it Chandra} quiescent luminosity for several years to decades. Using the fact that 13 years after the 1976--1978 outburst {\\it ROSAT} observed an order of magnitude lower quiescent luminosity we might set already an upper limit on the crust cooling time.  When assuming that shortly after the end of the 1976--1978 outburst the quiescent luminosity was similar to our measured {\\it Chandra} luminosity, the cooling time of the crust is at least about a factor of 10 in luminosity per decade. Due to differences in quiescent times, outburst times, and the time averaged accretion rates between MXB 1659--29 and KS 1731--260 the cooling curves calculated for KS 1731--260 by Rutledge et al. (2002) cannot be used for MXB 1659--29. However, if those MXB 1659--29 cooling curves resemble those of KS 1731--260, then tentatively it might be concluded that also the neutron star crust in MXB 1659--29 has a high conductivity (and enhanced core cooling is suggested). We await specifically calculated cooling curves for MXB 1659--29 and further monitoring observations using {\\it Chandra} or {\\it XMM-Newton} in order to be conclusive about the properties of the neutron star in this system.  \\subsection{Very low thermal emission in quiescent neutron star systems?}  For most quiescent neutron star X-ray transients it has been inferred that the thermal emission is in the range of a few times $10^{32-33}$ \\Lunit. However, recently, indications have been found (using {\\it XMM-Newton} data) that the thermal emission from the accretion driven millisecond X-ray pulsar SAX J1808.4--3658 in its quiescent state might be as low as a few times $10^{30}$ \\Lunit~(Campana et al. 2002; who suggested enhanced core cooling for this low thermal luminosity). Although the statistics of those results were not overwhelming and have to be confirmed with additional observations, it is an interesting possibility and if the processes which produce the quiescent emission (which has a power-law shape; Campana et al. 2002) for this system would become inactive, SAX J1808.4--3658 would become rather dim in quiescence.  Such weak quiescent neutron star systems might also be suggested by the indications found for enhanced cooling processes in certain systems, such as MXB 1659--29. The possibility of dim quiescent neutron star systems raises the question of how dim certain neutron star systems can become?  The answers to this question will have implications for our understanding of quiescent X-ray binaries. It has been found that those X-ray binaries which harbor a black hole instead of a neutron star can be at least one to two orders of magnitude less luminous in quiescence than the average neutron star system. This difference has been used as evidence that the black holes have event horizons (e.g., Garcia et al. 2001 and references therein) in contrast to the surfaces of neutron stars. However, if certain neutron star systems might become similarly dim (e.g., due to enhanced core cooling or very dim outbursts), than this luminosity difference will disappear and with it the evidence for event horizons in black hole systems. The quiescent spectrum of SAX J1808.4--3658 (Campana et al. 2002) indicates that the spectrum of such systems might not be dominated by a thermal component but might have a power-law shape, similar to what has seen for the black hole systems (Kong et al. 2002) removing also the spectral differences.  Another area in which it might be important to determine the full range of the luminosity distribution of quiescent neutron star systems, is that of the study of the low-luminosity X-ray sources in globular clusters. Based on their luminosities and their X-ray spectra (when enough statistics are available), those sources which have luminosities above a few times $10^{32}$ \\Lunit, have been classified as quiescent neutron star systems and those which have luminosities below $10^{32}$ \\Lunit, as another type of object (possible cataclysmic variables or millisecond radio pulsars). However, the observed low luminosity of $5 \\times 10^{31}$ \\Lunit~ of SAX J1808.4--3658 (Dotani, Asai, \\& Wijnands 2000; Campana et al. 2002) already shows (irrespective of what exactly causes the X-rays in this source) that those dim sources might be quiescent neutron star transients. The fact that the spectrum of SAX J1808.4--3658 appears to be considerably harder (Campana et al. 2002) than the average quiescent neutron star spectrum indicates that classifications on hardness ratio might lead to erroneous results. The possibility that the thermal emission of SAX J1808.4--3658 might be very low even suggests that those sources which have luminosities down to only a few times $10^{30}$ \\Lunit~might also be quiescent neutron star systems. Therefore, we conclude that classifying low-luminosity globular cluster sources based on their luminosity and broad band spectral shape (i.e., hardness ratio's) might possibly lead to misleading results. The full extent of those errors depends on the full luminosity distribution of the quiescent neutron star transients and how the spectrum correlates with the luminosity. Those properties have to be determined before a complete picture of the nature of the low-luminosity globular cluster sources can be understood\\footnote{Although we have only discussed quiescent neutron star systems, similar uncertainties in the luminosity distribution of CVs or millisecond radio pulsars exist. Those uncertainties in the properties of those systems have to be resolved before a full understanding of the low-luminosity globular cluster source population can be obtained.}."
    },
    "0207/astro-ph0207577_arXiv.txt": {
        "abstract": "{We report the first detection of doubly-deuterated methanol (CHD$_2$OH), as well as firm detections of the two singly-deuterated isotopomers of methanol (CH$_2$DOH and CH$_3$OD), towards the solar-type protostar IRAS16293$-$2422.  From the present multifrequency observations, we derive the following abundance ratios: [CHD$_2$OH]/[CH$_3$OH] = $0.2 \\pm 0.1$, [CH$_2$DOH]/[CH$_3$OH] = $0.9 \\pm 0.3$, [CH$_3$OD]/[CH$_3$OH] = $0.04 \\pm 0.02$.  The total abundance of the deuterated forms of methanol is greater than that of its normal hydrogenated counterpart in the circumstellar material of IRAS16293$-$2422, a circumstance not previously encountered. Formaldehyde, which is thought to be the chemical precursor of methanol, possesses a much lower fraction of deuterated isotopomers ($\\sim 20$\\%) with respect to the main isotopic form in IRAS16293$-$2422. The observed fractionation of methanol and formaldehyde provides a severe challenge to both gas-phase and grain-surface models of deuteration.  Two examples of the latter model are roughly in agreement with our observations of CHD$_2$OH and CH$_2$DOH if the accreting gas has a large (0.2-0.3) atomic D/H ratio.  However, no gas-phase model predicts such a high atomic D/H ratio, and hence some key ingredient seems to be missing. ",
        "introduction": "In the last few years, the study of doubly-deuterated molecules in the interstellar medium has gained considerable attention. This field was boosted by the discovery of an extremely large amount (D$_2$CO/H$_2$CO $\\sim$ 10\\%) of doubly-deuterated formaldehyde in the low mass protostar IRAS16293$-$2422 (hereafter IRAS16293 ; \\cite{Ceccarelli98} 1998), a fractionation about 25 times larger than in Orion (\\cite{Turner90} 1990). Follow-up observational studies of this first discovery confirmed this very large degree of deuteration in IRAS16293 (\\cite{Loinard00} 2000), and allowed a study of its spatial distribution (\\cite{Ceccarelli01} 2001). Subsequently, similarly large amounts of doubly-deuterated formaldehyde and ammonia have been observed towards another very young protostellar core, 16293E, which lies in the same molecular cloud (L1689N) as IRAS16293 (\\cite{Loinard01} 2001) and in the molecular cloud L1689N itself (\\cite{Ceccarelli02} 2002).  Finally, preliminary results of an ongoing project show that it is likely that {\\it all} low-mass protostars present similarly large abundance ratios of doubly deuterated formaldehyde with respect to H$_2$CO, whereas high-mass protostars do not (\\cite{Loinard02a} 2002, \\cite{Ceccarelli02} 2002).  All these observations suggest that such a large deuteration of formaldehyde is produced during the cold and dense pre-collapse phase of low-mass protostars.  Highly deuterated ices are very likely formed via active grain chemistry (\\cite{Tielens83} 1983), stored on the grain mantles, and eventually released into the gas during the collapse phase, when the heating of newly-formed protostars evaporates the CO-rich ices (\\cite{Ceccarelli01} 2001a). Methanol is also commonly believed to be formed on grain surfaces, because gas-phase models cannot account for the large detected abundances of methanol in hot cores (\\cite{Menten88} 1988).  If formaldehyde and methanol are produced on grain surfaces by simple successive hydrogenations of CO, then the reproduction of the abundance ratios between deuterated isotopomers and their normal counterparts is a crucial test for the grain-surface theory of deuteration (e.g. \\cite{Charnley97} 1997).  Although the grain picture seems qualitatively consistent with all the observations so far available towards protostars, the nature of the production of deuterated molecular species,  whether it occurs completely via active grain chemistry (starting from a high D/H atomic ratio derived from gas-phase chemistry) or at least partially via gas-phase formation, is still largely debated (see for example \\cite{Roberts00b} 2000b).  The debate has not been settled conclusively because of the relatively small body of available observations and the discovery of relatively large abundances of doubly-deuterated ammonia (NHD$_2$/NH$_3 \\sim$ 0.001; \\cite{Roueff00} 2000) in the molecular cloud L134N and triply-deuterated ammonia in the low-mass protostar NGC1333-IRAS4 (\\cite{vanderTak02} 2002) and in the dark cloud B1 (\\cite{Lis02} 2002).  The observed fractionation of ammonia in L134N can be accounted for by gas-phase models if a high degree of depletion of heavy materials onto the grain mantles is assumed (\\cite{Roberts00b} 2000b, \\cite{Rodgers01} 2001).  The ND$_{3}$ observations from \\cite{Lis02} (2002) can also be explained in the framework of gas-phase chemical models if the dissociative recombination of partially deuterated ions results in a somewhat higher probability for the ejection of hydrogen atoms than for deuterium atoms.  We wish to emphasize that this debate is not merely academic, as it involves our understanding of the chemistry of the interstellar medium and of ice formation in general and deuteration processes in particular.  Many observational studies use deuteration processes, which are supposedly well-understood, to derive key quantities such as the deuterium abundance (e.g.  in the Galactic Center ; \\cite{Lubowich00} 2000) or the degree of ionization (e.g. in protostars ; \\cite{Williams98} 1998).  The actual state of our comprehension of those processes has therefore a large impact.  In this Letter, we report the very first detection of a doubly-deuterated isotopomer of methanol (CHD$_2$OH), with 15 detected lines, towards the low-mass protostar IRAS16293$-$2422.  We also report the detection of the two singly-deuterated forms (CH$_2$DOH and CH$_3$OD) of methanol towards the same object.  We compare the derived fractionation ratios as well as the formaldehyde fractionation (\\cite{Loinard00} 2000) with predictions based on active grain chemistry. ",
        "conclusions": "Our most dramatic result is the detection in IRAS16293 of a form of doubly-deuterated methanol along with the detection of both possible singly-deuterated isotopomers of this molecule.  Up to now, only a tentative detection of CH$_3$OD has been reported in a low-mass protostellar source (\\cite{vanDishoeck95} 1995).  Singly-deuterated methanol has been definitely observed towards Orion (\\cite{Mauersberger88} 1988, where CH$_3$OD/CH$_3$OH $\\sim$ 0.01-0.06 and \\cite{Jacq93} 1993, where CH$_{2}$DOH/CH$_3$OH $\\sim$ 0.04) and SgB2 (\\cite{Gottlieb79} 1979, with CH$_3$OD/CH$_3$OH $\\sim$ 0.01). Equally strikingly, we find the deuterated forms of methanol to possess a total abundance greater than the main isotopomer in IRAS16293, even without the contribution of the doubly-deuterated isotopomer CH$_{2}$DOD!  To date, no other molecule has been observed to show such extreme deuterium fractionation.  As discussed in the Introduction, the abundances of deuterated methanol and deuterated formaldehyde provide a strong test of models involving active grain chemistry.  The basic hypothesis behind these models is that formaldehyde and methanol form by the hydrogenation of CO accreted onto the grains via reactions with atomic hydrogen (\\cite{Tielens82} 1982; \\cite{Charnley97} 1997, hereafter CTR97).  The enhanced deuteration is caused by an enhanced (atomic) D/H ratio in the gas during the era of mantle formation (\\cite{Tielens83} 1983).  The hydrogenation and deuteration of CO is predicted to form H$_2$CO first (CO $\\rightarrow$ HCO $\\rightarrow$ H$_{2}$CO) and subsequently CH$_3$OH (H$_{2}$CO $\\rightarrow$ H$_{3}$CO $\\rightarrow$ CH$_{3}$OH). With some simplifying assumptions, this leads directly to predictions for steady-state ratios of singly- and doubly-deuterated formaldehyde and methanol to their normal isotopic forms in terms of the relative accretion rates of H and D with respect to CO as free parameters (CTR97).  The relative accretion rate of H with respect to CO can be derived from the observation of the CO/CH$_3$OH and H$_2$CO/CH$_3$OH abundance ratios.  The predictions for the fractionation ratios then depend only on the relative accretion rate of D with respect to CO, or equivalently on the D/H atomic abundance ratio in the accreting gas. Figure \\ref{ratios} shows fractionation ratios predicted by the CTR97 model.  In particular, the calculated ratios of singly-deuterated isotopomers to normal species are plotted against the analogous ratios for doubly-deuterated species, both as functions of the D/H atomic ratio in the gas. As seen in the upper panel, the CH$_{2}$DOH and CHD$_{2}$OH observations are compatible with an atomic D/H ratio of 0.2 in the accreting gas. The CH$_3$OD abundance falls short of the predicted value for a D/H ratio of 0.2.  However, the gas phase abundance of CH$_3$OD in the hot core can be affected by gas phase ion-molecule reactions.  Specifically, protonation of methanol by H$_3^+$ or H$_3$O$^+$ followed by dissociative electron recombination back to methanol will drive the CH$_3$OD/CH$_3$OH ratio to the deuterium fractionation of molecules in the warm gas, which is very low (CTR97).  The timescale for this process is some 3$\\times$10$^4$ yrs which is comparable to the lifetime of IRAS16293 ($\\sim$ 2$\\times$10$^4$, \\cite{Ceccarelli00} 2000). We note that this chemical reshuffling of the deuterium will not affect the deuterium fractionation on the methyl group (eg., CH$_2$DOH, CHD$_2$OH). The CTR97 model, however, has some difficulties in explaining the observed formaldehyde fractionation ratios towards IRAS16293 (\\cite{Loinard00} 2000). In particular, the value of atomic D/H compatible with CH$_{2}$DOH and CHD$_{2}$OH is reasonably compatible with HDCO but results in too low a fractionation ratio for D$_{2}$CO by a factor of 5 or so. Possible gas-phase alterations have not been considered. \\begin{figure} \\includegraphics[width=9cm]{ratios_charn.eps} \\caption{Model predictions (solid lines, adapted from CTR97) of abundance ratios between singly-deuterated and doubly-deuterated isotopomers with respect to normal species are plotted vs one another for methanol and formaldehyde.  The predictions are obtained as functions of the gas-phase atomic D/H ratio.  The observed ratios of \\cite{Loinard00} (2000) are also shown. Upper panel: CH$_{2}$DOH/CH$_3$OH (diamond) and CH$_3$OD/CH$_3$OH (triangle) versus CHD$_{2}$OH/CH$_3$OH. Lower panel : HDCO/H$_2$CO versus D$_2$CO/H$_2$CO. } \\label{ratios} \\end{figure} The formaldehyde discrepancy suggests that at least some of the assumptions in the CTR97 model may be too drastic. For example, the CTR97 model contains the approximation that only the accreted H, D and CO are important in the formation of formaldehyde and methanol on the grain surfaces, and that no other reactions compete with formaldehyde and methanol formation.  Very recently, \\cite{Caselli02} (2002) [hereafter CSS02] proposed a somewhat more detailed but related model for the formation of formaldehyde, methanol, all of their deuterated isotopomers, and selected other species on grains. In their model, CSS02 consider accretion onto grains of H, D, CO and O, followed by a comprehensive set of reactions to form H$_2$O, H$_2$, CO$_2$, H$_2$CO, CH$_3$OH and all singly- and multiply-deuterated isotopomers of these species. Their model differs slightly from that of CTR97 in that it specifically includes small differences in the barriers to reaction between non-deuterated and deuterated species. Yet, in the so-called accretion limit and with the same assumption of rapid diffusion rates, the CSS02 model should yield approximately the same results as the CTR97 model given the same set of chemical reactions and physical conditions. Predictions for methanol and formaldehyde fractionation ratios are indeed similar to the CTR97 model and substantially the same discrepancies remain.  In particular, for D/H = 0.3, a temperature of 10 K, and so-called high-density conditions, CSS02 agree approximately with our observed fractionation ratios for CH$_{2}$DOH and CHD$_{2}$OH, but produce approximately 5 times too much CH$_{3}$OD, a value similar to that of CTR97 with D/H = 0.2 and no subsequent gas-phase chemistry.  For formaldehyde, CSS02 obtain a fractionation ratio for HDCO that is twice the observed value and a fractionation ratio for D$_{2}$CO that is a factor of two below the observed value.  In summary, neither model is in good agreement with all of our data. More importantly perhaps, the CTR97 and CSS02 models require a D/H atomic ratio in the range 0.2-0.3, which is a significantly larger value than can be produced by current gas-phase models, even in the presence of a large CO depletion (e.g. \\cite{Roberts00a} 2000a). Future progress will probably require more detailed chemical models in which  gas-phase and surface chemistry occur simultaneously. Although such models are currently in existence, they do not yet contain fractionation processes."
    },
    "0207/astro-ph0207431_arXiv.txt": {
        "abstract": "{\\it Chandra} observations of the core of the nearby starburst galaxy NGC~253 reveal a heavily absorbed source of hard X-rays embedded within the nuclear starburst region.   The source has an unabsorbed, 2 to 10 keV luminosity of $\\ge10^{39}$ erg s$^{-1}$ and photoionizes the surrounding gas. We observe this source through a dusty torus with a neutral absorbing column density of $N_{\\rm H}\\sim2\\times10^{23}$ cm$^{-2}$.  The torus is hundreds of pc across and collimates the starburst-driven nuclear outflow. We suggest that the ionizing source is an intermediate-mass black hole or a weakly accreting supermassive black hole, which may signal the beginnings or endings of AGN activity. ",
        "introduction": "In recent years, there has been increasing speculation about the connection between circumnuclear starbursts and active galactic nuclei (AGN).   Such speculation is due mostly to the fact that, as our instruments allow us to probe closer to the cores of nearby galaxies, we find that an an increasing number contain starbursts and AGN in close proximity (Levenson, Weaver and Heckman 2001, and references therein).  It is not clear whether proximity implies a physical connection, but a circumnuclear starburst could easily provide a pathway toward forming a supermassive black hole (and subsequent AGN), since it can processes as much as $\\sim10^{10}$ M$_{\\sun}$ of material in $10^7 - 10^8$ years (Norman and Scoville 1988). To date, however, there has been little direct evidence for this scenario.  X-rays can penetrate the dense cores of nearby galaxies and are thus crucial for probing the possible links between starburst and AGN activity. For this purpose, we have obtained {\\it Chandra} observations of the nearby ($\\sim2.6$ Mpc) starburst galaxy NGC 253. This galaxy possesses a strong circumnuclear starburst (Strickland et al. 2000) and evidence for a weak AGN (Turner \\& Ho 1995, Mohan, Anantharamaiah \\& Goss 2002). {\\it Chandra}, with its resolution of $\\sim1^{\\prime\\prime}$ ($\\sim12$ pc at NGC~253) allows us to untangle the X-ray emission processes due to stellar and non-stellar activity for the first time at the core of the galaxy. ",
        "conclusions": "{\\it Chandra} X-ray observations of the nearby starburst galaxy NGC~253 reveal what may be the beginnings or endings of AGN activity. The excellent spatial resolution allows us to isolate the optically-thick torus that collimates the starburst-driven nuclear outflow. At the center of the torus, along with the evolved, circumnuclear starburst, is a source of hard X-rays with an unabsorbed, 2 to 10 keV luminosity of $\\ge10^{39}$ erg s$^{-1}$. We suggest that this ionizing source is an intermediate-mass black hole or weakly accreting supermassive black hole. These data provide a unique look at the complex interplay between starburst and AGN activity.  Future tests of the starburst-AGN scenario will require studying older starburst populations in nearby galaxies.  In particular, it is important to look at normal stars in the optical or at X-ray binaries with {\\it Chandra} in the more evolved, more AGN-like composite (Seyfert/starburst) galaxies.   \\bigskip"
    },
    "0207/astro-ph0207560.txt": {
        "abstract": "Taylor--Couette flow in the presence of a magnetic field is a problem belonging to classical hydromagnetics and deserves to be more widely studied than it has been to date. %Literature in the nonlinear regime is scarce. In the nonlinear regime the literature is scarce. We develop a formulation %based on spectral methods, suitable for solution of the full three dimensional nonlinear hydromagnetic equations in cylindrical geometry, which is motived by the formulation for the magnetic field. It is suitable for study at finite Prandtl numbers and in the small Prandtl number limit, relevant to laboratory liquid metals. The method is used to determine the onset of axisymmetric Taylor vortices, and finite amplitude solutions. Our results compare well with existing linear and nonlinear hydrodynamic calculations and with hydromagnetic experiments. ",
        "introduction": "The motion of an incompressible viscous fluid between concentric rotating cylinders is one of the most important problems of fluid dynamics and is much studied as a benchmark to investigate issue of instability and nonlinear behaviour. \\cite{taylor23} found that if the rotation of the inner cylinder is greater than some critical value then circular--Couette flow (CCF) becomes unstable to axisymmetric perturbations. A secondary flow appears which has axial and radial motion in the form of pairs of toroidal vortices, now known as the Taylor--vortex flow (TVF). If the inner cylinder is driven further then this flow becomes unstable to non-axisymmetric perturbations.  Azimuthal waves appear in the Taylor--vortices and the whole pattern rotates at some wavespeed (wavy modes).  In his landmark 1961 book on stability theory, Chandrasekhar devoted equal attention to the hydrodynamic and the hydromagnetic Couette problems; the latter is the case in which the fluid is a conducting liquid (\\eg mercury, liquid gallium, liquid sodium) and a magnetic field is applied externally. Despite this early interest in the hydromagnetic Couette problem, which included experiments performed by \\cite{donnelly62} and by \\cite{donnelly63}, most of the activity of the following years was devoted to the hydrodynamic case. Among the few studies of the effects of the magnetic field it is worth remembering the works by \\cite{velikhov59}, \\cite{kurzweg63}, \\cite{roberts64}, who extended Chandrasekhar's theory to non-axisymmetric bifurcations from circular--Couette flow, \\cite{chang67} and \\cite{hunt71} at finite aspect ratio. Later \\cite{tabeling81}, using a method similar to Davey's (1962) amplitude expansion, calculated effective viscosity of axisymmetric flow in the Taylor vortex flow regime; he compared against Donnelly's (1962) experiments which indicate that the onset of wavy vortices is significantly inhibited by the magnetic field. \\cite{nagata96} has more recently investigated nonlinear solutions in the planar geometry, and \\cite{hollerbach00a} shows Taylor cells in spherical geometry.  The aim of this paper is to investigate effects induced on Couette flow by an externally applied magnetic field. This paper is meant to be the first of a series and is dedicated to the development of a suitable formulation for solving numerically the governing nonlinear three dimensional magnetohydrodynamic (MHD) equations in the cylindrical Couette geometry. The numerical method which we propose can be used for any value of radius ratio, is suitable for time stepping, has good stability features, is relatively easy to program and is more accurate than existing methods.  Our work is also motivated by the renewed interest in MHD flows in confined geometries which arises from current and planned experiments to produce dynamo action in the laboratory (Gailitis \\etal 2001; Stieglitz \\& Muller 2001). It must be stressed that our work does not apply directly to the dynamo problem for two reasons. Firstly the cylindrical Couette configuration is a possible geometry for these studies, but it has not been used in experiments yet, and it may prove not to be not the most efficient one (Laure, Chossat \\& Daviaud, 2000). Secondly, our investigation refers primarily to bifurcations at relatively small Reynolds numbers, while the dynamo experiments require rather large Reynolds numbers due to the small magnetic Prandtl number of liquid metals. Despite these two limitations, our work is related to the dynamo problem because it is important to have precise results at small Reynolds number MHD flows in order to develop and test modern acoustic flow visualisation techniques, \\cite{kikura99}, which offer the best chance to detect flow patterns in MHD dynamos. The lack of flow visualisation has clearly held back progress in the hydromagnetic Couette problem compared to the hydrodynamic case.   % ------------------------------------------------------------------------ % ------------------------------------------------------------------------ ",
        "conclusions": "In conclusion we have developed a formulation of the governing MHD equations of the cylindrical Couette geometry, suitable for timestepping in the nonlinear regime. Results agree well with experiments.  Although the equations do not decouple in the linear part, and we must treat mean-flows separately, the formulation is similar to that used by \\cite{glatzmaier84} in spherical geometry. We use potentials for the velocity yet do not eliminate the pressure. This has several advantages. Our motivation for adopting such a formulation is that the magnetic field then shares the \\emph{same} formulation as the velocity, dramatically reducing the potential for error. Only a relatively small part of our code is dedicated entirely to the magnetic field; this feature is important for testing, as there are fewer results against which to compare our results. Furthermore, it can also accomodate the small Prandtl number limit with only minor adjustments. The choice of governing equations which are only second order in $r$ makes the method accurate and matrices easily invertible. This feature also enables us to take the same radial truncation for all variables, if we desire, simplifying implementation a great deal.  We have opted to use potentials which ensure divergence-free fields.  Primative variable formulations for time integration of the Navier--Stokes equations, such as \\cite{marcus84} and \\cite{quartapelle95} in this geometry, do not in general extend naturally to the magnetic field.  In particular, they are not well suited to the small Prandtl number limit, relevant to liquid metals available in the laboratory.  For the axisymmetric case the expansion by potentials is essentially the same as that used by \\cite{barenghi91} and \\cite{jones85} and results appear to be very similar in terms of accuracy. Although we have one extra equation, the actual form of our equations is simpler because we avoided taking the second curl.  Our method is second order in time and exhibits good temporal stability, We have not encountered the difficulties experienced by \\cite{hollerbach00b} and \\cite{rudiger00} with three-dimensional potential formulations and no-slip boundaries.  Using the implict Euler method on the linear terms reduces the method to $O(\\dt/\\Rey,\\,\\dt^2)$.  With finite magnetic Prandtl numbers the magneto-rotational instability (R\\\"udiger \\& Zhang, 2001; Willis \\& Barenghi, submitted) %\\cite{rudiger01} leads to Reynolds numbers which can be surprisingly low and the $O(\\dt/\\Rey)$ error would dominate.  Results obtained using our method compare well with existing hydrodynamic literature with respect to the nonlinear equilibration of Taylor-vortex flow (Barenghi 1991), the onset of wavy modes (Jones 1985) and the wavespeed of wavy modes (Marcus 1984). In the presence of a magnetic field the results also compare well for the linear stability of circular--Couette flow (Roberts 1964), and in the nonlinear range the amplitude expansion of \\cite{tabeling81} and experiments of \\cite{donnelly62}.  In further work, we will use this method to analyse nonlinear three-dimensional hydromagnetic Taylor--Couette flow.  % ------------------------------------------------------------------------ % ------------------------------------------------------------------------"
    },
    "0207/astro-ph0207482_arXiv.txt": {
        "abstract": "The discovery of hard X-rays from tops of flaring loops by the HXT of YOHKOH represents a significant progress in the understanding of solar flares.  This report describes the properties of 20 limb flares observed by YOHKOH from October 1991 to August 1998, 15 of which show detectable impulsive looptop emission.  Considering the finite dynamic range (about a decade) of the detection it can be concluded that looptop emission is a common feature of all flares.  The light curves and images of a representative flare are presented and the statistical properties of the footpoint and looptop fluxes and spectral indexes are summarized.  The importance of these observations, and those expected from HESSI with its superior angular, spectral and temporal resolution, in constraining the acceleration models and parameters is discussed briefly. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207161_arXiv.txt": {
        "abstract": "The structure and evolution of central stars of planetary nebulae (CSPNe) is reviewed. CSPNe represent the rapid transitional stage between the Asymptotic Giant Branch (AGB) and the white-dwarf domain. It is shown that the whole evolution off the AGB through the central-star regime depends on the evolutionary history. The detailed evolution into a white dwarf is controlled by the internal stellar structure which, in turn, is determined by the duration of the preceding AGB evolution and therefore by the AGB mass-loss history. The evolution of hydrogen-deficient central stars has been a matter of debate since many years. Convective overshoot appears to be a key ingredient to model these objects. Various thermal-pulse scenarios with inclusion of overshoot are discussed, leading to surface abundances in general agreement with those observed for Wolf-Rayet central stars. ",
        "introduction": "Central stars of planetary nebulae (CSPNe) represent the rapid transitional stage between the Asymptotic Giant Branch (AGB) and the white-dwarf domain. Since the pioneering work of Paczy\\'{n}ski (1971), CSPNe have continued to be in the focus of evolutionary calculations. For instance, Sch\\\"onberner (1979, 1983) demonstrated how the detailed evolution into a central star depends on the previous AGB evolution, i.e.\\ on the thermal-pulse cycle at the tip of the AGB and the shut down of the heavy AGB mass loss. The dependence of the CSPN evolution on the thermal-pulse cycle was investigated in full detail later by Iben (1984), and existing model grid calculations were complemented by the computations of Wood \\& Faulkner (1986). The next stage of calculations included the consideration of appropriate initial-final mass combinations based on  empirical and semi-empirical AGB mass-loss prescriptions (Vassiliadis \\& Wood 1993, 1994; Bl\\\"ocker 1995a,b). Further reviews on the AGB and post-AGB evolution and comparisons of the above calculations are given, e.g., by Iben (1995), Habing (1996), Wood (1997) and Sch\\\"onberner (1997).  While the evolution of hydrogen-rich central stars appears to be adequately understood, the formation of hydrogen-deficient central stars has been a matter of debate for many years.  Convective overshoot turned out to be a prerequisite to model the abundances of objects as the Wolf-Rayet central stars (Herwig et al.\\ 1999, Bl\\\"ocker 2001, Herwig 2001a). ",
        "conclusions": ""
    },
    "0207/astro-ph0207357_arXiv.txt": {
        "abstract": "I will briefly review various analytic approaches to understanding the contents and properties of the hot X-ray emitting gas contained in clusters and groups. Special emphases are given to the following three issues: (1)Reconstruction of the gas distribution in groups and clusters from a (joint) analysis of X-ray, SZ and gravitational lensing observations; (2)Test of the analytic density profiles of dark halos suggested by numerical simulations and empirical models with current X-ray data; And (3)the effects of  preheating and radiative cooling on the X-ray properties of groups and clusters. ",
        "introduction": "Groups and clusters serve as a reservoir of baryons in the present-day universe. They exist in the form of hot plasma with temperature close to the virial temperature ($10^6$-$10^8$ K) of the underlying gravitational potential wells as a result of gravitationally-driven shocks and adiabatic compression. Since the discovery of diffuse X-ray emission associated with clusters of galaxies, many efforts have been made towards exploring the distribution and evolution of the hot gas in clusters. This is significant not only for our understanding of the matter composition and dynamical properties of clusters but also for test of various theories of structure formation. In recent years, a combination of multi-wavelength observations (e.g. optical, X-ray, radio, gravitational lensing, etc.), theoretical analysis and hydrodynamical simulations has greatly improved our knowledge of the intragroup/intracluster gas. In this review I will concentrate on the analytic approaches to determining and modeling the distribution and evolution of the diffuse X-ray emitting gas in groups and clusters. ",
        "conclusions": "We have entered a new era of exploration of matter and energy in groups and clusters, thanks to the high-sensitivity, high-resolution X-ray observations incorporated with optical, radio, SZ and gravitational lensing measurements.  A joint analysis of these independent measurements within next decade will undoubtedly allow us to reconstruct more precisely the distributions of both baryons and dark matter in groups and clusters, which is especially important with respect to the impact on the current debate on the density profiles of dark halos suggested by numerical simulations and empirical models. The physical process of the gas in the formation and evolution of groups and clusters has been another subject of a longstanding debate. The prevailing preheating and radiative cooling scenarios become indistinguishable in the context of current observations, numerical simulations or analytic models. This may indicate that our understanding of physical processes for gas is still incomplete. Clearly caution thus needs to be applied in using groups and clusters for cosmological purpose."
    },
    "0207/astro-ph0207027_arXiv.txt": {
        "abstract": "{ {SN~1993J is to date the radio supernova whose evolution has been monitored in greatest detail and the one which holds best promise for a comprehensive theoretical-observational analysis. The shell-like radio structure of SN~1993J has expanded in general accord with models of shock excited emission, showing almost circular symmetry for over 8 years, except for a bright feature at the south-eastern region of the shell that has been observed at every epoch. The spectrum of SN1993J has flattened from $\\alpha \\simeq -1$ to $\\alpha \\simeq -0.67$ $(S_{\\nu }\\propto \\nu ^{\\alpha })$. The decelerated expansion can be modeled well with a single slope but apparently better with two slopes. There are also intriguing hints of structure in the expansion curve. The results by the two VLBI groups carrying out this research show general agreement, but also some differences. A comparison of the optical and VLBI results about the details of the deceleration show some discrepancies.  } } ",
        "introduction": "Radio emission from supernovae has been successfully modeled in terms of the standard interaction model (SIM; Chevalier \\cite{che82}). This model considers fast-moving supernova ejecta with steep density profiles ($\\rho_\\mathrm{ej} \\sim r^{-n}$) sweeping a circumstellar medium (CSM) of density profile $% \\rho_\\mathrm{csm} \\sim r^{-s}$, resulting in the formation of a high energy-density shell. For $n > 5$, self-similar solutions exist, and the shell radius evolves in time with a power law $R \\sim t^m$, where $t$ is the time since explosion and $m = (n-3)/(n-s)$ is the deceleration parameter. The radio emission is attributed to synchrotron emission from relativistic electrons in the shell, partially suppressed by external free-free absorption from thermal electrons in the CSM.   Due to its proximity and its radio emission level, SN~1993J in M81 has offered an unprecedented occasion for VLBI studies. The harvest of results includes: a) an initial source detection and evolution by Marcaide et al. (\\cite{mar94}) and Bartel et al. (\\cite{bar94}), respectively; b) the discovery of shell-like radio structure (Marcaide et al. \\cite{mar95a}); c) the first ``movie'' of an expanding supernova (Marcaide et al. \\cite{mar95b}); d) determinations of the deceleration in the expansion by Marcaide et al. (\\cite{mar97}) and Bartel et al. (\\cite{bar00}); e) a determination of the center of explosion of SN1993J relative to the quasi-stationary core of M81 (Bietenholz et al. \\cite{bie01}). Results by the two groups carrying out this research show general agreement, but also some differences like the determination of the shell thickness ($\\sim$30\\% of the shell external radius by Marcaide et al. (\\cite{mar95b}); $\\sim$20\\% of the shell external radius by Bartel et al. (\\cite{bar00})). The angular expansion has so far been rather smooth and circular and in accord with the SIM model. However, the supernova shell has displayed, for every epoch and wavelength, an enhancement of emission at its south-eastern part (see, for instance, Fig.~\\ref{fig:sep99-6cm-image} and Fig. ~\\ref{fig:nov00-6cm-image}, corresponding to supernova images in September 1999 and November 2000, respectively, two of our last observing epochs) probably related to the existence of small anisotropies in the density distribution of the CSM. On the other hand, our maps do not show yet any structures or protrusions developing in the shell.  However, a closer look to the previous figures shows also something not uncommon, namely, the changing enhancement of that emission in the south-eastern part relative to other parts of the structure. To be sure: the emission from the south-eastern part is always enhanced, while the emission from other parts of the shell seems to slightly come and go. Bartel et al. (\\cite{bar00}) have even suggested that there may be a cyclic pattern of changes in shell azimuth. We do not have yet evidence of such thing from our data. Even so, there may be at least two possibilities: (a) the changing emission enhancements may be spurious due to, for example, imperfect closure phases or artifacts associated with the CLEAN algorithm, or (b) the changes are real and we should worry to understand them. As said, we do not have evidence from our data that the changes in different parts of the shell are any regular, but perhaps our source sampling has not been appropriate. In any case, the matter has to be systematically addressed with new observations.   To further complicate the picture, we show in Fig.~\\ref{fig:18cm-image} a  preliminary 18 cm image from November 2000. As expected, the shell is not yet clearly delineated at this wavelength. However, is the emission enhancement location well determined by the closure phases?  \\begin{figure}[htbp] \\vspace{240pt} \\special{psfile=\"marcaide_fig1.ps\" hoffset=-25 voffset=-60 hscale=45 vscale=45 angle=0} \\caption{ 6-cm global VLBI image of SN1993J in M81 corresponding to epoch 2369 days after the explosion. The contours are spaced by  $2^{1/2}$ factors, from a lower level of 160 \\,$\\mu$Jy\\,beam$^{-1}$ to a brightness peak of 1939\\,$\\mu$Jy\\,beam$^{-1}$. The convolving beam is circular, with FWHM diameter of 1.8\\,mas. The shell is still clearly defined, and though its shape is not perfectly circular, it does not show any evidence for strong asymmetries.} \\label{fig:sep99-6cm-image} \\end{figure}  \\begin{figure}[htbp] \\vspace{240pt} \\special{psfile=\"marcaide_fig2.ps\" hoffset=-25 voffset=-60 hscale=45 vscale=45 angle=0} \\caption{6-cm global VLBI image of SN1993J in M81 corresponding to epoch 2798 days after the explosion. The contours are spaced by  $2^{1/2}$ factors, from a lower level of 47 \\,$\\mu$Jy\\,beam$^{-1}$ to a brightness peak of 1327\\,$\\mu$Jy\\,beam$^{-1}$. The convolving beam is circular, with FWHM diameter of 1.8\\,mas. The emission enhancements appear somewhat different to Fig.~\\ref{fig:sep99-6cm-image}} \\label{fig:nov00-6cm-image} \\end{figure} ",
        "conclusions": "Over the coming years, we will try to carry on the VLBI monitoring at 6 cm until the source becomes undetectable ($\\sim 5 $ years), and simultaneously observe at 18 cm. With these observations, we will carefully monitor the deceleration, shell width, shell brightness, and spectral index distribution. We will be able to discern any possible dependence of those relevant parameters with frequency and time. Such essential information will constitute the input to our numerical simulation code (P\\'{e}rez-Torres et al. \\cite{per01}) and hence it will strongly influence our ability to characterize the physics involved. Also, our monitoring of the structure of SN1993J should help to understand how the spectral index changes in different parts of the structure.  Observations at 18cm are of outmost importance. Indeed, for SN1993J it can be expected that years after an emission decline (at 18 cm perhaps as long as 20 years) a nebular phase of expansion of progressively increasing emission will appear as in the case of SN1987A (Gaensler et al. \\cite{gae97}) and SN1979C (Montes et al.  \\cite{mon00}). Then, the appropriate wavelengths of observation will be the longest ones. Thus, it is relevant to establish a long record of 18 cm observations of SN~1993J of the maximum sensitivity now that the emission is optically thin and the shell structure is becoming conspicuous at this wavelength."
    },
    "0207/astro-ph0207211_arXiv.txt": {
        "abstract": "The compact elliptical galaxy M32 offers a unique testing ground for theories of stellar evolution. Because of its proximity, solar-blind UV observations can resolve the hot evolved stars in its center. Some of these late evolutionary phases are too rapid to study adequately in globular clusters, and their study in the Galactic field is often complicated by uncertainties in distance and reddening. Using the UV cameras on the Space Telescope Imaging Spectrograph, we have obtained a deep color-magnitude diagram (CMD) of the M32 center. Although the hot horizontal branch is well-detected, our CMD shows a striking scarcity of the brighter post-asymptotic giant branch (PAGB) and post-early AGB stars expected for a population of this size. This dearth suggests that the evolution to the white dwarf phase may be much more rapid than that predicted by canonical evolutionary tracks for low-mass stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207205_arXiv.txt": {
        "abstract": "A magnetically dominated plasma driven by motions on boundaries at which magnetic field lines are anchored is forced to dissipate the work being done upon it, no matter how small the electrical resistivity. Numerical experiments have clarified the mechanisms through which balance between the boundary work and the dissipation in the interior is obtained. Dissipation is achieved through the formation of a hierarchy of electrical current sheets, which appear as a result of the topological interlocking of individual strands of magnetic field. The probability distribution function of the local winding of magnetic field lines is nearly Gaussian, with a width of the order unity.  The dissipation is highly irregular in space as well as in time, but the average level of dissipation is well described by a scaling law that is independent of the electrical resistivity.  If the boundary driving is suspended for a period of time the magnetic dissipation rapidly drops to insignificant levels, leaving the magnetic field in a nearly force-free, yet spatially complex state, with significant amounts of free magnetic energy but no dissipating current sheets. Renewed boundary driving leads to a quick return to dissipation levels compatible with the rate of boundary work, with dissipation starting much more rapidly than when starting from idealized initial conditions with a uniform magnetic field.  Application of these concepts to modeling of the solar corona leads to scaling predictions in agreement with scaling laws obtained empirically; the dissipation scales with the inverse square of the loop length, and is proportional to the surface magnetic flux.  The ultimate source of the coronal heating is the photospheric velocity field, which causes braiding and reconnection of magnetic field lines in the corona. Realistic, three-dimensional numerical models predict emission measures, coronal structures, and heating rates compatible with observations. ",
        "introduction": "Magnetic fields are ubiquitous in astrophysical objects; circumstances where there is no magnetic field present are exceptions. Indeed, much of the non-thermal activity that is observed in astrophysical systems is probably related to the presence of magnetic fields.  Gravity is, indirectly, a major reason for the ubiquitous magnetic activity, because it tends to separate matter into dense and tenuous regions.  Magnetic fields that connect such regions are subjected to stress in the dense regions, and are forced to dissipate in the tenuous regions.  There, the magnetic field energy density can be many times higher than the thermal and kinetic energy density of the gas, and minor readjustments of the magnetic field may correspond to significant heating and acceleration of the gas.  Understanding the principles that control the dissipation of magnetic energy when the \\inx{plasma beta} ($\\beta=P_g/P_B$, where $P_g$ is gas pressure and $P_B$ is magnetic pressure) is low and the magnetic \\inx{Reynolds number} ($\\Rm = U L /\\eta$, where $U$ is velocity, $L$ size, and $\\eta$ magnetic diffusivity) is very high is a major challenge, and numerous research papers, review articles and books have been published on this subject over the years (e.g., Parker 1972, 1983, 1988, 1994; \\cite{Sturrock+Uchida81}; \\cite{vanBall86}; \\cite{Mikic+ea89}; \\cite{Heyvaerts+Priest92}; \\cite{Longcope+Sudan94}; \\cite{Galsgaard+Nordlund96a}b; \\cite{Nordlund+Galsgaard97}; \\cite{Gomez+00}; to mention just a few).  The solar corona is an ideal `test site' for theories and models of magnetic dissipation, since it provides rich opportunities for observing both the spatial and temporal structure of a dissipating low beta plasma. The Sun is indeed a `\\inx{Rosetta stone}' in the context of magnetic dissipation---once we understand how magnetic dissipation occurs under such well observed conditions we may be much more confident when extrapolating to more distant and less well observed circumstances.  Numerical experiments have emerged as a complementary and rich source of inspiration in the quest to understand magnetic dissipation. Below I briefly summarize conclusions from two types of experiments; generic experiments where a low beta plasma is driven from two opposing boundaries, and realistic experiments that attempt to model solar coronal conditions as closely as possible. ",
        "conclusions": "The results of the numerical experiments, and the properties of the scaling law derived from them, provide evidence that we are finally approaching a basic understanding of magnetic dissipation.  First of all, the basic form of the scaling law (\\ref{nordlund-eq-scaling}) follows from first principles, and agrees with previously proposed scaling laws (\\cite{Parker83d}, \\cite{vanBall86}).  In addition, it adds a prediction for the crucial inclination factor, for which Parker (1983) and van Ballegooijen (1986) had to make arbitrary assumptions.  Secondly, the \\inx{scaling law}  is robust in that the crucial inclination factor is bracketed from above and below.  It is bracketed from below because in general current sheets do not develop until the \\inx{local winding number} is of the order of unity. It is bracketed from above, because a local winding number much larger than unity inevitably leads to instabilities that rapidly dissipate the surplus magnetic energy.  Because of the spontaneous formation of a hierarchy of current sheets the scaling law is also robust against changes of the magnetic Reynolds number.  If anything, an increased Reynolds number could in principle lead to an {\\em increase} of the magnetic dissipation because, as pointed out by Parker (1988), if one assumes that a reduction of the magnetic diffusivity initially leads to a reduction of the magnetic dissipation the consequence is only that the boundary work for a while exceeds the dissipation, which leads to an increase of the average inclination and hence to a further increase of the boundary work.  When the magnetic dissipation eventually comes into balance with the boundary work again it happens at a {\\em higher} level than before.  But in practice, an increase of the magnetic Reynolds number $\\Rm$ just leads to an extension of the hierarchy of electrical current sheets to smaller scales, which makes it possible to dissipate at the same rate even without increasing the average angle of inclination at the boundaries.  Any claim to the contrary must be accompanied by a demonstration that it is possible to sustain a distribution of winding number that is substantially wider than unity at high $\\Rm$.  The scaling law (\\ref{nordlund-eq-scaling}) is furthermore consistent with \\inx{scaling laws derived from observations}, in that it predicts magnetic dissipation to scale as $L^{-2}$, in agreement with Porter \\& Klimchuk (1995).  Also, since the solar photosphere has a very intermittent distribution of magnetic field strength $B$, where $B$ is either very weak or of the order of 1 kG, the average heating is predicted to scale roughly as the \\inx{magnetic surface filling factor}.  This again agrees with the \\inx{observed scaling of coronal heating} (\\cite{Fisher+98})."
    },
    "0207/astro-ph0207519_arXiv.txt": {
        "abstract": "We present the results of a deep, optical/IR wide field imaging survey of selected fields in the nearby (d$\\sim 140$ pc) Taurus star-forming region. We report the discovery of 9 new members with spectral types M5.75-M9.5. We derive an Initial Mass Function encompassing 54\\% of the known members in Taurus. Comparison with dense regions like the Trapezium Cluster in Orion shows that Taurus has produced $\\times 2$ less brown dwarfs. We suggest that the lower frequency of brown dwarfs in Taurus may result from the low-density star-forming environment, leading to larger minimum Jeans masses. ",
        "introduction": "The variation of the Initial Mass Function (IMF) with the environment is a key problem in star formation, but any dependence of the IMF with the ambient conditions has been difficult to determine. Most surveys have concentrated on the most compact regions with the highest density of stars and gas, while extended, sparse groups of nearby young stars have not been studied with the same sensitivity, mainly because of the large area of the sky that they span. Therefore, meaningful comparisons of the IMF in differing conditions have not been readily available. This situation is changing with the advent of large-format CCD detectors, making possible deep surveys of the less embeded, more extended star-forming complexes. With a modest extinction (Av $\\la 4$) and low density of stars (N$\\sim 1-10$ pc$^{-3}$) Taurus is an example of the ``non-clustered'' mode of star formation, making it an ideal laboratory for comparison with high density regions like young clusters. ",
        "conclusions": ""
    },
    "0207/astro-ph0207033_arXiv.txt": {
        "abstract": "{\\small The cannonball model of GRBs is very overt (and, thus, falsifiable) in its hypothesis and results: all the considerations I review are based on explicit analytical expressions derived, in fair approximations, from first principles. The model provides a good description of {\\it all} the data on {\\it all} GRBs of known redshift, has made correct predictions, and is unprecedentedly self-consistent, simple and successful.} ",
        "introduction": "The cannonball (CB) model of GRBs \\cite{super,DD2000b,optical,radio} is based on {\\bf our} ignorance, for {\\bf we} (its authors) do not understand, e.g.: how the GRB engine works, how core-collapse supernovae (SNe) eject their ejecta, the transport of angular momentum in processes of collapse and/or accretion, relativistic magnetohydrodynamics, the relativistic ejections in quasars and microquasars... Thus, we base our starting hypothesis on analogy with the {\\it observations} of quasars and $\\mu$-quasars, overlooking current numerical simulations of these phenomena\\footnote{The definition \\#1.a of ``simulation'' in the OED is: ``The action or practice of simulating, with intent to deceive; false pretence, deceitful profession''.}. Quasars and $\\mu$-quasars appear to expel relativistic plasmoids when matter accretes abruptly from a disk or torus orbiting them. We assume the GRB engine to be similar: relativistic CBs are emitted axially from the recently made compact object in a core-collapse SN, as matter that has not been expelled as a SN shell (SNS) falls back \\cite{yo} to constitute an unstable disk. Most indications are that the plasmoids are made of ordinary matter, not some fancier substance such as $e^+\\,e^-$ pairs with some finely-tuned ``baryon-load'', as assumed in the conventional GRB scenarios: fireballs or their progeny (hereinafter ``the standard model (SM)''; for a balanced review, see \\cite{Ghis1}). ",
        "conclusions": ""
    },
    "0207/astro-ph0207669_arXiv.txt": {
        "abstract": "We have used the VLBA to produce a high dynamic range image of the nucleus of NGC 6251 at 1.6 GHz and snapshot images at 5.0, 8.4, and 15.3 GHz to search for emission from a parsec-scale counterjet.  Previous VLBI images at 1.6 GHz have set a lower limit for the jet/counterjet brightness ratio near the core at about 80:1, which is larger than expected given the evidence that the radio axis is fairly close to the plane of the sky.  A possible explanation is that the inner few pc of the counterjet is hidden by free-free absorption by ionized gas associated with an accretion disk or torus.  This would be consistent with the nearly edge-on appearance of the arcsecond-scale dust disk seen in the center of NGC 6251 by HST.  We detect counterjet emission close to the core at 1.6 GHz, but not at the higher frequencies. Given that the optical depth of free-free absorption falls off more rapidly with increasing frequency than the optically thin synchrotron emission from a typical radio jet, this result implies that the absence of a detectable parsec-scale counterjet at high frequencies is not due to free-free absorption unless the density of ionized gas is extremely high and we have misidentified the core at 1.6 GHz. The most likely alternative is a large jet/counterjet brightness ratio caused by relativistic beaming, which in turn requires the inner radio axis to be closer to our line of sight than the orientation of the HST dust disk would suggest. ",
        "introduction": "NGC 6251 (1637+826, J1632+8232, z=0.024) is an elliptical galaxy containing an apparently edge-on dust lane (\\citet{n83}; \\citet{cv97}) and a central stellar cusp and blue continuum light source (\\citet{y79}; \\citet{c93}). Spectroscopic observations with HST indicate the presence of a central black hole with a mass of $(6\\pm2) \\times 10^{8}\\ {\\rm M}_{\\odot}$ \\citep{ff99}. Although NGC 6251 is located near the edge of the cluster Zw 1609.0+8212, recent observations by \\citet{w00} show that this cluster contains at least three sub-clusters of galaxies that may not be physically related. NGC 6251 is associated with one of these sub-clusters. The distance to NGC 6251 is $72\\,h^{-1}$ Mpc, where $h = H_{\\circ}\\,/\\, {\\rm 100}\\ {\\rm km\\ s}^{-1}\\ {\\rm Mpc}^{-1}$.  The linear scale corresponding to 1 milliarcsec (mas) is $0.36\\,h^{-1}$ pc.  The line-of-sight velocity dispersion of the sub-cluster is less than 300 km s$^{-1}$, implying that it is a poor cluster with an X-ray atmosphere temperature of 0.7 keV \\citep{w00}.  Combined with plausible densities (consistent with the extended X-ray luminosity measured by \\citet{bw93}), this gas temperature is far too low to confine any part of the spectacularly linear (e.g., see Figure 1 in \\citet{w01}) kpc-scale radio jet.  The large-scale radio morphology of this source (\\citet{w77}; \\citet{w82}; \\citet{p84}; \\citet{j86a}) suggests that its radio axis is close to the plane of the sky.  Its high declination of +83$^{\\circ}$ allows very good (u,v) coverage by northern hemisphere VLBI arrays. Consequently, this is a good candidate for having a nearly edge-on inner accretion disk which could be detected via free-free absorption of radiation from a parsec-scale counterjet.  An early series of three global VLBI observations of NGC 6251 at 1.6 GHz was carried out during a five year period. The first 18-cm VLBI experiment in 1983 \\citep{j86a} showed a one-sided jet aligned with the VLA jet and containing a knot of emission approximately 25 mas from the core (we assume the core corresponds to the strong, unresolved peak at the eastern end of the jet).  No counterjet was detected at a limit of 80:1 measured $\\pm6$ mas from the core. Second and third epoch experiments were performed in 1985 and 1988 (\\citet{j86}; \\citet{jw94}) to look for motion of the 25-mas feature. No significant change in the separation between the core and the 25-mas knot was found (${\\rm{v/c}} < 0.23\\ h^{-1}$), although changes in jet morphology closer to the core were observed. The combination of low proper motion and large jet/counterjet brightness ratio implies, using the usual beaming model, that the 25-mas feature in the jet does not move with the bulk flow velocity or that the radio axis is much closer to our line of sight than expected from the large-scale radio structure and the nearly edge-on optical dust disk.  An alternative explanation for the lack of a detectable counterjet near the core is free-free absorption by ionized gas in front of the counterjet (plausibly in the form of a parsec-scale accretion disk or torus oriented perpendicular to the radio jets).  Evidence for absorption of radiation from the base of a counterjet has been seen in several other radio sources, including 3C84 (\\citet{v94}; \\citet{w94}), Centaurus A (\\citet{j96}; \\citet{tm01}), and NGC 4261 (\\citet{jw97}; \\citet{j00}; \\citet{j01}).  If a similar situation exists in NGC 6251, the counterjet should become visible farther from the core at 1.6 GHz, and closer to the core at higher frequencies.  The observations reported here were designed to detect the counterjet at 1.6 GHz, where its intrinsic brightness is likely to be larger, and at higher frequencies where it should be visible closer to the core.  Our ultimate goal is to understand the extent and structure of inner accretion disks on parsec and sub-parsec scales.  High resolution, high dynamic range, multi-frequency images of absorption by ionized gas are one of the best tools available for the study of these disks. NGC 6251 is one of the relatively few galaxies whose radio axis orientation and other properties make it a good candidate for such studies. ",
        "conclusions": "Free-free absorption does not appear to be responsible for the lack of a detectable parsec-scale counterjet at high frequencies in this source, unless high electron densities or path lengths are invoked and the short counterjet we see at 1.6 GHz is actually the (highly absorbed) core. The most likely alternative is a large jet/counterjet brightness ratio caused by relativistic beaming, which in turn requires the inner radio axis to be closer to our line of sight than the orientation of the HST dust disk would suggest.  In this case the smaller jet/counterjet brightness ratio seen farther from the core would be evidence for a gradual increase in the jet angle to our line of sight or a decrease in the bulk velocities in the jets."
    },
    "0207/astro-ph0207343_arXiv.txt": {
        "abstract": "Stars with very large mass loss on the red-giant branch can undergo the helium flash while descending the white-dwarf cooling curve.  Under these conditions the flash convection zone will mix the hydrogen envelope with the hot helium-burning core. Such ``flash-mixed'' stars will arrive on the extreme horizontal branch (EHB) with helium- and carbon-rich envelopes and will lie at higher temperatures than the hottest canonical (i.e., unmixed) EHB stars.  Flash mixing provides a new evolutionary channel for populating the hot end of the EHB and may explain the origin of the high gravity, helium-rich sdO and sdB stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207175_arXiv.txt": {
        "abstract": "{ We discuss the morphology and kinematics of five gigahertz-peaked spectrum (GPS) sources that have been observed with the VLBA.  We find a wide range of observed properties including core-jet structure, superluminal motion, variability, extended structure, and polarization, all of which appear to deviate from commonly-accepted GPS paradigms (e.g., O'Dea \\cite{OD98}). We suggest that the observed low frequency cutoff in GPS sources may be primarily due to free-free absorption rather than synchrotron self-absorption.  } ",
        "introduction": "GPS sources are characterized by their sharp low-frequency spectral cutoff and the absence of large scale structure.  Previous observations have suggested a simple double morphology with no evidence for significant relative component motions.  In this paper, we present new VLBA observations of five well-known GPS sources whose diverse properties are difficult to understand in terms of simple conventional models where the GPS sources are classical self-absorbed synchrotron sources that are the precursors of double-lobed radio galaxies. ",
        "conclusions": "All five of the sources we have studied have sharply bent jet structure. We find no other systematic properties which distinguish these GPS sources other than their peaked radio spectrum, which may be due to a combination of SSA and FFA from a surrounding ionized medium. CTA~102, 2134+004, and CTD~93 have bright cores plus fainter bent jets typical of what is observed in other quasars and AGN, whereas PKS 1345+125 and OQ~208 have double-sided jet structures that end in sharp bends.   The identically-peaked spectra of the separate components of individual GPS sources implies little differential Doppler shift of the spectral components if the cutoff is from SSA. This interpretation is supported by the absence of any differential component motions, possibly because all the components we are observing are moving at the same velocity with respect to an unidentified core below our level of detection.  But this interpretation is not consistent with observed the highly bent structure which would appear to require changes in the Lorenz factor corresponding to changes in the direction of the flow.  The common cutoff frequency across components of widely different surface brightness therefore suggests that some or all of the low frequency cutoff may be due to  FFA rather than SSA.  In this case the small angular size predicted from the early spectral observations and the assumption of SSA would be coincidental."
    },
    "0207/hep-ph0207260_arXiv.txt": {
        "abstract": "We calculate the time delays of neutrinos emitted in gamma ray bursts due to the effects of neutrino mass and quantum gravity using a time dependent Hubble constant which can significantly change the naive results presented hitherto in the literature for large redshifts, and gives some sensitivity to the details of dark energy. We show that the effects of neutrino mass, quantum gravity and dark energy may be disentangled by using low energy neutrinos to study neutrino mass, high energy neutrinos to study quantum gravity, and large redshifts to study dark energy. From low energy neutrinos one may obtain direct limits on neutrino masses of order $10^{-3}$ eV, and distinguish a neutrino mass hierarchy from an inverted mass hierarchy. From ultra-high energy neutrinos the sensitivity to the scale of quantum gravity can be pushed up to $E_{QG}\\sim 5\\times 10^{30}$ GeV. By studying neutrinos from GRBs at large redshifts a cosmological constant could be distinguished from quintessence. ",
        "introduction": "Gamma ray bursts (GRBs) are amongst the most distant, energetic and enigmatic astrophysical phenomena known. Understanding GRBs is arguably the most outstanding question in astronomy, and one which may be answered by a plethora of gamma ray observatories such as INTEGRAL, SWIFT and BATSE, and corresponding infrared and optical telescopes such as REM and LT \\cite{GRBreview}. It is well known that 99\\% of the energy of a supernova is emitted in the form of neutrinos, and therefore it is widely expected that GRBs are similarly a copious source of neutrinos which may be detected in future neutrino telescopes \\cite{Halzen:1996qw, Alvarez-Muniz:2000st, Halzen:1999xc, Gupta:2002zd}. Within this decade it is therefore likely that GRBs will become much better understood, and their exact nature and mechanisms which drive their internal engines will be revealed. For example it may turn out that a GRB results from the core collapse of a very massive supernova to a compact rotating black hole with the energy emitted in beamed relativistic fireball jets containing copious neutrino fluxes \\cite{Waxman:1997ti, Waxman:1998yy, Bahcall:1999yr, Waxman:1999ai, Bahcall:2000sa, Halzen:2002pg} Alternatively the GRB engine could result in the emission of beamed earth-sized cannonballs \\cite{Rujula:2002xj}.  In this paper we are not concerned with detailed models of GRBs, but instead regard them as a high intensity, high energy neutrino beam with a cosmological baseline. We shall be interested in the time delay of the arrival of neutrinos. We show that the time delay may be used as a probe of three physical effects: (i) neutrino mass, (ii) quantum gravity, (iii) dark energy. The time delay due to neutrino mass has been noted earlier in \\cite{Halzen:1996qw} while effects of quantum gravity on the time of flight have been considered in \\cite{Amelino-Camelia:1997gz, Biller:1998hg, Schaefer:1998zg, Ellis:1999sd} for high energy photons and in \\cite{Alfaro:1999wd, Ellis:1999sf} for neutrinos. However none of the above papers considers the time delays due to massive neutrinos in the presence of quantum gravity, although \\cite{Alfaro:1999wd} gives the dispersion relation for this case. Moreover no studies to date have calculated time delays using a formalism which correctly takes into account the time dependence of the Hubble constant due to matter and dark energy. In our study we consider the time dependence of the Hubble constant due to matter and dark energy and show that they can change the naive results by more than 100\\% for $z>1$. For large redshifts the results for time delays due to neutrino mass and quantum gravity are sensitive to the nature of dark energy. We show that the three effects may be disentangled by using low energy neutrinos to study neutrino mass, high energy neutrinos to study quantum gravity, and using large redshifts to study dark energy, leading to the results stated in the abstract.  For the determination of the neutrino mass, in principle one could compare the arrival time of the massive neutrinos with the arrival times of the photons emitted in the GRB, assuming them to be emitted at the same time. However this would be plausible only if the GRB were a point source {\\it in vacuo}, which it is not. In any realistic GRB model the photons are trapped inside the fireball and are released much later -- the exact amount of time delay being highly model dependent. Another strategy might be to compare the arrival times of the low energy neutrinos with that of the ultra-high energy ones, since (as we will show) the ultra-high energy neutrinos suffer negligible time delay due to mass. But again one would expect the low energy neutrinos to be produced thermally, as in supernovae, leading to a model dependent time delay.  However there are alternative strategies which could overcome these problems. To begin with, if the neutrinos are hierarchical in mass, $m_1\\ll m_2\\ll m_3$, then due to their mixing, neutrinos of the {\\em same energy} will arrive at the detector in three bunches, corresponding to the three mass eigenstates, the arrival time only depending on their mass. We can then compare the arrival times of the different neutrino mass eigenstates and put limits on the neutrino mass. This represents a ``clean strategy'' and indeed similar arguments for constraining neutrino mass from time delays using supernova neutrinos have been used in the literature (see \\cite{Beacom:1998yb} and references therein). However we will show that with GRB neutrinos the sensitivities possible are better by orders of magnitude.  For constraining models of quantum gravity and dark energy we can use arrival times of the ultra-high energy neutrinos, which are produced over many decades of neutrino energy. We can compare the arrival times of the high energy neutrinos of different energies -- which unlike the low energy thermal neutrinos are expected to be produced at almost the same time at the source -- and put limits from such observations. Thus, the results in this paper do not rely on the comparison of the arrival times between photons and neutrinos, which would involve the uncertainties discussed above. However for simplicity, we shall calculate arrival time delays relative to a hypothetical low energy photon which is assumed to be emitted at the same time. This is for convenience only; in a realistic search strategy what will be important will be the comparison of time delays between neutrinos of either the same energy, or between high energy neutrinos of two different energies as discussed above. ",
        "conclusions": "In this paper we have calculated the time delays of neutrinos emitted in gamma ray bursts due to the effects of neutrino mass and quantum gravity. Our results are based on the formula for $\\Delta t$ in Eq.\\ref{tdel2}, using the Hubble constant for a flat universe calculated using Eq.\\ref{hz} for different dark energy models. This formalism correctly takes into account the time dependence of the Hubble constant due to matter and dark energy, and can change the naive results in the literature by more than 100\\% for $z>1$. We have shown that the effects of neutrino mass, quantum gravity and dark energy may be disentangled by using low energy neutrinos to study neutrino mass, high energy neutrinos to study quantum gravity, and large redshifts to study dark energy. From low energy neutrinos one may obtain direct limits on neutrino masses of order $10^{-3}$ eV, and distinguish a neutrino mass hierarchy from an inverted mass hierarchy. From ultra-high energy neutrinos the sensitivity to the scale of quantum gravity can be pushed up to $E_{QG}\\sim 5\\times 10^{30}$ GeV. By studying neutrinos from GRBs at large redshifts a cosmological constant could be distinguished from quintessence.  For convenience we have calculated all time delays with respect to a hypothetical low energy photon, assumed to be emitted at the same time from a point source as the neutrino of a given energy. We emphasise that what is important in practical search strategies is not these time delays themselves, which will be unmeasurable due to the uncertainties in the emission characteristics of low energy photons, but rather the comparison of time delays between neutrinos of either the same energy, or between high energy neutrinos of two different energies. As already mentioned, if the neutrinos are hierarchical in mass then neutrinos of the {\\em same energy} will arrive at the detector in three bunches, and we can then compare the arrival times of the different neutrino mass eigenstates and put limits on the neutrino mass, as in the case of supernova neutrinos \\cite{Beacom:1998yb} but with sensitivities better by orders of magnitude. For constraining models of quantum gravity and dark energy one can use arrival times of the ultra-high energy neutrinos, and can compare the arrival times of the high energy neutrinos of different energies.    \\vskip 1in  \\begin{center} {\\bf Acknowledgement} \\end{center} {We particularly acknowledge very helpful discussions with Tom Marsh. We also wish to thank Juan Garcia-Bellido, Christian Kaiser, Ian McHardy, Tim Morris and Graham Shore for helpful discussions, and Francis Halzen for useful communications.}     \\vskip 1in"
    },
    "0207/astro-ph0207613_arXiv.txt": {
        "abstract": "We present high resolution HST imaging of the nebula associated with the galactic LBV candidate HD 168625, together with ISO imaging and AAT echelle spectroscopy. The overall nebular morphology is elliptical with the major axis at PA $\\simeq$ 120$^{\\deg}$. The dimensions of the nebula are 12$''$ x 16$''$.7 at H$_{\\alpha}$ and 15$''$.5 x 23$''$.5 at 4 $\\mu$m. In the HST H$\\alpha$ image, the nebula is resolved into a complex structure of filaments and arcs of different brightness. The asymmetry is lost in the HST continuum image where the nebula appears more diffuse and richer in filaments and clumps with the shape of cometary tails. At 11.3 $\\mu$m the nebular emission peaks in two diametrically opposite lobes, placed on the nebula boundaries and along its major axis. A very faint loop is also visible at optical wavelengths, north and south of the shell. We suggest that the nebula is an ellipsoid with projected sizes of 14$''$ and 9$''$ (0.19 pc $\\times$ 0.12 pc) along the RA and DEC directions, respectively. This ellipsoid is expanding at 19 km s$^{-1}$ and is dynamically as old as $\\simeq$ 4800 yrs; it probably interacts with the stellar wind and the loop so that PAH emission is detected from its caps, i.e. the lobes seen in the ISO images. The chemistry of the loop suggests that it is composed of un-processed material, probably from the local interstellar medium swept by the stellar wind. ",
        "introduction": "\\label{sec:intro} It is widely recognized that Luminous Blue Variables (LBVs) represent a post-main sequence phase in which massive stars (M$_i \\geq$ 20 M$_{\\odot}$, Langer et al. 1994) lose a considerable amount of mass via giant eruptions and minor outbursts. The ejected gas and dust build up a circumstellar nebula chemically enriched by the central star nucleosynthesis. From the expansion velocity of known LBV nebulae, a dynamical age of a few 10$^4$ years is usually inferred, which points to a very short-lived evolutionary phase. For this reason, LBVs are rare objects; indeed, only 40 are classified as such in the whole Local Group (Humphreys \\& Davidson 1994). P Cygni and AG Carinae are considered the prototypes of the LBV class. \\par\\noindent The LBV nebulae provide us with a wealth of details about their central stars. Their chemical composition is used to determine the evolutionary phase when the LBV instability triggers the nebula ejection (cf. Smith et al. 1997 and Waters et al. 1999 in the case of AG Carinae), and their morphology constrains the physics of the central star wind. Except P Cygni, all LBV nebulae display an asymmetric morphology, progressing from elliptical (e.g. AG Carinae) to bipolar (e.g. Eta Carinae and HR Carinae). Nota et al. (1995) reproduced the observed shapes through an interacting wind model, where a spherical stellar wind interacts with a pre-existing density contrast between the equatorial and polar direction. This density contrast could be, for example,  induced by mass transfer in a binary or rotation in a single star  (cf. Bjorkman \\& Cassinelli 1993, Owocki et al. 1998). From the H$\\alpha$ imaging it is also possible to estimate the ionized gas  mass in the LBV nebulae which is a key parameter in the understanding of the total mass ejected and of crucial importance  to  constrain  the evolutionary models for massive stars. \\par Unfortunately, the diagnostic power of the nebular morphology is limited by the spatial resolution accessible from the ground. Coronographic imaging has so far been the observational technique achieving the highest resolution possible from the ground: it has been able to resolve the global symmetry of LBV nebulae out to the LMC  (Nota et al. 1995) and the nebular fine structure only for close by objects,  such as AG Carinae. However, the  comparison with hydrodynamic models (cf. Frank 1997, Garcia-Segura et al. 1997) obviously requires higher levels of morphological details   to properly constrain the shaping mechanism at work  in LBV nebulae. Further, a more complete analysis should rely on high resolution imaging of ionized/neutral gas and dust; this would also assess the total  (gas + dust) mass of LBV nebulae and hence the total mass lost by the central star during the outburst. For these reasons, we have re-observed a complete sample of LBV nebulae in the Galaxy and in the Large Magellanic Cloud (Schulte-Ladbeck et al. 2002) with HST/WFPC2 and ISO/ISOCAM, among which is the galactic LBV candidate HD 168625. \\par HD 168625 is known to be variable with an amplitude of 0.06 mag (van Genderen et al. 1992), although its variability does not closely follow the typical pattern observed in the case of bona-fide LBVs. Nevertheless, its LBV candidacy was proposed when Hutsemekers et al. (1994), for the first time, resolved an associated  circumstellar nebula. Nota et al. (1996) imaged the nebula with the STScI Coronograph and resolved it into an  elliptical shell surrounded by two faint filaments forming a northern and southern loop. A gas mass of $\\sim$ 0.5 M$_{\\odot}$ was derived from images in the light of H$\\alpha$. Nota et al. also acquired two sets of spectroscopic data, six months apart, where HD 168625 was seen to fade by 0.3 mag and cool from T$_{eff}$ = 15,000 K to 12,000 K at a constant mass-loss rate of $\\sim$ 1.1 $\\times$ 10$^{-6}$ M$_{\\odot}$yr$^{-1}$. In addition, the spectra were used to derive the plasma properties of the nebula: for an assumed T$_e$ of 7,000 K, the average density is 1000 cm$^{-3}$ and the nitrogen content is Log(N/H)$+$12 = 8.04 indicating that the nebula is composed of stellar ejecta. Unfortunately, the spectra were taken at low resolution so that the kinematic structure of the nebula was not fully resolved. Nota et al. could measure the brighter, blueshifted edge of the nebula which apparently defined an expansion motion of $\\sim$ 40 km s$^{-1}$ and a dynamical age of $\\sim$ 10$^3$ yrs. \\par\\noindent The nebula surrounding HD 168625 was also observed in the mid infrared by Skinner (1997) and Robberto \\& Herbst (1998). Their images (taken between 4.7$\\mu$m and 20$\\mu$m) revealed emission by warm dust in the eastern and western edges of the optical nebula. Skinner (1997) derived a dust mass of 9.5 $\\times$ 10$^{-5}$ M$_{\\odot}$ for a distance of 2.2 kpc. Robberto \\& Herbst (1998) revised the distance of HD 168625 from 2.2 to 1.2 kpc and estimated the dust mass of the nebula to be $\\sim$ 0.003 M$_{\\odot}$. \\par We have ``revisited'' the nebula associated with HD 168625 employing the high spatial resolution of WFPC2 onboard HST and the high spectral resolution of UCLES on the AAT. These new data are complemented with ISO/ISOCAM observations in order to derive a multi-wavelength, detailed analysis of the nebular morphology and kinematics which are then used to trace back the outburst history of HD 168625. The data are presented in Section~\\ref{sec:data}. The HST and ISO images are discussed in Sections~\\ref{sec:hst} and ~\\ref{sec:iso} respectively, and the nebular  properties and kinematics are found in  Sections~\\ref{sec:mass} and \\ref{sec:neb}, and the stellar spectrum from the echelle data-set is presented in Section~\\ref{sec:star}. In Section~\\ref{sec:puzzle} we assemble and discuss the overall morphology of the HD 168625 nebula. ",
        "conclusions": "\\label{sec:puzzle} Multi-colour imaging of the nebula associated with HD 168625, obtained from space with HST and ISO, shows an elliptical  nebula with the major axis oriented at PA $\\sim$ 120$^\\circ$. Its main structure is a shell which HST has resolved into nesting filaments. Both shell morphology and size depend on wavelength: \\par\\noindent {\\it i)} in the H$\\alpha$ light, the southern edge is well defined while the northern part dissolves into diffuse circumstellar matter. The shell surface brightness is higher along the southern edge and peaks between PA $\\sim$ 140$^\\circ$ and PA $\\sim$ 210$^\\circ$. \\par\\noindent {\\it ii)} The shell  disappears in the V continuum light, where it is replaced by diffuse, filamentary emission within a faint outline which traces the boundaries of the gaseous nebula. At these wavelengths ($\\lambda_c \\sim$ 5480 \\AA) the bulk of the emission seems to originate from a substructure, the {\\it paddle},  to the NE of the central star. \\par\\noindent {\\it iii)} The shell lights up again at mid-infrared wavelengths, 3.1 -- 5.2 $\\mu$m and 10.8 -- 11.9 $\\mu$m, but this time the emission peaks in two lobes which are adjacent to the inner edge of the shell, in correspondence to {\\it dark or obscured} regions in the optical image. The {\\it paddle} and the southern portion of the shell, which dominate the continuum and H$\\alpha$ emissions, are here far less pronounced. \\par\\noindent {\\it iv)} The size of the nebula is different at optical and mid-IR wavelengths. In the light of H$\\alpha$, the nebula has an extension of 12$''$ $\\times$ 16.$''$7 (0.16 $\\times$ 0.23 pc), if we exclude the northern loop. At 3.1 -- 5.2 $\\mu$m, the nebula has a larger extension of 15.$''$5 $\\times$  23.$''$5 (0.21 $\\times$ 0.32 pc), and at 10.8 -- 11.9 $\\mu$m, an even larger size of 31$''$ $\\times$ 35.$''$5 (0.42 $\\times$ 0.48 pc). \\par Aperture photometry of the optical images indicates that the {\\it paddle} and the shell southern edge may differ in the gas-to-dust content, with the {\\it paddle} more dusty and the southern edge of the shell more gas-rich. Aperture photometry of the infrared images suggests that the eastern and western lobes of the shell are dominated by dust and PAH molecules, in agreement with the findings of Skinner (1997) and Robberto \\& Herbst (1998) who resolved the lobes up to 20 $\\mu$m. \\par\\noindent We believe that the nebular chemistry and kinematics derived here can explain the observed composite morphology of HD 168625. \\par\\noindent Nota et al. (1996) already derived for the nebula a Log(N/H) $+$ 12 of $\\simeq$ 8.04, showing that the nebula is indeed N-enriched and hence of stellar origin. They could not resolve the nebula into its components, as we did in this work with the help of echelle data. The higher spectral resolution of our data has allowed us to separate the shell from the loop and measure for each component the [NII]6584/H$\\alpha$ intensity ratio, that we use as a N-abundance indicator. It turns out that the [NII]6584/H$\\alpha$ ratio is, on average, 0.51 $\\pm$ 0.15 for the shell and 0.22 $\\pm$ 0.05 for the loop, which means that the shell is N-enriched relative to the loop. In particular, the [NII]6584/H$\\alpha$ ratio measured for the loop is very similar to what is observed for galactic HII regions. We have indeed selected a number of galactic HII regions from the sample of Shaver et al. (1983) which are at the same distance of HD 168625 and/or have a plasma temperature T$_e$([NII]) close to what assumed for HD 168625 (T$_e \\sim$ 7000 K, Nota et al. 1996). We have computed their relative [NII]6584/H$\\alpha$ intensity ratios and determined the mean value of 0.26 $\\pm$ 0.08, which compares well to a [NII]6584/H$\\alpha \\simeq$ 0.22 in the loop. This therefore implies that the loop is composed of un-processed material which has been blown by the stellar wind. It is unlikely that the loop is an interstellar bubble such the one detected NW of HR Carinae (Nota et al. 1997): its dynamical age is at most $\\simeq$ 7300 yrs old, too young for a HII region. The possibility that the loop is a previous stellar ejection also seems improbable: in this scenario, about 3000 years (i.e. the age difference between the loop and the shell) would have been enough for the star to self-enrich in N by almost a factor of 2. This time interval appears to be quite short with respect to evolutionary models of B stars (Lamers et al. 2001). Therefore, we suggest that the loop is local interstellar medium swept by the stellar wind. The fact that we have been able to detect the loop only in the DEC direction may imply that the loop lies preferentially on a plane, maybe the equatorial plane of the star. \\par\\noindent The radial velocities measured for the shell indicate that it is an ellipsoid expanding at 19 km s$^{-1}$ in both the RA and DEC directions. Its axes are 14$''$ (0.19 pc) and 9$''$ (0.12 pc) along the RA and DEC directions, respectively, in agreement with that estimated from the HST images. Such a morphology is very common among LBV nebulae, the best known example being AG Carinae. Nota et al. (1995) explained the elliptical shape of AG Carinae by invoking a density contrast between the stellar equator and poles which would restrict the nebula expansion on the stellar equatorial plane and produce a ``waist'' in the circumstellar nebula. The density contrast would leave the stellar polar axis as the only free direction to the nebular expansion and therefore would shape the nebula into an ellipsoid. There exist several ways to produce a density contrast (Livio 1995), such as stellar rotation and stellar binarity. Although no solid evidence for stellar rotation or binarity is available for HD 168625, one of these mechanisms could be responsible for the overall morphology of the nebula associated with HD 168625 and also would preferentially direct the present stellar wind along the polar axis of the star, i.e. the RA axis of the nebular ellipsoid, so that it would interact with the ``caps'' of the ellipsoid. The interaction would destroy the CO molecules and give rise to PAH emission as detected by ISO in the lobes of the nebula. Moreover, the HST images suggest that the shell caps may be interacting with the loop and this could also produce PAH emission."
    },
    "0207/astro-ph0207339_arXiv.txt": {
        "abstract": "This poster is a summary of a paper on the subject (Bayer-Kim et al. 2002) that has been submitted for publication in MNRAS. We present X-ray, H$\\alpha$, radio, and optical data, both imaging and spectroscopy, of the cluster RX\\,J0820.9+0752 and its peculiar central cluster galaxy (CCG). We announce the discovery of several isolated off-nuclear patches of blue light with strong line-emssion (total L(H$\\alpha)\\sim 10^{42}\\mbox{ergs}^{-1}$) and distinct ionizational and kinematic properties, embedded in a region of H$\\alpha$/X-ray emission. We propose  and investigate a scenario in which a secondary galaxy also featured in our observations has moved through a cooling wake produced by the CCG, thereby producing the observed clumpy morphology and helping to trigger star-formation within the clumps. ",
        "introduction": "From the strongly peaked profile in our 9.4\\thinspace ks {\\sl Chandra} observation (Fig. 1a), we derive that the hot intracluster gas in the cluster RX\\,J0820.9+0752 is cooling within a radius of $r\\approx 20$\\thinspace kpc. The mass deposition rate of a few tens of solar masses per year is consistent with the new, reduced cooling flow scenario proposed by Voigt et al. (2002) as a possible solution to the problem of the missing soft X-ray luminosity in cooling flow clusters (Peterson et al. 2001). The X-ray emission appears extended to the NW and is coincident with a luminous H$\\alpha$ nebula seen in an AAT image taken through a narrow-band filter in the H$\\alpha$ light (Fig. 1b). The cluster contains only a weak radio source (flux density of $0.45\\pm 0.13$\\thinspace mJy at 4.89\\thinspace GHz as determined from a complete VLA snapshot campaign; Edge et al. 2002, in prep.), but in an HST image its central galaxy shows highly peculiar morphology with two arcs of clumped emission to the NE and a secondary elliptical to the SE, opposite the X-ray/H$\\alpha$ feature (Fig. 1c). Several other isolated blobs are also apparent scattered further out to the NW and E.  \\begin{figure} \\centering \\epsfig{file=r0821_x-opt.eps,height=60truemm, width=90truemm}\\\\ \\epsfig{file=ttfcol.eps,height=60truemm, width=90truemm}\\\\ \\epsfig{file=hstcol.eps,height=60truemm, width=90truemm} \\caption{(a) The slightly smoothed $0.5-2$\\thinspace keV {\\sl Chandra} image of the cluster emission from RX~J0820.9+0752 with the optical contours overlaid (top). (b) Continuum-subtracted H$\\alpha$ emission around the central galaxy. The image shown is $20.5\\times 16$ arcsec (middle). (c) F606W HST image ($\\Delta\\lambda \\approx 4490-5910$\\AA\\ at the redshift of the object) of the central region of the cluster on the same scale as the H$\\alpha$ image (bottom). North is to the top and East to the left in all the figures.} \\end{figure} ",
        "conclusions": "The association of the H$\\alpha$ nebula with the X-ray emission is reminiscent of a similar feature found recently in A1795 (Fabian et al. 2001b). The most probable explanation given in this paper is that the filament represents a ``cooling wake'', produced by the central cluster galaxy moving through the hot ICM. We show that the same explanation can be applied to RX\\,J0820.9+0752, where the weakness of the radio source practically rules out the possibility that it has played a significant role in producing the wake. We find that the energy output through the strong line-emission is consistent with ionization from energetic photons produced by young stars. However, we observe a velocity offset of up to 200\\thinspace kms$^{-1}$ between the young stellar components and the line-emitting gas in the filament, suggesting they are not co-spatial. This could indicate that massive stars are not the only ionization source for the nebula. Maybe processes such as the ``cold mixing model'' proposed by Fabian et al. (2001a) can account for at least part of the observed line-emission.  An important factor in producing the clumpy morphology and triggering the star-bursts seems to be the secondary elliptical galaxy, whose velocity properties relate it to the cooling gas. We propose a scenario in which the secondary galaxy passes through the cooling wake produced by the CCG.  \\begin{figure} \\centering \\caption{Fit to one of the blue blobs. The spectrum is represented by the dashed line, the model by the solid one. The dotted spectrum is the difference between the late component of the model and the blob spectrum, clearly showing the blue excess and the Balmer and [OII]$\\lambda 3727$\\AA{} emission lines.} \\epsfig{file=HaSB.eps,height=90truemm,width=45truemm,angle=270} \\label{Fit} \\end{figure}"
    },
    "0207/astro-ph0207425_arXiv.txt": {
        "abstract": "{ \\ion{H}{i} observations of high-velocity clouds (HVCs) indicate, that they are interacting with their ambient medium. Even clouds located in the very outer Galactic halo or the intergalactic space seem to interact with their ambient medium. In this paper, we investigate the dynamical evolution of high velocity neutral gas clouds moving through a hot magnetized ambient plasma by means of two-dimensional magnetohydrodynamic plasma-neutral gas simulations. This situation is representative for the fast moving dense neutral gas cloudlets in the Magellanic Stream as well as for high velocity clouds in general. The question on the dynamical and thermal stabilization of a cold dense neutral cloud in a hot thin ambient halo plasma is numerically investigated. The simulations show the formation of a comet-like head-tail structure combined with a magnetic barrier of increased field strength which exerts a stabilizing pressure on the cloud and hinders hot plasma from diffusing into the cloud. The simulations can explain both the survival times in the intergalactic medium and the existence of head-tail high velocity clouds. ",
        "introduction": "High-velocity clouds (HVCs) -- first discovered by Muller et al. (\\cite{muller}) -- are defined as neutral atomic hydrogen clouds with radial velocities that cannot be explained by simple galactic rotation models. After almost 40 years of eager investigations there is still no general consensus on the origin and the basic physical parameters of HVCs (e.g.~Bregman \\cite{bregman}). The most critical issue of HVC research is the distance uncertainty. Danly et al. (\\cite{danly}), Keenan et al. (\\cite{keenan}) and Ryans et al. (\\cite{ryans}) consistently determined an upper distance limit of d $\\leq$  5 kpc to HVC complex M. The most important step forward is the very recently determined distance bracket of 4 $\\leq$ d $\\leq$ 10 kpc towards HVC complex A by van Woerden et al. (\\cite{van Woerden}). These results clearly place the HVC complexes M and A in the gaseous halo of the Milky Way.  Parallel to the growing evidence that a significant fraction of the HVC complexes are located in the Milky Way halo, Blitz et al. (\\cite{blitz}) supported the hypothesis that some HVCs are of extragalactic origin. They argued, that it is reasonable to assume that primordial gas -- left over from the formation of the Local Group galaxies -- may appear as HVCs. Braun \\& Burton (\\cite{bb99}) identified 65 compact and isolated HVCs and argued that this ensemble represents a homogeneous subsample of HVCs at extragalactic distances. Observational evidence for extragalactic HVCs may also be found by the detection of the highly ionized high-velocity gas clouds by Sembach et al. (\\cite{sembach}), because of its very low pressure of about p k$^{-1} \\approx 5 \\ \\mathrm{K cm}^{-3}$.  The Magellanic Stream (MS) and the Leading Arm (LA) (Putman et al. \\cite{putman}) both form coherent structures over several tens of degrees having radial velocities in the HVC regime. They represent debris most likely caused by the tidal interaction of the Magellanic Clouds with the Milky Way. The distances to these features are of the order of 50 kpc.  Meyerdierks (\\cite{meyerdierks}) detected a HVC that appears like a cometary shaped cloud with a central core and an asymmetric envelope of warm neutral atomic hydrogen (the particular HVC is denoted in literature as HVC A2). He interpreted this head-tail structure as the result of an interaction between the HVC and normal galactic gas at lower velocities. Towards the HVC complex C, Pietz et al. (\\cite{pietz96}) discovered the so-called {\\sc Hi} ``velocity bridges'' which seem to connect the HVCs with the normal rotating interstellar medium. The most straight forward interpretation for the existence of such structures is to assume that a fraction of the HVC gas was stripped off the main condensation. Br\\\"uns et al. (\\cite{bruens00}) extended the investigations of Meyerdierks (\\cite{meyerdierks}) and Pietz et al. (\\cite{pietz96}) over the entire sky covered by the new Leiden/Dwingeloo {\\sc Hi} 21-cm line survey (Hartmann \\& Burton \\cite{hartmann97}) and found head-tail structures in all HVC-complexes including the Magellanic Stream, except for the very faint HVC-complex L. Their analysis revealed that the absolute value of the radial velocity of the tail is always lower than the value for the head of the HVC ($|v_{\\rm LSR,tail}| < |v_{\\rm LSR,head}|$). In addition, it was shown that the fraction of HVCs showing a head-tail structure increases proportional to the peak column density and increasing radial velocity $|v_{\\rm GSR}|$.  HVCs mostly appear as ``pure'' neutral atomic hydrogen clouds. Absorption line studies provide information on the ionization state and the metalicity of HVCs. The results indicate that the bright and very extended HVC complexes consist (at least partly) of processed material, having $\\le$ 1/3 of the solar abundances (Wakker \\cite{wakker-met}). \\\\ Thus, the question of the dynamical and thermal stabilization of a cold dense neutral cloud in a hot thin ambient plasma arises. To confine a HVC in the Magellanic Stream by pressure a relatively high halo gas density is required (Mirabel et al.~\\cite{mirabel}, Weiner \\& Williams \\cite{weiner}). Therefore, the lifetime of the HVCs should be significantly limited by evaporation (Murali \\cite{murali}) since the necessary pressure for confinement is associated with a high energy transfer. An alternative mechanism of dynamical stabilization is the magnetic confinement of the cloud. In this contribution, we numerically investigate the formation of a magnetic barrier around the HVC in the Magellanic Stream and its stabilizing effects on the neutral gas cloud. ",
        "conclusions": "We presented the first plasma-neutral gas simulations of the interaction of HVCs with the Galactic halo gas. Previous numerical studies (Santilla\\'an \\cite{santi} and Quilis \\cite{quilis}) focused on the dynamics in external, Galactic gravitational fields. The questions of thermal and magnetic insulation of the neutral gas clouds in the hot ambient medium as well as their stability were not addressed. Our two-fluid approach provides rather detailed information on the interaction region of the cold almost neutral HVC gas  and the hot almost fully ionized ambient plasma. The Galactic magnetic field is locally draped around the HVC, i.e., a magnetic barrier forms similar to magnetopauses around comets (McComas \\cite{comas}) and unmagnetized planets (Luhman \\cite{luhma}). This barrier is of great importance for the thermal insulation of the cold, fast traveling cloud from the hot surroundings. Furthermore, it prevents the plasma from diffusing into the neutral gas cloud. We conclude that the formation of magnetic barriers contributes significantly to the dynamical stability of HVCs. Even for relatively weak initial magnetic fields in the Halo the magnetic barrier sufficiently stabilizes the cloud. In particular, unstable Kelvin-Helmholtz modes are not excited due to the stabilizing effect of a magnetic field component parallel to the flow. Thermal radiative losses that are not included in our contribution may further enhance the stability of the boundaries against Kelvin-Helmholtz modes (Vietri et al.~\\cite{vie97}). Within the regions of discontinuity and in the adjacent inner layers the physics becomes rather involved. In particular, some fraction of the stored magnetic field energy can be converted into heat by magnetic reconnection and thus can be responsible for the observed x-ray emission (Zimmer et al.~\\cite{zimmer}, Birk et al.~\\cite{birk2}). Moreover, one may note that the effect of ionization caused by the relative cloud halo movement (Konz et al.~\\cite{konz}) may enhance the formation of magnetic barriers by ionized boundary layers. \\\\ A self-consistent three-dimensional simulation including all relevant physical effects seems to be highly desirable and is a promising task for the future.  Finally, we note that the interaction scenario studied in this paper may also be of relevance in the context of cloudlets in the magnetospheres of Active Galactic Nuclei as well as in cluster cooling flows."
    },
    "0207/astro-ph0207080_arXiv.txt": {
        "abstract": "{We complete previous investigations on the thermodynamics of self-gravitating systems by studying the grand canonical, grand microcanonical and isobaric ensembles. We also discuss the stability of polytropic spheres in connexion with a generalized thermodynamical approach proposed by Tsallis. We determine in each case the onset of gravitational instability by analytical methods and graphical constructions in the Milne plane. We also discuss the relation between dynamical and thermodynamical stability of stellar systems and gaseous spheres. Our study provides an aesthetic and simple approach to this otherwise complicated subject. ",
        "introduction": "\\label{sec_introduction}  The statistical mechanics of systems interacting via long-range forces exhibits peculiar features such as negative specific heats, inequivalence of statistical ensembles and phase transitions. These curious behaviours have been first discussed in the astrophysical literature during the elaboration of a thermodynamics for stars (Eddington 1926), globular clusters (Lynden-Bell \\& Wood 1968), black holes (Hawking 1974) and galaxies (Padmanabhan 1990). They have been recently rediscovered in different fields of physics such as nuclear physics, plasma physics, Bose-Einstein condensates, atomic clusters, two-dimensional turbulence... The main challenge is represented by the construction of a thermodynamic treatment of systems with long-range forces and by the analogies and differences among the numerous domains of application (see Chavanis 2002e and other contributions in that book).  Gravity provides a fundamental example of unshielded long-range interaction for which ideas of statistical mechanics and thermodynamics can be developed and tested.  For systems with long-range interactions, the mean-field approximation is known to be {\\it exact} in a suitable thermodynamic limit. Therefore, the structure and stability of self-gravitating systems at statistical equilibrium can be analyzed in terms of the maximization of a thermodynamical potential. This thermodynamical approach leads to isothermal configurations which have been studied for a long time in the context of stellar structure (Chandrasekhar 1942) and galactic structure (Binney \\& Tremaine 1987). As is well-known, isothermal spheres have infinite mass so that the system must be confined within a box (Antonov problem) in order to prevent evaporation and make the thermodynamical approach rigorous. It is also well-known that isothermal configurations only correspond to {\\it metastable} equilibrium states (i.e., local maxima of the thermodynamical potential), not true equilibrium states. These metastable equilibrium states are expected to be relevant, however, for the timescales contemplated in astrophysics. In particular, globular clusters described by Michie-King models are probably in such metastable states.  The series of equilibria of finite isothermal spheres can be parametrized by the density contrast between the center and the boundary of the system. For sufficiently low density contrasts, the system is  thermodynamically stable. However, instability occurs at sufficiently large concentrations: at some point in the series of equilibria, the solutions cease to be local maxima of the thermodynamical potential and become unstable saddle points. The crucial point to realize is that the onset of instability depends on the statistical ensemble considered: microcanonical (MCE), canonical (CE) or grand canonical (GCE). This contrasts with ordinary systems, with short range interactions, for which the statistical ensembles are equivalent at the thermodynamic limit, except near a phase transition. This suggests that gravitating systems in virial equilibrium are similar to normal (extensive) systems at the verge of a phase transition (Padmanabhan 1990).  The thermodynamical stability of self-gravitating systems can be studied by different technics. Horwitz \\& Katz (1978) use a field theory and write the density of states, the partition function and the grand partition function in MCE, CE and GCE respectively as a path integral for a formal field $\\phi$. In the mean-field approximation, the integral is dominated by the distribution $\\phi_{0}$ which {maximizes} a specific action $A\\lbrack \\phi\\rbrack$. In the grand canonical ensemble, $A_{GCE}\\lbrack \\phi\\rbrack$ is the Liouville action. Horwitz \\& Katz (1978) solve the problem numerically and find that the series of equilibria become unstable for a density contrast $1.58$, $32.1$ and $709$ in GCE, CE and MCE respectively. It is normal that the critical density contrast (hence the stability of the system) increases when more and more constraints are added on the system (conservation of mass in CE, conservation of mass and energy in MCE). This field theory has been rediscussed recently by de Vega \\& Sanchez (2002) who confirmed previous results and proposed interesting developements.  The thermodynamical stability of self-gravitating systems can also be settled by studying whether an isothermal sphere is a maximum or a saddle point of an appropriate thermodynamical potential: the entropy in MCE, the free energy in CE and the grand potential in GCE. The change of stability can be determined very easily from the topology of the equilibrium phase diagram by using the turning point criterion of Katz (1978) who has extended Poincar\\'e's theory on linear series of equilibria (see also Lynden-Bell \\& Wood 1968). This method is very powerful but it does not provide the form of the perturbation profile that triggers the instability. This perturbation profile can be obtained by computing explicitly the second order variations of the thermodynamical potential and reducing the problem of stability to the study of an eigenvalue equation. This study was first performed by Antonov (1962) in MCE and revisited by Padmanabhan (1989) with a simpler mathematical treatment.  This analysis was extended in CE by Chavanis (2002a) who showed in addition the equivalence between thermodynamical stability and dynamical stability with respect to Navier-Stokes equations (Jeans problem). Remarkably, this stability analysis can be performed analytically by using simple graphical constructions in the Milne plane.  In the present paper, we propose to extend these analytical methods to more general situations in order to provide a complete description of the thermodynamics of spherical self-gravitating systems. In Sec. \\ref{sec_thermo}, we review the stability limits of isothermal spheres in different ensembles by using the turning point criterion. In Sec. \\ref{sec_stab}, we consider specifically the grand canonical and grand microcanonical ensembles and evaluate the second order variations of the associated thermodynamical potential.  In Sec. \\ref{sec_field}, we briefly discuss the connexion between thermodynamics and statistical mechanics (and field theory) for self-gravitating systems. In Sec. \\ref{sec_tsallis}, we consider the case of generalized thermodynamics proposed by Tsallis (1988) and leading to stellar polytropes (Plastino \\& Plastino 1997, Taruya \\& Sakagami 2002a,b, Chavanis 2002b). In Sec. \\ref{sec_isobaric}, we discuss the stability of isothermal and polytropic spheres under an external pressure (Bonnor problem). We provide a new and entirely analytical solution of this old problem. Finally, in Sec. \\ref{sec_st} we discuss the relation between dynamical and thermodynamical stability of stellar systems and gaseous spheres. ",
        "conclusions": "\\label{sec_conc}  We have performed an exhaustive study of the thermodynamics of self-gravitating systems in various ensembles. This paper completes previous investigations on the subject and all ensembles have now been treated. Contrary to ordinary (extensive) systems, we have to perform a specific study in each ensemble since the stability limits differ from one to the other. Remarkably, the thermodynamical stability problem can be studied analytically or with simple graphical constructions. The dynamical stability of isothermal gaseous spheres can also be studied analytically both in Newtonian mechanics for the Euler-Jeans equations (Chavanis 2002a) and in general relativity for the Einstein equations (Chavanis 2002d).  These results can be of relevance both for astrophysicists and statistical mechanicians and could make a bridge between these two communities. On an astrophysical point of view, they show that we must be careful to precisely define the ensemble in which we work since they are not equivalent. This does not affect the structure of the equilibrium configuration but it may affect its stability. On a physical point of view, this study fills an important gap in the statisical mechanics literature since the case of self-gravitating systems is not discussed at all in standard textbooks of statistical mechanics and thermodynamics.  We have also discussed the relevance of Tsallis generalized thermodynamics for stellar systems. Collisionless stellar systems such as elliptical galaxies can achieve a metaequilibrium state as a result of a violent relaxation.  The Boltzmann entropy $S_{B}[\\overline{f}]$ is the correct entropy for these systems (in a coarse-grained sense and assuming that the system is non-degenerate). It measures the disorder where the disorder is equal to the number of microstates consistent with a given macrostate. However, $S_{B}[\\overline{f}]$ is not maximized by the system because of incomplete relaxation (this is independant on the fact that $S_{B}[\\overline{f}]$ has no maximum!). In any case, the state resulting from incomplete violent relaxation is a nonlinearly stable solution of the Vlasov equation (on a coarse-grained scale). If $\\overline{f}=f(\\epsilon)$, it maximizes a H-function $S[f]=-\\int C(f)d^{3}{\\bf r}d^{3}{\\bf v}$, where $C(f)$ is a convex function, at fixed mass and energy. Therefore, we can use a thermodynamical analogy to study the dynamical stability of collisionless stellar systems. Tsallis entropies $S_{q}[f]$ are a particular class of H-functions leading to stellar polytropes. Stellar polytropes do not give a good fit of elliptical galaxies. A better model of incomplete violent relaxation, motivated by precise physical arguments, consists of an isothermal core and a polytropic halo with index $n\\simeq 4$ (Hjorth \\& Madsen 1993). This can be fitted by  a composite Tsallis (polytropic) model with $q=1$ in the core (complete mixing) and $q\\simeq 7/5$ in the envelope (incomplete mixing). More generally, the state resulting from incomplete violent relaxation does not necessarily maximize a H-function at fixed mass and energy. Due to the Jeans theorem, the distribution function can depend on other integrals than the stellar energy or the Jacobi energy. It is possible, nevertheless, that relevant distribution functions arising from incomplete violent relaxation maximize a $H$-function at fixed mass, energy, angular momentum and additional (adiabatic) invariants.  \\vskip0.2cm \\noindent{\\it Acknowledgements.} I am grateful to J. Katz and T. Padmanabhan for stimulating discussions and encouragement.   \\appendix"
    },
    "0207/astro-ph0207049_arXiv.txt": {
        "abstract": "In this letter, we present the results of our study of galaxy-galaxy lensing in massive cluster-lenses spanning $z = 0.17$ to $0.58$, utilizing high-quality archival {\\it Hubble Space Telescope} ({\\it HST}\\,) data. Local anisotropies in the shear maps are assumed to arise from dark matter substructure within these clusters. Associating the substructure with bright early-type cluster galaxies, we quantify the properties of typical $L^*$ cluster members in a statistical fashion.  The fraction of total mass associated with individual galaxies within the inner regions of these clusters ranges from 10--20\\% implying that the bulk of the dark matter in massive lensing clusters is smoothly distributed.  Looking at the properties of the cluster galaxies, we find strong evidence ($>3$-$\\sigma$ significance) that a fiducial early-type $L^\\ast$ galaxy in these clusters has a mass distribution that is tidally truncated compared to equivalent luminosity galaxies in the field. In fact, we exclude field galaxy scale dark halos for these cluster early-types at $>10$-$\\sigma$ significance.  We compare the tidal radii obtained from this lensing analysis with the central density of the cluster potentials and find a correlation which is in excellent agreement with theoretical expectations of tidal truncation: $\\log [r_t*] \\propto (-0.6\\pm 0.2) \\log [\\rho_0]$. ",
        "introduction": "Galaxy-galaxy lensing provides a powerful tool to statistically measure the mass and the details of the mass distribution for field galaxies (Tyson et al.\\ 1984; Brainerd et al.\\ 1996). These studies confirm the existence of massive dark matter halos around typical field galaxies, extending to beyond 100\\,kpc\\footnote{We adopt h=H$_o$/100km\\,s$^{-1}$\\,Mpc$^{-1}$=0.5 and $q_o=0.5$, $\\Omega_o = 1$. Our results however, are not sensitive to values of the cosmological parameters.}(Brainerd et al.\\ 1996; Fischer et al.\\ 2000; Smith et al.\\ 2001b; McKay et al.\\ 2001).  The same technique can be modified and implemented within clusters to constrain the masses of cluster galaxies (Natarajan \\& Kneib 1997, NK97; Geiger \\& Schneider 1998). Successful application of the same to the rich, lensing cluster AC\\,114 at $z = 0.31$, suggests that the average $M/L$ ratio and spatial extents of the dark matter halos associated with early-type galaxies in such dense environments may differ from those of comparable luminosity field galaxies (Natarajan et al.\\ 1998, NKSE98).  The technique applied by NKSE98 quantifies the local weak distortions in the observed shear field of massive cluster-lenses, as perturbations arising from the massive halos of cluster galaxies (for details see NK97).  By associating these perturbations with bright early-type cluster members, the relative mass fraction in their halos is constrained using a combined $\\chi^2$-maximum likelihood method. The strength of this approach is the simultaneous use of constraints from the observed strong and weak lensing features.  The fractional mass in clusters associated with individual galaxy halos has important consequences for the frequency and nature of interactions (Moore et al.\\ 1996; Ghigna et al.\\ 1998; Okamato \\& Habe 1999; Merritt 1983). The theoretical expectation is that the global tidal field of a massive, dense cluster potential well should be strong enough to truncate the dark matter halos of galaxies that traverse the cluster core. In this letter, we test this expectation using well calibrated mass models for rich clusters at $z\\sim 0.17$--0.58 that utilize the observed strong lensing features -- positions, magnitudes, geometry of multiple images and measured spectroscopic redshifts as well as the shear field. ",
        "conclusions": "We have statistically extracted characteristic parameters for typical $L^\\ast$ cluster galaxies that inhabit massive, dense lensing cluster-lenses ranging in redshift from 0.17--0.58. This has been achieved by combining strong and weak lensing {\\it HST} observations in conjunction with an assumed parametric mass model.  We find that the inferred mass distribution of a fiducial $L^\\ast$ is extremely compact, although the inferred $r_t^\\ast$'s lie well outside the optical radii and correspond to roughly between 5--$10\\,R_e$.  Our analysis also shows that the halos of individual cluster galaxies contribute at most 10--20\\% of the total mass of the cluster within the central 1\\,Mpc, covered by the {\\it HST} {\\it WFPC2} imaging using the results of our likelihood analysis alongwith the best-fit paramters that characterize the smooth clump. Therefore, in the inner regions of these clusters the bulk of the dark matter is in fact smoothly distributed. Similar lensing studies of field galaxies, e.g.\\ Wilson et al.\\ (2001), typically find a non-zero signal for the radially averaged stacked tangential shear out to 200\\,kpc. In contrast our study of the halos of galaxies in clusters detects a finite $r_t^\\ast$, which we attribute to the tidal truncation induced by the motion of these cluster galaxies inside the potential well. From the contours in the likelihood plots, the presence of field galaxy scale dark halos can, in fact, be excluded at $>10$-$\\sigma$ significance.  The clusters we study here are all rich systems spanning a range in central density, which may explain why the best-fit values of $r_t^\\ast$ obtained vary by a factor of 2--3.  To test this suggestion we plot in Fig.~2 the variation of the central density of the cluster dark matter with $r_t^\\ast/\\sigma^\\ast$ based on our lens models and evaluated at the cluster core radius.  We see a good correlation and derive a best-fit slope of $-0.6 \\pm 0.2$. This compares well with the theoretical expected value from a tidal stripping model (Merritt 1983) of $-0.5$: \\begin{eqnarray} r_t^\\ast\\,\\,\\approx\\,\\,40\\,\\,(\\frac {\\sigma_{*}}{180{\\rm \\,km\\,s}^{-1}})\\, ({\\frac {\\rho_0(r_c)}{3.95 \\times {10^6}\\,M_\\odot{\\rm\\,kpc}^{-3}}})^{-\\frac{1}{2}}\\,\\,{\\rm kpc}. \\end{eqnarray}  Dark halos of the scale detected here indicate a high probability of galaxy encounters over a Hubble time within a rich cluster.  However, since the internal velocity dispersions of these cluster galaxies ($<250$\\,km\\,s$^{-1}$) are much smaller than their orbital velocities, these interactions are unlikely to lead to mergers, suggesting that the encounters of the kind simulated in the `galaxy harassment' picture (Moore et al.\\ 1996) are frequent and likely.  In fact, high resolution cosmological N-body simulations of cluster formation and evolution (Ghigna et al.\\ 1998; Moore et al.\\ 1996), find that the dominant interactions are between the global cluster tidal field and individual galaxies after $z = 2$. The cluster tidal field significantly tidally strips galaxy halos in the inner $0.5$\\,Mpc and the radial extent of the surviving halos is a strong function of their distance from the cluster center. Much of this modification is found to occur between $z = 0.5$--0. Detailed comparison of these results with tidal stripping of dark matter halos in cosmological N-body simulations will be presented in a forthcoming paper.  The prospects for extending this technique to larger scales within clusters in order to study the efficiency of halo stripping as a function of radius (variation of $r_t^\\ast$ as a function of radius) and morphological type are very promising with new instruments such as the {\\it Advanced Camera for Survey} on {\\it HST}. Multi-band imaging will enable photometric redshift determination for the background sources which will reduce one of the significant sources of noise for future analyses."
    },
    "0207/astro-ph0207401.txt": {
        "abstract": "An $R$ \\& $K^\\prime$ atlas of the {\\em IRAS} 1-Jy sample of 118 ultraluminous infrared galaxies (ULIGs) was presented in a companion paper (Kim, Veilleux, \\& Sanders 2002; Paper I). The present paper discusses the results from the analysis of these images supplemented with new spectroscopic data obtained at Keck.  All but one object in the 1-Jy sample show signs of a strong tidal interaction/merger. Multiple mergers involving more than two galaxies are seen in no more than 5 of the 118 ($<$ 5\\%) systems.  None of the 1-Jy sources is in the first-approach stage of the interaction, and most (56\\%) of them harbor a single disturbed nucleus and are therefore in the later stages of a merger.  Seyfert galaxies (especially those of type 1), warm ULIGs ($f_{25}/f_{60} > 0.2$) and the more luminous systems ($>$ 10$^{12.5}$ $L_\\odot$) all show a strong tendency to be advanced mergers with a single nucleus.  The individual galaxies in the binary systems of the 1-Jy sample show a broad distribution in host magnitudes (luminosities) with a mean of --21.02 $\\pm$ 0.76 mag. (0.85 $\\pm$ $^{0.86}_{0.43}$ $L^\\ast$) at $R$ and --23.98 $\\pm$ 1.25 mag. (0.90 $\\pm$ $^{1.94}_{0.61}$ $L^\\ast$) at $K^\\prime$, and a $R$- or $K^\\prime$-band luminosity ratio generally less than $\\sim$ 4.  Single-nucleus ULIGs also show a broad distribution in host magnitudes (luminosities) with an average of --21.77 $\\pm$ 0.92 mag. (1.69 $\\pm$ $^{2.25}_{0.97}$ $L^\\ast$) at $R$ and --25.03 $\\pm$ 0.94 mag. (2.36 $\\pm$ $^{3.24}_{1.38}$) at $K^\\prime$.  These distributions overlap considerably with those of quasars. The same statement applies to $R - K^\\prime$ colors in ULIG and quasar hosts.  An analysis of the surface brightness profiles of the host galaxies in single-nucleus sources reveals that about 73\\% of the $R$ and $K^\\prime$ surface brightness profiles are fit adequately by an elliptical-like $R^{1/4}$-law.  These elliptical-like 1-Jy systems have luminosity and $R$-band axial ratio distributions that are similar to those of normal (inactive) intermediate-luminosity ellipticals and follow with some scatter the same $\\mu_e - r_e$ relation, giving credence to the idea that some of these objects may eventually become intermediate-luminosity elliptical galaxies if they get rid of their excess gas or transform this gas into stars.  These elliptical-like hosts are most common among merger remnants with Seyfert 1 nuclei (83\\%), Seyfert 2 optical characteristics (69\\%) or mid-infrared ($ISO$) AGN signatures (80\\%).  The mean half-light radius of these ULIGs is $4.80 \\pm 1.37$ kpc at $R$ and 3.48 $\\pm$ 1.39 kpc at $K^\\prime$, typical of intermediate-luminosity ellipticals. These values are in excellent agreement with recent quasar measurements obtained at $H$ with $HST$, but are systematically lower than other $HST$ measurements derived at $R$. The reason for this discrepancy between the two quasar datasets is not known.  In general, the results from the present study are consistent with the merger-driven evolutionary sequence ``cool ULIGs $\\rightarrow$ warm ULIGs $\\rightarrow$ quasars.'' However, many exceptions appear to exist to this simple picture (e.g., 46\\% of the 41 advanced mergers show no obvious signs of Seyfert activity).  This underlines the importance of using a large homogeneous sample like the 1-Jy sample to draw statistically meaningful conclusions; the problems of small sample size and/or inhomogeneous selection criteria have plagued many studies of luminous infrared galaxies in the past. ",
        "introduction": "Local ultraluminous infrared galaxies (ULIGs; log [L$_{\\rm IR}$/$L_\\odot$] $\\ge$ 12; $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and $q_0$ = 0) represent some of the best laboratories to study in detail the violent aftermaths of galaxy collisions and their possible connection with quasars and normal (inactive) elliptical galaxies. Recent deep surveys with the Infrared Space Observatory ($ISO$) and sub-mm ground-based facilities have revealed several distant ($z \\approx$ 0.5 -- 4.0) infrared-luminous galaxies which appear to share many of the properties of local ULIGs ($ISO:$ Kawara et al. 1998; Puget et al. 1999; Matsuhara et al. 2000; Efstathiou et al. 2000; Serjeant et al. 2001; Sanders et al. 2002, in prep.; $SCUBA:$ Smail, Ivison, \\& Blain 1997; Hughes et al. 1998; Blain et al. 1999; Eales et al. 1999; Barger, Cowie, \\& Sanders 1999). This extragalactic population of high-$z$ infrared bright galaxies appears to dominate the far-infrared extragalactic background and is probably a major contributor to the overall star formation and metal enrichment history of the universe (e.g., Smail et al. 1997; Hughes et al. 1998; Barger et al. 1998; Genzel \\& Cesarsky 2000; Rupke, Veilleux, \\& Sanders 2002; Blain et al. 2002).  Two crucial questions need to be answered to properly address the issue of the origin and evolution of ULIGs: (1) What is the dominant energy source in ULIGs: Starbursts or active galactic nuclei (AGNs)? (2) Is the dominant energy source a function of the interaction/merger phase in these systems?  Considerable progress has been made in recent years in answering the first of these questions. Ground-based optical and near-infrared spectroscopic survey of the 1-Jy sample has shown that at least 25 -- 30\\% of ULIGs show genuine signs of AGN activity (Kim, Veilleux, \\& Sanders 1998; Veilleux, Kim, \\& Sanders 1999a; Veilleux, Sanders, \\& Kim 1997, 1999b; see also Goldader et al. 1995; Murphy et al. 2001b). This fraction increases to 35 -- 50\\% among the objects with log [$L_{\\rm IR}$/$L_\\odot$] $\\ge$ 12.3. Comparisons of the dereddened emission-line luminosities of the BLRs detected at optical or near-infrared wavelengths in the ULIGs of the 1-Jy sample with those of optical quasars indicate that the AGN/quasar in ULIGs is the main source of energy in at least 15 -- 25\\% of all ULIGs in the 1-Jy sample. This fraction is closer to 30 -- 50\\% among ULIGs with $L_{\\rm IR} > 10^{12.3}\\ L_\\odot$. These results are compatible with those from recent mid-infrared spectroscopic surveys carried out with {\\it ISO} (e.g., Genzel et al. 1998; Lutz et al. 1998; Rigopoulou et al. 1999; Tran et al. 2001). Indeed, a detailed object-by-object comparison of the optical and mid-infrared classifications shows an excellent agreement between the two classification schemes (Lutz, Veilleux, \\& Genzel 1999).  These results suggest that strong nuclear activity, once triggered, quickly breaks the obscuring screen at least in certain directions, thus becoming detectable over a wide wavelength range.  Much of the research effort now focusses on answering the second, more difficult question of a possible dependence of the energy source on the interaction/merger phase. A large dataset already exists in the literature on the morphology of luminous and ultraluminous infrared galaxies. Optical studies have shown that the fraction of strongly interacting/merger systems increases with increasing infrared luminosities, reaching $>$ 95\\% among the ultraluminous systems (e.g., Sanders et al. 1988a; Melnick \\& Mirabel 1990; Murphy et al. 1996; Clements et al. 1996; although see Lawrence et al. 1989; Zou et al. 1991; Leech et al. 1994). The improved angular resolution and lower optical depth in the near-infrared as compared to the optical provides a cleaner view of the nuclear stellar distribution of ULIGs. The infrared morphologies of ULIGs often look significantly different from the optical images (e.g., Carico et al. 1990; Graham et al. 1990; Eales et al. 1990; Majewski et al. 1993).  Over the past few years, high-resolution imaging with {\\em HST} has contributed significantly to our knowledge of ULIGs (e.g., Surace et al. 1998; Zheng et al. 1999; Scoville et al. 2000; Borne et al. 2000; Cui et al. 2001; Farrah et al. 2001; Colina et al. 2001; Bushouse et al. 2002). These studies have had great success characterizing the central cores and massive stellar clusters in these objects, but the images often do not reach faint enough flux limits to fully characterize the host galaxies and associated tidal features. Adaptive optics imaging with ground-based telescopes offers a promising new way to acquire deep high-resolution of ULIGs, but this technique has so far been used for only a limited number of objects (e.g., Surace \\& Sanders 1999, 2000; Surace, Sanders, \\& Evans 2000, 2001).  In Kim, Veilleux, \\& Sanders (2002; Paper I), we presented an $R$ \\& $K^\\prime$ atlas of the {\\em IRAS} 1-Jy sample of 118 ULIGs.  This large and homogeneous dataset on the nearest (median redshift = 0.145) and brightest such objects in the universe is particularly well suited to address the issue of the origin and evolution of ULIGs (see Kim 1995 and Kim \\& Sanders 1998 for a detailed description of the 1-Jy sample; note that Galactic extinction is negligible for the 1-Jy sample since $\\vert b \\vert$ $>$ 30$^\\circ$ for all these sources). Results derived from the 1-Jy sample should serve as a good local baseline for studies on distant luminous infrared galaxies planned in the years to come with $SIRTF$, $SOFIA$, and other ground-based infrared and submm facilities.  In the present paper, we analyze these images and combine the results of this analysis with those derived from our optical and near-infrared spectroscopic data to look for possible trends with optical and near-infrared morphological and spectrophotometric parameters.  When possible, the results from our analysis of the ground-based data on the 1-Jy sample are combined with published results obtained at other wavelengths and compared with the quickly growing set of high-quality data on optical and infrared-bright QSOs (e.g., McLure et al. 1999; M\\'arquez et al. 2000; McLeod \\& McLeod 2001; Canalizo \\& Stockton 2000a, 2000b, 2001; Surace et al. 2001; Dunlop et al. 2002 and references therein; see Stockton 1999 for a review of the data before 1999).  Given the high frequency of mergers among ULIGs, the data on the 1-Jy sample also allow us to examine the disk-disk merger scenario for elliptical galaxy formation (e.g., Toomre 1977; Schweizer 1982; Wright et al. 1990; Scoville et al. 1990; Stanford \\& Bushouse 1991; Kormendy \\& Sanders 1992; Doyon et al. 1994; Genzel et al. 2001). The results of our analysis are described in \\S\\S 2 -- 4. These results are then used in \\S 5 to discuss the origin and evolution of ULIGs and their possible evolutionary link with quasars and elliptical galaxies. A summary is presented in \\S 6. Each object in the 1-Jy sample is described in detail in an Appendix. Preliminary results of this study were presented in Veilleux (2001).   We adopt $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and $q_0$ = 0.0 throughout this paper and have converted the results from published papers to this cosmology to facilitate comparisons. We also compare our absolute magnitudes with the magnitudes corresponding to a $L^\\ast$ galaxy in a Schechter function description of the local field galaxy luminosity function. The reader should be cautious when comparing luminosities from various papers as they may be based on different definitions of $L^\\ast$. The values of the absolute magnitudes of a $L^\\ast$ galaxy in various bands, $M^\\ast$, are now better constrained than in the past thanks to large-scale surveys such as the 2dF Galaxy Redshift Survey (2dFGRS). Here we adopt $M^\\ast_R = -21.2$ mag. and $M^\\ast_{K^\\prime} = -24.1$ mag.  The value of $M^\\ast_R$ was derived using a $L^\\ast$ magnitude of $M_r = -20.9$ mag. based on the results from the Las Campanas Redshift Survey (Lin et al. 1996) and assuming a typical $r - R$ = 0.3 mag.  for a moderately old stellar population at $z \\approx 0.1$ (Surace 2002, private communication). A similar value for $M^\\ast_R$ is derived using $M^\\ast_B = -20.1$ mag. (e.g, Mobasher et al. 1993; Zucca et al. 1997; Ratcliffe et al. 1998; Folkes et al. 1999) and adopting $B - R \\approx 1.0$ mag. for typical field galaxies at $z \\approx 0.1$ (e.g. Fukugita et al. 1995), but keeping in mind that this color term is strongly dependent on morphological type and more specifically on the level of star formation activity in the galaxy. The $K^\\prime$ magnitude is based directly on the results of Cole et al. (2001) from an infrared-selected subsample of the 2dFGRS. These results are also consistent to within $\\pm$ $\\sim$ 0.2 mag. with the numbers quoted from optically selected samples (e.g., Mobasher, Sharples, \\& Ellis 1993; %Szokoly et al. 1998; Loveday 2000) and from the K-band surveys by Glazebrook et al. (1995) and Gardner et al. (1997), after taking into account the various $k-$ and photometric corrections mentioned in Table 3 of Cole et al. (2001). ",
        "conclusions": "An $R$- and $K^\\prime$-band atlas of the {\\em IRAS} 1-Jy sample of 118 ULIGs was presented in a companion paper (Kim, Veilleux, \\& Sanders 2002). The present paper discusses the results from the analysis of these images and combines them with the results from published spectroscopic studies of ULIGs at optical, near-infrared, and mid-infrared wavelengths as well as new Keck spectroscopy. The results on the 1-Jy sample are compared with those from optical and near-infrared studies of quasars and normal ellipticals.  The main conclusions are as follows:  \\begin{itemize}  \\item[1.] All but one object in the 1-Jy sample show signs of a strong tidal interaction/merger in the form of distorted or double nuclei, tidal tails, bridges, and overlapping disks.  Interactions involving more than two galaxies are seen in only 5 (4\\%) of the 118 systems. These results confirm those of previous studies.  \\item[2.] Objects with red $R - K^\\prime$ colors and large nuclear-to-total luminosity ratios show a tendency to host Seyfert nuclei, have warm $IRAS$ 25 $\\mu$m/60~$\\mu$m colors and be AGN-dominated according to the $ISO$ classification scheme.  \\item[3.] Using a classification scheme first proposed by Surace (1998) and based on the results of published numerical simulations of the mergers of two galaxies, we classified the 1-Jy sources according to their overall morphology and apparent nuclear separations.  None of the 1-Jy sources appears to be in the early stages (first approach or first contact) of a merger. Most (56\\%) of them harbor a single disturbed nucleus with and without tidal tails; they are therefore in the late stages of a merger.  The fraction of advanced mergers with a single nucleus increases above infrared luminosities of 10$^{12.5}$ $L_\\odot$ and decreases below 10$^{12}$ $L_\\odot$.  \\item[4.]  The strengths of the H$\\beta$ and Mg Ib stellar features measured in the nuclei of ULIGs are not particularly good indicators of the merger phase or epoch of the merger event.  \\item[5.] All Seyfert 1s and most of the Seyfert 2s are advanced mergers, either based on their overall morphology or their small ($<$ 5 kpc) nuclear separations.  A similar result is found when we consider the warm objects with $f_{20}/f_{60}$ $>$ 0.2 or the AGN-dominated objects based on $ISO$ classification. LINERs and H~II region-like galaxies show no preference between pre-merger and advanced merger phases.  \\item[6.]  The individual galaxies making up the binary systems of the 1-Jy sample show a broad distribution in host absolute magnitudes (luminosities) with a mean of --21.02 $\\pm$ 0.76 mag. (0.85 $\\pm$ $^{0.86}_{0.43}$ $L^\\ast$) at $R$ and --23.98 $\\pm$ 1.25 mag. (0.90 $\\pm$ $^{1.94}_{0.61}$ $L^\\ast$) at $K^\\prime$, and a luminosity ratio at $R$ or $K^\\prime$ generally less than $\\sim$ 4.  The hosts of single-nucleus ULIGs have mean absolute magnitudes (luminosities) of --21.77 $\\pm$ 0.92 mag. (1.69 $\\pm$ $^{2.25}_{0.97}$ $L^\\ast$) at $R$ and --25.03 $\\pm$ 0.94 mag. (2.36 $\\pm$ $^{3.24}_{1.38}$) at $K^\\prime$.  These magnitudes are similar to those found in previous studies of ULIGs, except those derived from shallow $HST$ snapshots which underestimate the contribution from low surface brightness features. Correlations are observed between host galaxy luminosity and infrared luminosity, $IRAS$ 25 $\\mu$m/60~$\\mu$m color, and optical spectral type. The trend with infrared luminosity is due to a redshift bias. There is considerable overlap between the host galaxy luminosity distribution of single-nucleus ULIGs and that of quasars, although the hosts of the quasars studied by Dunlop et al.  (2002) are slightly more luminous on average than the 1-Jy ULIG hosts.  This luminosity shift is not observed when comparing with the results of McLeod \\& McLeod (2001) and Surace et al. (2001). The $R - K^\\prime$ colors of ULIG hosts are similar to those of quasars.  \\item[7.]  An analysis of the surface brightness profiles of the host galaxies in single-nucleus ULIGs reveals that about 35\\% and 2\\% of the $R$ and $K^\\prime$ surface brightness profiles are fit adequately by an elliptical-like $R^{1/4}$-law and an exponential disk, respectively. Another 38\\% are equally well fit by either an exponential or an elliptical-like profile. The remainder (26\\%) of the single-nucleus sources cannot be fit with either one of these profiles.  Combining these results, we find that a de Vaucouleurs profile is an adequate fit to $\\sim$ 73\\% of the single-nucleus ULIGs in the 1-Jy sample. These elliptical-like hosts are most common in merger remnants with Seyfert 1 nuclei (83\\%), Seyfert 2 characteristics (60\\%) or mid-infrared ($ISO$) AGN signatures (80\\%).  \\item[8.] The hosts of ULIGs have half-light radii ($<r_e> = 4.80 \\pm 1.37$ kpc at $R$ and $<r_e>$ = 3.48 $\\pm$ 1.39 kpc at $K^\\prime$) which are similar to those measured by McLeod \\& McLeod (2001) in quasar hosts, but are significantly smaller than the quasar hosts studied by Dunlop et al. (2002). The origin of this apparent discrepancy between the two quasar datasets is not clear at present. The hosts of 1-Jy systems follow with some scatter the $\\mu_e - r_e$ relation of normal ellipticals (especially if Seyfert 1s are excluded from the analysis due to possible AGN residuals). The distributions of luminosities, $R$-band axial ratios, and half-light radii in single-nucleus ULIGs are also similar to those of normal ellipticals of intermediate luminosities. The results at $K^\\prime$ are more uncertain because of the smaller sample size and possible PSF subtraction residuals. Elliptical-like hosts in the 1-Jy sample show a broader range of boxiness values than normal ellipticals.  \\item[9.] Results \\#7 and \\#8 provide strong support to the idea that some of the ultraluminous infrared mergers in the 1-Jy sample may eventually become intermediate-luminosity elliptical galaxies under the condition that they get rid of their excess gas or transform this gas into stars. The significant fraction of single-nucleus ULIGs with ambiguous surface brightness profiles or with large boxiness parameters indicate that these objects are still feeling the effects of the recent mergers.  These results confirm the predictions of numerical simulations that violent relaxation is very efficient at the center of the merger but less so in the outer regions.  \\item[10.] The results from this study are generally consistent with the evolutionary scenario in which ULIGs are the results of a merger of two gas-rich galaxies which first goes through a starburst-dominated pre-merger phase when the system is seen as a binary, next reaches a dust-enshrouded AGN-dominated merger phase once the two nuclei have merged into one, and then finally ends up as an optical (post-ULIG) AGN where the host is elliptical-like and shows only limited signs of the merger.  However, our results also indicate that many ULIGs in the 1-Jy sample may not follow this exact scenario. For instance, approximately 46\\% of the 41 advanced mergers in the 1-Jy sample show no obvious signs of Seyfert activity, while seven of the 45 pre-mergers already show Seyfert 2 (but not Seyfert 1) activity. The possible correlations between host luminosities, $IRAS$ 25 $\\mu$m/60~$\\mu$m color, and optical spectral types may add another twist to the merger scenario, suggesting that merger-induced quasar activity may require the merger of massive ($\\ga$ $L^\\ast$ + $L^\\ast$) galaxies.  These trends may also be due to a brightening of the circumnuclear starburst in the hosts of the warm Seyfert ULIGs. However, one needs to be cautious when interpreting these apparent trends because of the difficulty associated with removing the contribution of the central AGN in some of the warm Seyfert ULIGs. \\end{itemize}  Deep near-infrared $HST$ and adaptive-optics imaging of ULIGs and quasars will help further clarify the nature of their host galaxies. These observing techniques should be applied to a large sample of warm ULIGs to refine the PSF subtraction in these objects and verify the possible trends with host luminosities reported in the present paper. The spatial resolution of the data presented here is not sufficient to determine the core properties of the elliptical-like ULIG hosts and the bulge-to-disk ratio of the galaxies in the pre-merger phase.  Both high spatial resolution and large dynamic range will be needed to fully characterize the host galaxies of these ULIGs. A similar program should be carried out on the classical quasars to help us understand the apparent inconsistencies between the various quasar datasets. Finally, three-dimensional spectrographs on large telescopes will provide the sensitivity and two-dimensional spatial coverage needed to derive detailed kinematic maps in several elliptical-like ULIG hosts. Comparisons with data on normal old (low-$z$) and young (high-$z$) ellipticals will help quantify the similarities and differences between these two classes of objects.  \\clearpage"
    },
    "0207/astro-ph0207224_arXiv.txt": {
        "abstract": "Most astrophysical plasmas entail a balance between ionization and recombination. We present new results from the powerful R-matrix method which yields in an ab initio manner: (I) self-consistent photoionization and recombination cross sections using identical wavefunction expansions, (II) unified e-ion recombination rates at all temperatures of interest, incorporating non-resonant and resonant processes, radiative and dielectronic recombination (RR and DR), and (III) level-specific recombination rates for many excited atomic levels. RR and DR processes can not be measured or observed independently. In contrast to the shortcomings of simple (but computationally easy) approximations that unphysically treat RR and DR separately with different methods, emphasizing marginal effects for selected ions over small energy ranges, the R-matrix method naturally and compeletely accounts for e-ion recombination for all atomic systems.  Photoionization and recombination cross sections are compared with state-of-the-art experiments on synchrotron radiation sources and ion storage rings. Overall agreement between theory and experiments is within 10-20 \\%. For photoionization the comparison includes not only the ground state but also the metastable states, with highly resolved resonance structures. The recent experiments therefore support the estimated accuracy of the vast amount of photoionization data computed under the Opacity Project (OP), the Iron Project (IP), and related works using the R-matrix method. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207154_arXiv.txt": {
        "abstract": "Optically thin and geometrically thick accretion flows are known to be responsible for the observed radio/X-ray luminosity relation of the X-ray bright galactic nuclei. It has also been suggested that supermassive black hole masses can be estimated  from measurements of the core radio luminosity and the X-ray luminosity by using the advection-dominated accretion flow (ADAF) model. In this study we increase the number of data available by compiling the radio/X-ray fluxes and the mass in published literatures, and compare the observed ratio of the luminosities with predictions from various models of optically thin accretion flows. Semi-analytically derived relations of the luminosities are presented in cases of the standard ADAF model and modified ADAF models, in which a truncation of inner parts of the flows and winds causing a reduction of the infalling matter are included. We show that the observed relation can be used indeed to estimate the supermassive black hole mass, provided that properties of such accretion flows are known. Having investigated sensitivities of the method on modifications of the 'standard' ADAF model, we find that a general trend of model predictions from the 'standard' ADAF, the truncated ADAF and the 'windy' ADAF are somewhat indistinguishable. We also find, however, that the extreme case of the windy model  is inconsistent with currently available observational data, unless microphysics parameters are to be substantially changed. High resolution radio observations, however, are required to avoid the contamination of non-disk components, such as, a jet component, which, otherwise, results in the over-estimated SMBH mass. ",
        "introduction": "Supermassive black holes (SMBHs) have been considered as the most likely power sources of the activity in quasars and Active Galactic Nuclei  (Lynden-Bell 1969; Rees 1984). SMBHs at the centers of all galaxies are now recognized as ubiquitous, whose mass  $M_{\\rm SMBH}$ is proportional to the spheroidal bulge mass of the host galaxy or the galactic bulge luminosity (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Richstone et al. 1998) and is strongly correlated with the velocity dispersion of the host galaxy (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a; Ferrarese 2002). Recently, evidence for the existence of a SMBH in the center of our Galaxy has been added (Eckart \\& Genzel 1997; Genzel et al. 1997; Ghez et al. 1998). Though searches for SMBHs are primarily based on spatially resolved kinematics, another possible way to infer the presence of SMBHs and to shed light on the physical condition is to examine spectral energy distribution over a wide range from the radio to the hard X-ray frequencies. This emission spectrum is produced by an accreting matter, as the surrounding gas accretes onto the central SMBH (e.g., Frank et al. 1992; Ho 1999).  Several attempts have been made in utilizing this fact to estimate SMBH masses (Yi \\& Boughn 1998, 1999; Franceschini et al. 1998; Lacy et al. 2001; Ho 2002; Alonso-Herrero et al. 2002; Cao \\& Jiang 2002). It has been also suggested that advection-dominated accretion flows (ADAFs) are responsible for the observed radio and X-ray luminosities of some of the X-ray bright galactic nuclei (Fabian \\& Rees 1995; Di Matteo \\& Fabian 1997; Yi \\& Boughn 1998, 1999). Moreover, Yi \\& Boughn  (1999) discussed that the observed radio/X-ray luminosity relation in terms of the 'standard' ADAF model and demonstrated that it could be used as a relatively effective and consequently suitable tool to estimate unknown black hole masses. In this study we further explore this possibility both by increasing the number of data available and by deriving similar relations for the models with additional features, such as, truncations, winds, and investigate sensitivities of the method Yi \\& Boughn suggested on modifications of the 'standard' ADAF model.  This paper begins with model descriptions for the standard ADAF model and modified versions of the model in $\\S$  2. We derive analytical expressions to describe the radio/X-ray  luminosities for various models. We present the calculated ratios of the luminosities and compare them with the observational data in $\\S$ 3. We discuss and conclude in $\\S$ 4. ",
        "conclusions": "One of important predictions of the ADAF model is the radio/X-ray luminosity relation which can be  used to estimate central SMBH masses. We compute the theoretical predictions of the radio/X-ray luminosity  relation for the ADAF models with  various modifications and compare the predicted relation with that from the available SMBH candidates. Several nearby extra galaxies have consistent ratios of radio to X-ray luminosities with the ADAF predictions for an estimated mass of the central SMBH. If the type of the accretion flow is provided by observations with a high angular resolution, the inherent radio/X-ray luminosity relation provides a direct estimate of the central SMBH mass. We show the logarithm of the ratio of radio to X-ray luminosities against the logarithm of the X-ray luminosity for several SMBH masses in Figure 1. The observations for several extra galactic objects are consistent with the predictions of  ADAF models, as can be checked in Figures 2 and 3. It is, however, interesting to note that the model representing the  ADAFs with the strong wind can be ruled out unless the microphysics in the standard ADAF model has to be modified significantly.  In this study, we implicitly assume that all the SMBHs are non-rotating. When the black hole rotates, the last stable circular orbit is no longer 3 $r_g$. It becomes smaller than this value and eventually the event horizon when the black hole spin parameter $a \\sim 1$. If this is the case the boundary condition should change as well as the dynamical properties of the flows in inner parts (Gammie \\& Popham 1998; Popham \\& Gammie 1998). Besides dynamics of the flows, a rotating black hole may power a jet by the Blandford-Znajek process (Blandford \\& Znajek 1977). Therefore, even if the black hole rotates slowly, the rotation should be properly taken into account in order to make this method robust. Another crucial modification could be made in terms of microphysics we simplified. We have assumed that the viscously generated energy mainly heats the ions, heat exchanges between ions and electrons occur only due to the Coulomb interaction, other physical processes, such as, MHD turbulence, pair production, are ignored. Many works have been done to quantify the effects of such processes (e.g., \\\"{O}zel et al. 2000). Once clear understanding is incorporated in the method we present here, its implication can be more robust. On the other hand, at the moment, the peak luminosity of the radio is unlikely to be modified though the 'wings' of the radio spectrum are modified by the nonthermal electrons. Therefore, as long as we use the peak radio luminosity instead of whole structure, this method is not seriously defected."
    },
    "0207/astro-ph0207362_arXiv.txt": {
        "abstract": "From ROSAT data, the bright X-ray cluster Abell 85 was found to show a roughly symmetrical shape, on to which are superimposed several features, among which: (1) a blob or group of galaxies falling on to the main cluster, with the gas in the impact region probably hotter and the galaxies in that zone showing enhanced star formation; (2) a filament, either diffuse or made of small groups of galaxies, extending at least 4 Mpc from the cluster (Durret et al. 1998). Preliminary results obtained from XMM-Newton and Chandra observations of Abell 85 and its filament will be presented. ",
        "introduction": "The hot ($1 \\la T \\la 10\\,$keV; central density, $n_{0}\\approx 10^{-3} ${cm}${}^{-3}$) X-ray emitting gas found in rich clusters of galaxies is an excellent tool to probe the cluster dynamics, morphology, and history.  We present here an X-ray study of Abell 85 ($z = 0.0556$, richness class $R=0$, B-M type III) using new XMM-Newton data obtained in January 2002 and public data available from the Chandra archive.  At the distance of Abell 85, 1 arcmin corresponds to $65 h_{70}^{-1}$kpc (assuming $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$). ",
        "conclusions": "The high resolution and sensibility of XMM-Newton and Chandra satellite provides us with a new vision of an old cluster (quoting a meeting in Marseille five years ago...):  \\begin{itemize} \\itemsep = 0pt \\item The X-ray emission in the central region is not totally homogeneous or symmetric: the maximum is displaced relatively to the outer isophotes, a ``hole'' seems to exist south of the center;  \\item The temperature is cooler and the metallicity is higher in the central regions of the cluster than was previous measured;   \\item The existence of an X-ray filament, or at least diffuse X-ray emission south east of the cluster is confirmed; The south blob is definitely not a relaxed structure.  \\end{itemize}"
    },
    "0207/astro-ph0207648_arXiv.txt": {
        "abstract": "We report proper motion measurements of water masers in the massive-star forming region W~51A and the analyses of the 3-D kinematics of the masers in three maser clusters of W51A (W51 North, Main, and South). In W~51 North, we found a clear expanding flow that has an expansion velocity of $\\sim$70\\kms\\  and indicates deceleration. The originating point of the flow coincides within 0\\arcsec.1 with a silicon-monoxide maser source near the H{\\rm II} region W~51d. In W51 Main, no systematic motion was found in the whole velocity range (158\\kms\\   $\\leq\\; $\\vlsr  $<$ $-$58\\kms) although a stream motion was reported previously in a limited range of the Doppler velocity (54\\kms\\   $\\leq\\; $\\vlsr  $<$ 68\\kms). Multiple driving sources of outflows are thought to explain the kinematics of W51 Main. In W51 South, an expansion motion like a bipolar flow was marginally visible. Analyses based on diagonalization of the variance-covariance matrix of maser velocity vectors demonstrate that the maser kinematics in W~51 North and Main are significantly tri-axially asymmetric. We estimated a distance to W51 North to be 6.1$\\pm$1.3 kpc on the basis of the model fitting method adopting a radially expanding flow. ",
        "introduction": "\\label{sec:introduction}  Water maser emission is one of the most important phenomena in the study of star formation, often based on data obtained using very long baseline interferometry (VLBI) with high angular and velocity resolution (e.g., \\cite{rei81,eli92}). Analyses of spatial positions, Doppler velocities, and proper motions of individual maser features with a typical size of 1 AU \\citep {rei81} have revealed the 3-D gas kinematics around young stellar objects (YSOs) (e.g., \\cite{gen81a}a, hereafter G81; \\cite{gen81b}b; \\cite {sch81}, hereafter S81; \\cite{rei88,gwi92,cla96,fur00,ima00}, hereafter Paper {\\rm I}; \\cite{tor01}). In practice, details of the gas kinematics are complicated but seem to depend mainly on the evolutionary status of YSOs. Water maser sources sometimes enable measurement of the internal motions of giant molecular clouds by measuring relative bulk motions between clusters of water masers that are very close to each other (Paper {\\rm I}; \\cite{tor01}). Such bulk motions may be owe to, e.g.,  propagation of shock layers from newly-formed H{\\rm II} regions, cloud contraction by the self-gravitation due to the huge mass of a giant molecular cloud.  The massive-star forming region W~51A contains at least five independent clusters of water masers: W51 North, West, Main, and two clusters in South (e4 and e3). All of the maser clusters are independently associated with H{\\rm II} regions and dense and young cloud cores exhibiting several species of molecular emission (\\cite{gen78, dow79}; G81; S81; \\cite{gau87, gau93, zha95}, 1997; \\cite{zha98,lep98}, hereafter LLD; \\cite{eis02}, hereafter EGHMM). In W51N, silicon-monoxide masers have been detected (\\cite{mor92}; EGHMM). Proper motions of these water masers were measured two decades ago (G81; S81). However, details of the kinematics are still obscured because of limited numbers of measured proper motions ($\\leq$30, S81; G81). In W51 North, a separation motion has been recognized between two dominant \"sub-clusters\" of maser features (see Sect.\\ \\ref{sec:results-summary}), so-called the \"NW Cluster\" and the \"Dominant Center Reference Cluster\" (S81). EGHMM reconfirmed the similar separation motion in the patterns of the sub-clusters, which are similar to bow shocks, by comparing the patterns observed in 1998 with those of S81. In W51 Main, LLD found a stream from the \"Double-Knot\" sub-cluster in the south-east direction. The observed velocity range, however, was very limited (54\\kms\\   $<$ \\vlsr  $<$ 68\\kms). Throughout the entire range ($-$50\\kms\\   $\\leq$ \\vlsr  $\\leq$ 130\\kms) the kinematics are predominated by random motions (G81). In other maser clusters, their kinematics have never been well understood.  Here, we report monitoring observations of the W~51A water masers with the Japanese domestic VLBI network (J-Net) \\footnote {J-Net includes the 45-m telescope of Nobeyama Radio Observatory (NRO), which is a branch of the National Astronomical Observatory, an interuniversity research institute operated by the Ministry of Education, Culture, Sports, Science and Technology.} \\citep{omo94}. Sect.\\ 2 describes the VLBI observation and data reduction. Sect.\\ 3 summarizes the revealed 3-D kinematics of the individual clusters of water masers. Sect.\\ 4 discusses the origins and related issues of the maser kinematics. The estimation of the distance to W~51A and relative 3-D bulk motions among the maser clusters are also described. ",
        "conclusions": "Figure \\ref{fig:PM-W51N} shows several examples of measured relative proper motions of water maser features. Maser features fundamentally seem to move with constant velocities. The deviations from fit lines assuming constant velocity motions are within several tenths of a milliarcsecond. Some features have a large deviation from the fit lines because they are located together with other features within a small range, 1 mas in space and 1\\kms\\   in velocity, in which we were not able to correctly trace the same feature from one epoch to another.  Tables \\ref{tab:pmotionsN}, \\ref{tab:pmotionsM}, and \\ref{tab:pmotionsS} give parameters of maser features with measured proper motions. The numbers of measured proper motions were 123, 48, and 10 in W51 North, Main, and South (e4), respectively, which are larger than those of previous observations for W51 North by S81 and for W51 Main and South (e4) by G81. An important difference between the previous and the present measurements is the difference in time separations between the successive observing epochs: two years and, at minimum, only one month, respectively.  Even with a much shorter time separation of our observations, it has not been possible to measure proper motions in a large fraction of detected maser features ($>$ 50 \\%) mainly because of growth and decay of maser features among the observing epochs. These results imply that many of the maser features have lifetimes shorter than 1--2 months. Unfortunately, no maser proper motion has been identified in the W 51 South (e3) region.  Usually, each of water maser clusters consists of several groups of maser features with a size of 100--1000 AU. In this paper, we define such a group of features as a \"sub-cluster\".   \\subsection{W51 North} \\label{sec:W51N} \\subsubsection{Overview} \\label{sec:W51Noverview}  Figure \\ref{fig:W51N-color} presents the angular distribution of water masers in W51 North and West, which revives those found by previous observations (\\cite{gen78,dow79}; S81; EGHMM). Figure \\ref {fig:3D-W51N} presents the detailed 3-D motions of water masers around the two dominant sub-clusters, the red-shifted (north--west) and the blue-shifted (south--east, Dominant Center Reference) sub-clusters, which clearly exhibit a bipolar expanding flow. We found a \"bow shock\" pattern in the SE sub-cluster, which opens toward the NW sub-cluster and has been found since 1977 (S81; EGHMM). On the basis of the locations of the two sub-clusters and this bow-shock pattern, we estimated a location of SiO maser emission in W51N on our map with an uncertainty of less than 100 mas, whose position was measured by EGHMM with respect to the water masers with an uncertainty of less than 50 mas. The originating point of the outflow seems to be located around the middle of the two maser sub-clusters and to be roughly coincident with the location of the SiO maser. While only a separation motion between the two sub-clusters has been confirmed by S81 and EGHMM, the present result reveals that the individual sub-clusters themselves are also expanding.  On the other hand, motions of maser features far from the two sub-clusters do not exhibit any systematic motion but random motions with velocities up to 100\\kms. Some peculiar motions are found in the two dominant sub-clusters and are not due to misidentification of the proper motions. Thus, although the expanding flow has been visible, the kinematics of the W51 North region is heavily disordered dynamically by fast random motions. The details of the maser kinematics around W51 West were also obscured because only two proper motions were measured. It is difficult to compare the maser distribution with that of S81 because of too sparse time separation ($\\sim$20 yr).  \\subsubsection{Model fitting the maser kinematics} \\label{sec:model-fit} In order to estimate kinematical parameters of the expanding flow and a distance to W51N, we made model-fitting analyses for the 3-D motions of maser features. The procedure was performed on the basis of the least-square fitting of the observed kinematics to a radially-expanding flow model and in almost the same way as that applied to the W3 IRS~5 water masers (Paper {\\rm I}), which is not repeated here. One difference is only the assumed speed of the radial expansion of a maser feature {\\it i}, $V(i)$, as a function of distance of a maser feature from the originating point of the outflow, $r_{i}$, which is expressed more simply as $V_{\\mbox{exp}}(i)\\;=\\;V_{0}(r_{i}/r_{0})^{\\alpha}$, and where $V_{0}$ is an expansion velocity at a unit distance $r_{0}$, $\\alpha$ is a power-law index indicating the apparent acceleration of the flow. We made the fitting step-by-step, excluding maser features having unreliably large positive or negative expansion velocities or distances from the outflow origin. After such several iterations, we used 68 proper motion data and obtained best solutions that are given in Table \\ref{tab:w51model-fit}.  Figure \\ref{fig:3D-W51N} shows an estimated position of the outflow origin in the maser motions (plus), which coincides with the location of the SiO maser emission (filled square, EGHMM) within the position uncertainly ($\\sim$100 mas). A systemic line-of-sight velocity of the flow is almost equal to that of the ambient molecular cloud ($\\simeq$56\\kms, e.g., S81; \\cite {cox87, rud90, zha95, zha98,oku01}) and roughly coincident with that of the SiO maser emission ($\\simeq$47\\kms) within a velocity width of the cloud ($\\simeq$26\\kms).  Figure \\ref{fig:W51N-3DV} presents the maser feature motions projected onto three different planes. The maser kinematics indicate no rotation of the expanding flow, suggesting that ballistic motions predominate the kinematics. The best fit model and an expansion velocity plot against distance from the outflow origin (Figure \\ref{fig:expansion}) indicate that the expanding flow decelerates in the water maser region ($r=$ 200--500 mas or 1200--3000 AU from the outflow origin), where the expansion velocity decreases from $\\simeq$90\\kms\\   to $\\simeq$50\\kms. This is a controversial case against these expanding flows that apparently exhibit the accelerations in water maser kinematics (Orion KL, \\cite{gen81b}b; W~49N, \\cite{gwi92}; W3 IRS~5, Paper {\\rm I}).  \\subsection{W51 Main} \\label{sec:W51M}    Figure \\ref{fig:W51M-color}a presents the angular distribution of water masers in W51 Main, which also revives those found by previous observations (\\cite{gen78, gen79}; G81) shown in Figure \\ref {fig:W51M-color}b. Four maser sub-clusters have been identified by G81: \"Double Knot\", \"Middle High Velocity Cluster\", \"Northern High Velocity Cluster\", and \"Southern High Velocity Cluster\", all of which seem to be stable for at least 20 years.  LLD identified a \"cocoon\" in the Double Knot, which is more clearly seen by superposing three maps of G81, LLD, and the present work around the coordinate (45, 25) in unit of mas in Figure \\ref {fig:W51M-color}b. Assuming a distance to W51M of 6 kpc, this cocoon has an inner and an outer radii of approximately 12 AU and 60 AU, respectively. A rotation of the cocoon has been also identified by LLD, but it was not found in the present work, probably because of too small a number of maser features detected around the cocoon.  On the other hand, the remaing maser sub-clusters have large position offsets up to 20 mas (120 AU at a distance of 6 kpc) between the previous and the present maps. Likely these are not true motions of sub-clusters but the \"Christmas tree\" effect due to appearance and disappearance of maser features during 20 years. Unlike the bow shock pattern seen in W~51N, no clear feature alignment was found except for the cocoon mentioned above.  Figure \\ref{fig:3D-W51M} presents the detailed 3-D motions of water masers around the four maser sub-clusters. Most of all maser features with measured proper motions are red-shifted with respect to the systemic velocity of this region (50--60\\kms, e.g., G81; \\cite {cox87,rud90, zha95, ho96, zha98, oku01}). Looking at the whole Doppler-velocity range (158\\kms\\  $\\leq\\; $\\vlsr $<$ $-$58\\kms), the kinematics of water masers is apparently random. On the other hand, LLD measured 26 proper motions of water masers with a one-month time baseline and found a stream in the SW direction from the cocoon in a limited range of the Doppler velocity (54\\kms\\  $\\leq\\; $\\vlsr $<$ 68\\kms). We note that the Southern High Velocity Cluster seems to independently have an expanding flow. Maser features on the north--east side of this sub-cluster (around the coordinate [40, $-$100] in unit of mas in Figure \\ref{fig:3D-W51M}) have proper motions toward the cocoon with velocities of 40--80\\kms, probably which are not a flow contracting toward the cocoon but an expanding flow from the point around the coordinate (20, $-$140) in unit of mas in Figure \\ref {fig:3D-W51M}. The systemic Doppler velocity of the candidate originating this expanding flow is expected to be around \\vlsr $\\sim$90\\kms\\ on the basis of the mean Doppler velocity of maser features in the sub-cluster.  \\subsection{W51 South (e4)} \\label{sec:W51S}   Figure \\ref{fig:3D-W51S} presents the 3-D motions of water maser features in W51 South, close to the 3.6-cm continuum source W51 e4 \\citep{gau93} and 2-mm continuum source W~51 e8 \\citep{zha98}. The distribution of water masers in this region seems to have been roughly stable and aligned in the north--west to south--east direction (G81). We found that the water masers exhibit marginally a Doppler-velocity gradient along the maser alignment and a systematic separation motion between the red-shifted and the blue-shifted masers with bipolarity. We attempted a model fitting for the water masers using only a model such as that applied as Step 1 for W51 North water masers, in which we estimated only the originating point of the outflow (see Sect.\\ \\ref{sec:model-fit} and Paper {\\rm I}). Table \\ref{tab:w51smodel-fit} gives its solution.    We have obtained data on the 3-D kinematics of three clusters of water masers in W~51A with VLBI monitoring observations composed of five epochs with an 8-month time baseline. The main conclusions of this paper are as follows.  1. In W51N, a bipolar outflow is clearly exhibited by the maser kinematics, modelled using a radially expanding flow model. The position of the driving source of the outflow coincides with that of the SiO maser source within 0\\arcsec.1. Expansion velocity is consistent with that obtained by EGHMM on the basis of the change in maser distribution for 15 years but decreasing with distance from the driving source from 90\\kms\\ to 50\\kms.  2. In W51M, random motions predominate the maser kinematics. Multiple outflows are considered. Unlike W51N, no systematic change in the maser distribution has been seen during 20 years due to the \"Christmas tree effect\" for maser appearance and disappearance.  3. In W51S, a bipolar outflow was found marginally. The driving source of the outflow has a large offset ($>$ 1\\arcsec) from both of the continuum sources W51 e4 and e8.  4. The maser kinematics in W~51A are heavily biased due to possible concentration of young massive stars and their embedding clouds that destroy symmetry or systematic motion of the outflows. Duration of active massive-star formation in each of the maser clusters is shorter than 10$^{5}$ yr in W51A, which is suggested by the fact that the \\h2o masers associated with the different driving sources of outflows are simultaneously observed.  5. The distance to W51A was estimated to be 6.1$\\pm$1.3 kpc on the basis of the model fitting method applied to the maser kinematics of W51N. On the estimated distance, the kinematical model for the W51N outflow is consistent with that previously proposed on the basis of VLA observations (EGHMM) and exhibits simple deceleration in the flow. This distance value is smaller than that previously adopted ($\\sim$7 kpc) but consistent with the kinematical distance ($\\sim$5.5 kpc).  6. Measurements of relative 3-D bulk motions between the maser clusters have been attempted and showed an apparent separation motion between W~51 North and Main/South. To elucidate the true motions, random motions and biases in the maser kinematics should be carefully taken into account. Reliable kinematical models are indispensable to find accurately systemic bulk motions of maser clusters.  \\bigskip We gratefully acknowledge all staff members and students who have helped in array operation and in data correlation of the J-Net. We also thank Drs. T.~Sasao, Y.~Asaki, and M.~Miyoshi for providing the previous J-Net data. We are also grateful to Drs. M.~J.~Reid, L.~J.~Greenhill, T.~Liljestr\\\"om, and S.~Inutsuka for valuable comments and Dr. J.~A.~Eisner for providing information on SiO masers in W51 North. H. I. was financially supported by the Research Fellowship of the Japan Society of the Promotion of Science for Young Scientist.  \\clearpage"
    },
    "0207/astro-ph0207404_arXiv.txt": {
        "abstract": "High resolution spectroscopic observations of Seyfert galaxies with \\chandra\\ and \\xmm\\ allows us to study the detailed ionization and thermal structures of the X-ray absorbing/emitting material in the circumnuclear environment.  The vast improvement in the spectral resolving power by more than an order of magnitude enables us, for the first time, to unambiguously distinguish the dominant line emission mechanisms and to measure its dynamical properties as well.  The X-ray band harbors spectral transitions from a wide range of ionization states, including valence-shell transitions in K-shell and L-shell ions from most cosmically abundant elements, as well as inner-shell transitions of iron and other mid-$Z$ elements, which can be probed through absorption measurements.  The X-ray spectrum, therefore, provides simultaneous velocity and column density constraints of highly ionized to only slightly ionized gas harbored in many of these systems.  We summarize recent results that have emerged from observations of \\sytwo\\ galaxies with the grating spectrometers onboard \\chandra\\ and \\xmm.  We give particular emphasis to an empirical physical model that we have developed based on the observed spectra, and how it can be used for comparative studies with \\syone\\ galaxies to test the \\agn\\ unification scenarios. ",
        "introduction": "In the past several years, X-ray observations have played an important role in the development of a ``unified model'' of \\agn, in which the observational properties of the various classes (BL Lac objects, \\syone\\ and 2, etc.) are explained solely in terms of their inclination angle with respect to the observer \\citep{miller83, antonucci85, miller90, antonucci93}.  Such issues are naturally related to the understanding of the structure of matter immediately surrounding the central engine, where a variety of physical processes are expected to take place.  Spectral modeling and observations both suggest that the soft X-ray band should contain a wealth of information about the circumnuclear environment, which harbors regions ranging from relatively cool absorbing and reflecting media to hot and tenuous ionized regions.  X-ray spectra of many \\agn\\ exhibit strong emission lines, especially in \\sytwo\\ galaxies where the central continuum source is blocked by a torus of obscuring material and emission lines of large equivalent width are produced, both in the hard (2 -- 10 keV) and soft X-ray bands (0.3 -- 2 keV).  However, owing to the low spectral resolution spectra available prior to the deployment of \\chandra\\ and \\xmm, physical parameters that may, in principle, be derived from the soft X-ray spectrum were not well-constrained.  \\asca\\ and \\sax\\ observed strong soft X-ray line emission in many \\sytwo\\ galaxies.  The nature of this line emission, however, has remained rather controversial, and models involving both photoionized and collisionally ionized plasmas yielded acceptable fits to the data (see, e.g., \\citealt{ueno94, iwasawa94, netzer97, turner97, griffiths98, sako00a}).  \\chandra\\ and \\xmm\\ spectroscopic observations of Seyfert galaxies have provided us with a better understanding of the physical nature of the circumsource medium.  \\citet{sako00b} have shown using \\chandra\\ \\hetg\\ data of Mrk~3 that the soft X-ray emission line spectrum is consistent with that produced in a warm absorbing medium seen in re-emission, providing further evidence that support the unified picture of \\agn.  \\citet{kinkhabwala02a, brinkman02a} and \\citet{ogle02a} have performed quantitative analyses of the X-ray spectra of the archetypal \\sytwo\\ NGC~1068, and placed tight constraints on the column density and velocity distibution of the circumnuclear medium, as well as strict upper limits on the amount of collisionally ionized gas in this object.  \\begin{figure} \\centerline{\\psfig{figure=fig1.ps,width=8.8cm,angle=0}} \\caption[]{The soft X-ray spectrum of NGC~1068 obtained with the \\asca\\ SIS0 (top), the \\xmm\\ \\rgs1 (middle), and the \\chandra\\ \\hetg\\ (bottom).  The spectrum is dominated by a forest of emission lines, which cannot be resolved with the resolving power capabilities of the \\asca\\ SIS ($E/\\Delta E \\sim 20$). \\label{fig:1068}} \\end{figure}  Although the X-ray spectra of \\sytwo\\ galaxies look very different from those of \\syone, there is much overlap in the observable parameter space that characterize the physical nature of the circumsource medium.  In \\syone, the central continuum source is used as a back-lighter to study the properties of the medium along the line of sight through absorption spectroscopy.  In \\sytwo, on the other hand, where the direct view to the continuum source is blocked by the putative molecular torus, the properties of the absorbing medium can be studied by detailed investigation of light that is reprocessed and scattered into our line of sight.  As mentioned above, absorption spectroscopy provides information pertaining to only a particular line of sight along which the observer happens to be looking.  This line of sight may or may not be representative of the entire circumnuclear region, and, as a consequence, it is not clear whether the derived parameters are representative of the global properties of the \\agn. In principle, detection of emission lines superimposed on the background continuum provide some rough estimates of the covering fraction, since, unlike absorption lines, they are produced in regions that lie outside the line of sight as well, but a detailed investigation of the distribution of material is extremely difficult.  A similar complication exists in interpreting the spectra of \\sytwo.  In these sources, the observed emission lines are produced in the entire circumsource medium, and the resulting spectrum is a sum over all the possible lines of sight from the central continuum source along the ionized medium.  The derived column densities are, therefore, biased towards regions of high covering fraction and column density (for reasons to be discussed in detail below).  However, even given these complications, the information content of high-resolution X-ray spectra is certainly revolutionary.  In the remaining sections, we discuss what we can measure from the spectrum and the assumptions that go into the modeling and interpretation.  We identify and compare the parameter space spanned by spectra of both Seyfert 1s and 2s, and discuss how X-ray observations can be used to test the unified model of \\agn.  Since much of the results on the X-ray spectral analyses of \\syone\\ galaxies are discussed elsewhere in this {\\it Proceedings}, here I will focus mainly on high resolution spectral data of \\sytwo\\ galaxies and its relation to those of \\syone. ",
        "conclusions": ""
    },
    "0207/astro-ph0207318_arXiv.txt": {
        "abstract": "We summarize the main results from MODEST-1, the first workshop on MOdeling DEnse STellar systems.  Our goal is to go beyond traditional population synthesis models, by introducing dynamical interactions between single stars, binaries, and multiple systems.  The challenge is to define and develop a software framework to enable us to combine in one simulation existing computer codes in stellar evolution, stellar dynamics, and stellar hydrodynamics.  With this objective, the workshop brought together experts in these three fields, as well as other interested astrophysicists and computer scientists. We report here our main conclusions, questions and suggestions for further steps toward integrating stellar evolution and stellar (hydro)dynamics. \\vspace{1pc} ",
        "introduction": "\\label{intro}  Population synthesis models have been used successfully in comparisons with observations of the global properties of stars, star clusters, and galaxies. The simplest models are constructed from a weighted sum of individual stellar evolution tracks, while more detailed models incorporate some additional information about binary stellar evolution.  For some stellar environments such a synthesis approach is perfectly adequate, and there the main challenge is to deal with the considerable complexities of binary star evolution.  However, the situation is very different for the class of {\\it dense stellar systems}, defined as environments in which a typical star has a significant chance to interact and possibly collide with another star during its lifetime.  In such an environment stars of different ages can exchange mass, disrupt each other or merge, and their merger products can get involved in similar interactions; binary stars can encounter single stars as well as other binaries, where one or more of the stars may already be a merger product; and so on.  Examples of dense stellar systems are star-forming regions and the dense cores of open and globular clusters, as well as galactic nuclei.  It is clear that the possibilities are almost endless.  While population synthesis based on single-star evolution can easily be exhaustive, and synthesis based on a mixture of single stars and binaries can at least aim to be reasonably complete, there is no way that one can anticipate and tabulate all possible multiple-star interactions in dense stellar systems.  Detailed attempts at population synthesis for such systems by necessity have to be dynamical, taking into account the particular ways that stars encounter one another in a given simulation.  During the last few years, several dynamical population synthesis studies have appeared ({\\it cf.} Portegies Zwart {\\it et al.} 2001, Hurley {\\it et al.} 2001).  In these studies, the dynamics of a dense stellar system is modeled through direct $N$-body integration, while the stellar evolution is modeled through fitting formulae that have been obtained from large numbers of individual stellar evolution tracks.  Binary stellar evolution is modeled through the use of semi-analytic and heuristic recipes (Hurley {\\it et al.} 2002).  Astrophysically these results are novel and exciting, but their reliability is not so easy to assess.  Validation is a core issue here, requiring not only detailed internal checks but also comparison between different codes run by different groups.  This question was discussed at some length last year at IAU Symposium 208 in Tokyo, resulting in the specification of a well defined set of initial cluster and stellar parameters (Heggie 2002).  Given the fact that the necessary codes are rather complex, requiring years of development, so far few groups have been able to confront this new challenge.  This stands in contrast to the first collaborative experiment (Heggie {\\it et al.}  1998), which was confined to stellar dynamics (without stellar evolution), and attracted \"entries\" from about 10 groups.  We hope that our new MODEST initiative will stimulate more groups to engage also in the friendly competition of the second collaborative experiment.  Further improvement to the more comprehensive simulations referred to above will require the use of ``live'' stellar evolution models before too long, in order to deal with the unusual types of new stars that can be formed by mergers in dense stellar systems.  However, the challenges of coupling existing stellar evolution codes and stellar dynamics codes are quite daunting.  The first workshop specifically organized to address these challenges was held during July 17-21, 2002 at the American Museum of Natural History in New York City.  The workshop brought together a group of experts in stellar evolution, stellar dynamics, stellar hydrodynamics and other fields of astrophysics, as well as computer scientists.  Originally, the workshop was announced to a small group of people who were known to work on the interface of dynamics and evolution, under the title ``Integrating Stellar Evolution and Stellar Dynamics''.  We originally expected to see a handful of participants for an informal round-table discussion.  The fact that instead 34 attendants convened is a clear sign of the timeliness of the meeting, and the desirability to form a concerted effort to bridge the gap between the stellar evolution and dynamics communities.  This paper offers a summary of the week-long series of discussions held during the workshop, distilled by the organizers (Piet Hut and Mike Shara) and eight of the participants representing a cross section of expertise available during the meeting.  In addition, we have created a web site\\footnote{\\tt http://www.manybody.org/modest.html} where the name `modest' reflects our renaming of the meeting during the last day to MODEST-1, the first workshop on MOdeling DEnse STellar systems.  We plan to hold biannual follow-up meetings, MODEST-2 in Amsterdam in December 2002, and MODEST-3 in Australia in July 2003. In addition, we have started an email list to facilitate ongoing discussions about technical details of dynamical population synthesis simulations.  Further information can be found on our web site.  As a summary of our workshop, this paper contains the input of all of the participants, which are listed below under the acknowledgments. While many of the authors have contributed to various sections, each section has one or two main authors, as follows.  \\S1 and \\S4 were written by Piet Hut, \\S2 by Michael Shara, \\S3 and \\S6 by Piet Hut and Jun Makino, \\S5 by Onno Pols and Ronald Webbink, \\S7 and \\S8 by James Lombardi, \\S9 by Sverre Aarseth and Ralf Klessen, \\S10 by Steve McMillan and Peter Teuben, and \\S11 by Steve McMillan.  In order to make the discussion concrete we have provided specific code fragments in \\S6 and \\S8 below.  We see this paper as the start of a discussion that will ultimately result in the definition of clear standards for interfaces between stellar dynamics, evolution, and hydrodynamics.  However, the current fragments are for illustration only, and are {\\em not} necessarily intended to become part of any future standard. ",
        "conclusions": "\\label{conclusions}  \\newcommand{\\myskip}{\\medskip} \\newcommand{\\myem}{\\myskip\\noindent\\em}  Dynamical simulations of dense star clusters have reached the point where detailed treatments of many aspects of stellar physics must be included.  A significant fraction of stars in globular clusters and galactic nuclei are expected to experience close encounters or actual physical collisions with other stars at some time during the evolution of their parent system.  At the same time, collisions and the effects of stellar and binary evolution can strongly influence cluster dynamics, and may lead to the formation of objects whose properties provide key insights into a cluster's past.  Population synthesis studies have reached a similar conclusion from the opposite direction: dynamical interactions can be vitally important in determining the observed properties of dense stellar systems.  The dynamics of dense stellar systems is also essential for understanding star cluster formation. While protostars in a dense cluster environment build up, they are likely to interact strongly or even merge, and in general they will compete with each other for gas accretion. This has important consequences for the stellar mass spectrum and for the subsequent dynamical evolution of the cluster.  In the workshop MODEST-1 (for MOdeling DEnse STellar systems) the participants discussed many possible avenues for combining stellar physics with stellar dynamics.  Options considered ranged from simple rules and heuristic recipes, to extensive look-up tables using precomputed data, to full-blown ``live'' simulations of stellar and binary evolution and stellar hydrodynamics embedded in a dynamical code.  The following is a consensus view of the current state of the art and an assessment of feasible future developments in the various subfields represented at the meeting.  {\\myem Dynamics.} Traditionally, treatments of stellar and binary evolution and simple recipes for collisions have been realized as modules attached to existing dynamical integrators.  In part, this is historical---dynamicists have had the most pressing reasons to incorporate these effects into their simulations.  However, it is also a fairly natural way to proceed, as the dynamical portion of a large N-body calculation is usually also the part principally concerned with large-scale structure, scheduling, and the orchestration of ``local'' events, such as binary formation and destruction, stellar interactions, stellar evolution, and so on.  One might imagine constructing a more democratic system in which the dynamics, stellar evolution, and hydrodynamics are handled on an equal footing.  However it seems likely that, for the foreseeable future, the dynamical integrator will continue to provide the framework within which other physical effects are incorporated.  {\\myem Evolution of isolated stars.} For ``canonical'' stars that start their lives on the main sequence with more or less normal compositions and never experience close encounters with other stars, there seems to be no strong reason to perform on-the-fly computations of stellar evolution.  Such calculations will almost certainly be of lower precision and contain less physics than existing published calculations.  Rather, the most practical approach involves the use of look-up tables and fitting formulae based on precomputed tracks, essentially as already implemented in current N-body codes.  {\\myem Evolution of isolated binaries.} Binary evolution is too complex for live binary evolution programs, and is expected to remain so for the foreseeable future.  No such programs currently exist, and even simplified versions would likely be too fragile for standalone use.  The physics can be very sensitive to small perturbations and in many cases is not sufficiently well defined for encapsulation in a program to be possible; the number of binary configurations in which the detailed physics is simply unknown is depressingly large.  For the same reasons, no definitive precomputed binary evolutionary tracks exist.  The parameter space is probably too large for look-ups analogous to those used in stellar evolution to be practical in any case.  We thus expect continued use of recipes and heuristic rules of increasing sophistication, again more or less as implemented in existing N-body codes.  We note that this approach has the added benefit of allowing an investigator to identify and parametrize key binary properties, and to vary and study their effects in a controlled way.  {\\myem Hydrodynamics.} Some integrated treatment of stellar collisions is clearly required.  Many collisions involving main-sequence stars can be adequately handled by rules and recipes currently under development, but it seems inevitable that others will have to be performed on the fly, probably using SPH as the description of fluid dynamics best suited to incorporation into a dynamical integrator. Existing codes do not include such modules; most resort to (over)simplified ``sticky'' criteria for stellar mergers.  Basic self-contained SPH (or shortcuts such as entropy-sorting) treatments of two-body collisions could in principle be added to existing codes in a relatively straightforward way.  Integration of arbitrary stellar encounters within a full N-body environment is probably a feasible, but much longer-term, goal.  {\\myem Collision Products.} Collisions---either direct, between unbound stars, or indirect, resulting from binary evolution or temporary capture of stars in binaries---will give rise to ``non-canonical'' stars quite unlike those normally studied by stellar evolution codes or reported in the literature.  They will be out of thermal equilibrium, will probably be rapidly rotating, and will have unusual composition and entropy profiles.  We will not be able to precompute and interpolate all the possibilities.  Here we really do need live stellar evolution codes to study the appearance and evolution of the collision products.  However, such studies pose a severe challenge to existing techniques, and lie beyond the capabilities of current stellar evolution codes.  The creation of a robust, standalone module to handle the evolution of collision products is a high priority.  \\myskip How to make the pieces communicate?  It is unrealistic to expect researchers to completely rewrite their codes (no matter how attractive such a prospect might be...) in order to merge them with other programs.  Rather, it is better to create modular programs by encapsulating parts or all of existing computer codes and define robust interfaces specifying clearly the functionality of each module and the data that must be provided and returned for each to work. Such an approach is vital, as it will facilitate controlled comparison of competing techniques.  Behind the interface, the structure of each module will be entirely up to the programmer, so long as it conforms rigorously to the agreed-upon interface specifications.  We have begun a study of the interfaces, data structures, and communication protocols needed to realize this goal.  The external representation of simulation data is also an important and unresolved issue---we need to share data between programs in an efficient, extendible, and non-destructive way.  The first MODEST workshop was successful in bringing together the three astrophysics communities of researchers working the fields of stellar evolution, stellar hydrodynamics, and stellar dynamics.  We will continue to hold these workshops twice yearly, thereby providing a meeting point for those who are actively involved in simulating dense stellar systems.  For further details, see the MODEST web site.\\footnote{\\tt http://www.manybody.org/modest.html}   \\bigskip  {\\it Acknowledgments.}  We acknowledge the input of the participants of the MODEST-1 workshop.  Here is the complete list of those who attended part or all of the workshop:  \\begin{table}[h] \\begin{tabular}{ll} Sverre Aarseth     &  Jun Makino       \\\\ David Chernoff     &  Marc Hemsendorf \\\\ Scott Fleming      &  Rosemary Mardling \\\\ Marc Freitag       &  Steve McMillan \\\\ Yoko Funato        &  David Merritt \\\\ John Fregeau       &  John Ouellette \\\\ Mirek Giersz       &  Onno Pols \\\\ Paul Grabowski     &  Dina Prialnik  \\\\ Atakan G\\\"urkan    &  Fred Rasio \\\\ Jarrod Hurley      &  Helmut Schlattl \\\\ Piet Hut           &  Mike Shara \\\\ Vicky Kalogera     &  Shawn Slavin \\\\ Ralf Klessen       &  Rainer Spurzem \\\\ Attay Kovetz       &  Jerry Sussman \\\\ Yuexing Li         &  Peter Teuben \\\\ James Lombardi     &  Dany Vanbeveren \\\\ Mordecai Mac Low   &  Ron Webbink \\\\ \\end{tabular} \\end{table}  In addition, we also acknowledge comments on the manuscript by Douglas Heggie and Pavel Kroupa. R.S.K. acknowledges financial support by the Emmy Noether Program of the Deutsche Forschungsgemeinschaft (DFG, grant KL1358/1). J.C.L. acknowledges support from NSF grant AST 00-71165.  His work was also partly supported by the National Computational Science Alliance under grant AST 98-0014N and utilized the NCSA SGI/Cray Origin2000 parallel supercomputer. S.M. acknowledges support from NASA ATP grant NAG5-10775. R.F.W.'s participation is supported in part by NASA grant NAG 5-11016 to the University of Illinois.  \\def\\araa{{\\em Ann.\\ Rev.\\ Astron.\\ Astrophys.}} \\def\\aas{{\\em Astron.\\ Astrophys.\\ Suppl.\\ Ser.}} \\def\\aj{{\\em Astron.\\ J.}} \\def\\anap{{\\em Ann.\\ Astrophys.}}   % \\def\\apj{{\\em Astrophys.\\ J.}} \\def\\apjs{{\\em Astrophys.\\ J.\\ Suppl.\\ Ser.}} \\def\\aap{{\\em Astron.\\ Astrophys.}} \\def\\jcam{{\\em J.\\ Comput.\\ Appl.\\ Math.}} \\def\\jcp{{\\em J.\\ Comput.\\ Phys.}} \\def\\jfm{{\\em J.\\ Fluid Mech.}} \\def\\mnras{{\\em Mon.\\ Not.\\ R.\\ Astron.\\ Soc.}} \\def\\nat{{\\em Nature}} \\def\\pta{{\\em Phil.\\ Trans.\\ A.}} \\def\\ptp{{\\em Prog.\\ Theo.\\ Phys.}} \\def\\prd{{\\em Phys.\\ Rev.\\ D}} \\def\\prl{{\\em Phys.\\ Rev.\\ Lett.}} \\def\\prsa{{\\em Proc.\\ R.\\ Soc.\\ London A}} \\def\\pasp{{\\em Pub.\\ Astron.\\ Soc.\\ Pac.}} \\def\\zp{{\\em Z.\\ Phys.}} \\def\\za{{\\em Z.\\ Astrophys.}}"
    },
    "0207/astro-ph0207632_arXiv.txt": {
        "abstract": "Using the standard prescription for the rates of supernovae type~II and type~Ia, we compare the predictions of a simple model of star formation in galaxies with the observed radial gradients of abundance ratios in a sample of early-type galaxies to infer the relative contribution of each type of supernova. The data suggests a correlation between the fractional contribution of Type~Ia to the chemical enrichment of the stellar populations ($1-\\xi$) and central velocity dispersion of order $1-\\xi\\sim -0.16\\log\\sigma_0+0.40$, so that the type~Ia contribution in stars ranges from a negligible amount in massive ($\\sigma_0\\sim 300$ km s$^{-1}$) galaxies up to $10\\%$ in low-mass ($\\sim 100$ km s$^{-1}$) elliptical galaxies. Our model is parametrized by a star formation timescale ($t_{\\rm SF}$) which controls the duration of the starburst. A correlation with galaxy radius as a power law ($t_{\\rm SF}\\propto r^\\beta$) translates into a radial gradient of the abundance ratios. The data implies a wide range of formation scenarios for a simple model that fixes the luminosity profile, ranging from  inside-out ($\\beta=2$), to  outside-in formation ($\\beta =-1$), as is consistent with  numerical simulations of elliptical galaxy formation. An alternative scenario that links $t_{\\rm SF}$ to the dynamical timescale favours  inside-out formation over a smaller range $0.4<\\beta < 0.6$. In both cases, massive galaxies are predicted to have undergone a more extended period of star formation in the outer regions with respect to their low-mass counterparts. ",
        "introduction": "Type~Ia supernovae (SNIa) describe stellar explosions whose spectra show lines of elements or intermediate mass such as silicon, and of the iron group, but no hydrogen lines. They currently occupy a very special place in cosmology since they can be used as standard candles. Even though their luminosities span an order of magnitude, empirical correlations between their absolute magnitude and the shape of their light curves can be used in order to determine distances (Phillips 1993) in an analogous way to classical Cepheids. Furthermore, the absolute luminosity of SNIa is orders of magnitude brighter than variable stars, enabling us to observe them at high redshift. This is a tool currently being used to constrain the cosmological parameters (e.g. Perlmutter et al. 1999; Riess et al. 1998). However, SNIa also play a very important role in the process of galactic chemical enrichment since they contribute a large amount of iron to the interstellar medium (Thielemann, Nomoto \\& Yokoi 1986). In fact Type~Ia supernovae might even be the main producers of iron in the Universe (Ishimaru \\& Arimoto 1997; however see Gibson, Loewenstein \\& Mushotzky 1997). The observed properties of SNIa's suggest a binary model in which at least one of the stars is a white dwarf that  reaches the Chandrasekhar limit ($\\sim 1.4M_\\odot$) by accretion or by merging with another white dwarf. The timescale for the explosion is thereby limited by the lifetimes of stars which end up as white dwarfs, i.e. with masses $M\\simlt 8M_\\odot$. This implies  that Type~Ia's can be observed long after star formation subsided. Indeed, all supernovae observed in early-type galaxies --- which feature no ongoing star formation --- are Type~Ia, whereas late-type galaxies display a mixture of Type~Ia and Type~II supernovae (Cappellaro et al. 1997)  On the other hand, Type~II supernovae (SNII) show hydrogen lines in their spectra and arise from the core collapse of a single, massive ($M\\simgt 8M_\\odot$) star at the end of its lifetime, which occurs between $1-50$~Myrs after the core hydrogen burning phase started. Hence, the timescales of either type of supernovae are remarkably different. Furthermore, the yields of chemical elements are also in sharp contrast, since SNIa produce much more iron than SNIIs, so that stars born during the first phases of star formation --- when the contribution from SNIa to the interstellar medium was negligible --- display an enhancement of $\\alpha$ elements over iron, with respect to the younger generations of stars, such as the Sun, which are born in an environment polluted by both types of supernovae.  Observations of radial gradients in the colours (Franx, Illingworth \\& Heckman 1989; Peletier et al. 1990; J\\o rgensen, Franx \\& Kj\\ae gaard 1995) and line indices (Gonz\\'alez 1993; Davies, Sadler \\& Peletier 1993; Peletier et al. 1999) in elliptical galaxies show gradients, being mostly redder and more metal rich at their centres, although some early-type galaxies display blue cores (Menanteau, Abraham \\& Ellis 2001). Broadband photometry has been a technique repeatedly used to infer the star formation history of galaxies notwithstanding the age-metallicity degeneracy (Worthey 1994), which prevents us from getting a well-defined picture of galaxy formation using colours alone. A combined analysis of line indices seem to break the degeneracy since their age and metallicity dependence can be rather different (Kuntschner 2000; Trager et al. 2000a, 2000b). Abundance ratios represent an alternative observable since the timescales for SNIa and SNII are remarkably different. Giant early-type galaxies feature an overabundance of $\\alpha$ elements over iron (Peletier 1989; Worthey, Faber \\& Gonz\\'alez 1992; Trager et al. 2000a; Kuntschner 2000),  indicative of a short duration of the star formation stage, so that mostly SNII contribute to the metallicity of the stellar component.  This has been used in several models of galaxy formation in order to constrain the star formation history of galaxies. Matteucci (1994) analysed the observed [Mg/Fe] ratios in ellipticals to constrain the star formation history to timescales shorter than $0.1$~Gyr. Furthermore, the trend of [Mg/Fe] with galaxy mass was used to imply either an increasing star formation efficiency with galaxy mass, or a top-heavy initial mass function, so that more massive stars were formed in ellipticals compared to a more quiescent environment such as in a disk galaxy. A detailed analysis of abundance ratios in metal-poor stars is a valuable tool for  determining the star formation history in the solar neighbourhood (Gratton et al. 2000), and this can be extended to stellar populations in globular clusters. In a recent analysis of abundance ratios in a sample of six red giant stars in the $\\omega$~Cen cluster, Pancino et al. (2002) found evidence for the contribution of SNIa to the composition of younger, more metal rich red giants.  In this paper we explore the radial gradients observed in abundance ratios such as [Mg/Fe] in elliptical galaxies as a function of the star formation timescale. The relative contribution from SNIa and SNII can be compared with  observations to infer the formation process of the stellar components, enabling us to connect the star formation history and the dynamical history of early-type galaxies. In \\S2 and \\S3 we describe the model used to predict the rates of either type of supernovae and their contribution to chemical enrichment. \\S4 describes our model predictions and compares them to observed data. Finally, in \\S5 we discuss the implications of the comparison between our simple model and the observed data. ",
        "conclusions": "Our model links the abundance ratios to the star formation timescale, which cannot be directly observed. Instead, the data available give radial gradients. A simple power law, as described in equation (\\ref{eq:beta}), can be used in order to throw light on the possible correlation between the infall rate or the star formation efficiency and the dynamical properties of the galaxy. Figure~5 compares the observed radial gradients with the predicted ones, as a function of [Mg/Fe]. The hollow squares represent the data from Trager et al. (2000a), measured at two projected radial positions: $R_e/8$ and $R_e/2$, which enables us to determine a gradient. A clear trend is seen towards positive gradients in more [Mg/Fe]-enhanced (i.e. more massive) galaxies. The other data points come from the compilation of Kobayashi \\& Arimoto (1999) from which we selected the work of Gonz\\'alez (1993, GON); Davies et al. (1993, DSP), and Carollo et al. (1993, CDB) estimated by the authors of the compilation to have the most reliable data. There is a very large scatter in this compilation, with a wide range of slopes, both negative and {\\sl positive}. In agreement with the data from Trager et al. (2000a), a correlation can be seen so that galaxies with positive slopes tend to have the largest super-solar abundances, which correspond to the most massive galaxies as seen in figure~2. The solid line gives a least-squares fit to the data from Trager et al. (2000a). The dashed line also shows a similar trend in the compilation of Kobayashi \\& Arimoto (1999). Estimating abundance ratios from line indices is still a very delicate issue especially because of the difficulty in finding a suitable set of stars for calibration. Hence, a theoretical approach is called for, using the computations of Tripicco \\& Bell (1995), who recomputed all of the Lick/IDS spectral indices from a grid of theoretical spectra and atmospheres with varying abundance ratios. The response functions to nonsolar ratios is used to correct the standard population synthesis models of Worthey (1994), calibrated for solar abundance ratios. The uncertainties in such translation frmo abundances to line indices is discussed in Trager et al. (2000a), and an error bar including these uncertainties is shown in figure~5.  \\begin{figure} \\epsfxsize=3.5in \\begin{center} \\leavevmode \\epsffile{cmnsn1a_f5.eps} \\bigskip \\caption{ Observed abundance gradients estimated from the measurements of Trager et al. (2000a; TRG; hollow squares) at $R_e/8$ and $R_e/2$. A typical error bar is shown for illustration. The solid line is a linear squares fit to this data. A rough estimate of [Mg/Fe] gradients from Kobayashi \\& Arimoto (1999) is also shown. The data points from this compilation correspond to the sample of Gonz\\' alez (1993, GON); Davies et al. (1993, DSP); and Carollo et al. (1993, CDB). The trend clearly seen in the data from Trager et al. (2000a) as well as hinted at by the void in the lower right corner of the data from Kobayashi \\& Arimoto (1999), and illustrated by the dashed line, can be interpreted as a hint towards a correlation between galaxy mass (which scales with [Mg/Fe]) and the lower limit of the gradient.} \\end{center} \\end{figure}  A negative gradient implies an inside-out formation process, which is suggestive of a monolithic collapse scenario, with gas falling towards the centres of dark matter halos, with a gas density profile that decreases outwards and a temperature that decreases inwards, thereby generating a process of star formation that starts at the centre and spreads outwards. On the other hand, recent TreeSPH numerical simulations of galaxy formation (Sommer-Larsen, Gotz \\& Portinari 2002) find early-type galaxies (both ellipticals and lenticulars) in which star formation proceeds in the opposite way, namely outside-in, mainly triggered by merging. Furthermore, recent observations of field spheroidals in the Hubble Deep Field (Menanteau et al. 2001) show blue cores, result which is currently interpreted as secondary bursts of star formation, but which could be an indication of an outside-in formation process, which implies a positive gradient.  In the light of these data, we infer a varying correlation between the star formation timescale ($t_{\\rm SF}$) and galaxy radius. A simple approach, fixing $\\gamma$ as in equation (\\ref{eq:bfg}) gives a slope for this correlation in the range  $-1.3 < \\beta < 2$. A better approximation, relating $\\beta$ and $\\gamma$ through the dynamical timescale, assuming a fixed mass-to-light ratio (\\ref{eq:ml}) gives $0.4 < \\beta < 0.6$ with a mass or luminosity profile slope of $1.8 < \\gamma < 2.2$. The higher values of $\\gamma$ correspond to positive gradients in the abundance ratio.  The scatter is rather large, and it can only imply a very weak correlation between $\\beta$ or $\\gamma$ and some dynamical parameter such as central velocity dispersion or mass. This scatter gives another hint of merging as the major mechanism in the assembly process of early-type systems. In a purely monolithic collapse scenario, one would expect a well-defined correlation between the dynamical timescale and the star formation timescale, in such a way that most of the early star formation would occur at the centre. The void in the lower right corner of figure~5 (illustrated by the thick dashed line) shows that we can infer a correlation between [Mg/Fe] (or galaxy mass as shown in Figure~2) and a lower limit to the gradient. This can be interpreted as monolithic collapse being possible only in low mass systems. Higher mass galaxies must be assembled through merging. We emphasize here that the inside-out versus outside-in classification may be an oversimplification. Depending on the simple model we choose --- i.e. fixed luminosity profile, as in (\\ref{eq:bfg}), or correlating $\\beta$ with the dynamical timescale, as in (\\ref{eq:ml}), massive galaxies give respectively real outside-in star formation ($\\beta <0$) or a shallower slope for the correlation between star formation timescale and radius. In either case, the correlation suggested by the dashed line in figure~5 hints at more star formation in the  outskirts of massive galaxies compared to low mass systems during the major stage of star formation.  It is worth mentioning that the slope of the abundance ratios is a more robust estimator than the absolute abundance. The latter is strongly dependent on the amount of gas ejected in outflows, which is an important mechanism in early-type galaxies as hinted at by the mass-metallicity relation (Arimoto \\& Yoshii 1987; Ferreras \\& Silk 2001), whereas the slope would depend on the scaling of the outflows with radial distance, which should be similar to the scaling of the star formation efficiency and infall timescale, thereby reducing the effect of outflows on the slope.  The radial gradients in abundance ratios represent an alternative way of studying the universality of the IMF, compared to analyses of the metallicities of local stars or the intracluster medium (Wyse 1997). Figure~3 shows that a non-universal IMF would translate into a slope change of the abundance ratio [Mg/Fe]. For instance, if we expect a top-heavy IMF in environments with a high star formation rate (i.e. with a short $t_{\\rm SF}$) then the correlation between [Mg/Fe] and star formation timescale should be steeper. However, this calls for a more detailed model which is beyond the scope of this paper. We have also neglected the effect of early galactic winds which could eject $\\alpha$-enhanced material out of the galaxy. A correlation of these winds with the local escape velocity could explain the observed radial gradient of Mg abundance in ellipticals (Martinelli, Matteucci \\& Colafrancesco 1998).  The model presented here relies on the fact that we understand the mechanisms that trigger SNIa and we are able to predict their rates or at least their ratio with respect to SNII. Short bursts of star formation imply that most of the contamination from SNIa will go to the ISM and not so much to the stellar component. Hence, we expect gas in ellipticals to have solar abundance ratios. There has been a long controversy over this point. The study of {\\sl ASCA} observations of a few clusters hinted at a significant SNIa contamination of the gas in rich clusters (Ishimaru \\& Arimoto 1997). However, Gibson et al. (1997) showed that a different calibration could invalidate this hypothesis. Arimoto et al. (1997) even raised the doubt of whether abundance estimates using the standard iron L-line complex in X-rays is giving wrong metallicities. The latest analysis of six early-type galaxies in Virgo using {\\sl ROSAT} and {\\sl ASCA} observations (Finoguenov \\& Jones 2000) seem to agree with a high iron content, i.e. an important contribution from SNIa to the enrichment of the ISM in elliptical galaxies.  In this paper, we conclude that a simple treatment of the abundance ratios allows us to infer the star formation history and its connection to the dynamical formation history. The data seems to rule out a monolithic or secular formation scenario for the stellar component of massive ellipticals. On the other hand, the large scatter in the slope of the radial dependence of [Mg/Fe] in low-mass galaxies implies both a merging and a monolithic ``mechanism'' should be invoked for these systems. A combined analysis of abundance ratios both in the stellar populations and in the interstellar medium will enable us to explore the connection between these two histories, and to quantify the importance of merging events {\\sl both} in the star formation and dynamical histories of early-type galaxies."
    },
    "0207/astro-ph0207068_arXiv.txt": {
        "abstract": "Using linear kinetic plasma theory the relation between electron density and magnetic field fluctuations for low-frequency plasma waves for Maxwellian background distribution functions of arbitrary temperatures in an uniform magnetic field is derived. By taking the non-relativistic temperature limit this ratio is calculated for the diffuse intercloud medium in our Galaxy. The diffuse intercloud medium is the dominant phase of the interstellar medium with respect to radio wave propagation, dispersion and rotation measure studies. The differences between the relation of electron density and magnetic field fluctuations from the linear kinetic theory as compared to the classical MHD theory are established and discussed. ",
        "introduction": "Important input quantities for the quasilinear test-particle description of cosmic ray transport in weakly-turbulent astrophysical plasmas are the wavenumber power spectra of magnetic field fluctuations. Within the plasma wave viewpoint, the plasma irregularities are usually modelled as a superposition of linear waves well below the ion cyclotron frequency, such as Alfven and magnetosonic waves. However, the observed turbulence properties in the more distant interstellar and intergalactic plasmas are obtained from radio propagation measurements as dispersion measures, rotation measures and interstellar scintillation. These turbulence diagnostics are biased towards the high-density ionized interstellar phases with large volume filling factors, i.e. the diffuse intercloud gas and HII envelopes. In particular, dispersion measure and scintillation data are primarily diagnostics of density and only secondarily of magnetic field. These diagnostics demonstrate the existence of interstellar density irregularities with Kolmogorov-type ($ \\propto \\omega ^{-s}$, $s=const.$) frequency power spectra extending over 11 decades in frequency, much below ($\\omega <<\\Omega _p$) the proton gyrofrequency (Rickett \\cite{r90}, Armstrong et al. \\cite{ars95}).  Often the electromagnetic fluctuations are described within magnetohydodynamic (MHD) theory (e.g. Sturrock \\cite{s94}, Goldreich \\& Sridhar \\cite{gs95}, Hollweg \\cite{h99}, Lithwick \\& Goldreich \\cite{lg02}) which is appropriate at large turbulence scalelength $l\\ge l_{MHD}$. However, the plasma parameter of the diffuse intercloud medium $g=\\nu _{ee}/\\omega _{p,e}\\simeq 10^{-10}$ is much smaller than unity, so that a kinetic description of the electromagnetic turbulence seems to be necessary. It is the purpose of the present paper to provide the relation between electron density and magnetic field fluctuations on the basis of the linear kinetic plasma theory. In paper II of this series we will use the kinetic turbulence relations to calculate frequency power spectra of electron density fluctuations from anisotropic power spectra of magnetic field fluctuations in the form of Alfven and magnetosonic waves. Such anisotropic interstellar magnetic field power spectra are required in order to be in accord with the heating/cooling balance of the diffuse intercloud medium (Lerche \\& Schlickeiser \\cite{ls1}). ",
        "conclusions": "Using linear kinetic plasma theory we have calculated the relation between electron density and magnetic field fluctuations for low-frequency plasma waves for Maxwellian background distribution functions of arbitrary temperatures in an uniform magnetic field. By taking the non-relativistic temperature limit we determined this ratio for the diffuse intercloud medium in our Galaxy. The diffuse intercloud medium is the dominant phase of the interstellar medium with respect to radio wave propagation, dispersion and rotation measure studies. We have found differences between the relation of electron density and magnetic field fluctuations from the linear kinetic theory as compared to the classical MHD theory.  Whereas shear Alfven waves are incompressive in MHD theory, linear kinetic theory yields the non-zero relation (\\ref{nealfv3}) even in the limit of vanishing electron kineticity $\\rho _e=0$. Only at very small wavenumbers $k<<(2/3)|\\Omega _p|\\sin \\theta /V_A=(2/3)(\\omega _{p,i}/c)\\sin \\theta $ the kinetic result agrees with the MHD result.  For magnetosonic waves the kinetic ratio of the normalised density and magnetic field fluctuations is modified from the MHD ratio by the factor $f_{kin}=(3/2)(1+\\beta \\sin ^2\\theta )$ which is independent of wavenumber and varies within values of $1.5$ and $1.5(1+\\beta )$.  In the next paper of this series we will use these kinetic turbulence relations to calculate frequency power spectra of electron density fluctuations from anisotropic power spectra of magnetic field fluctuations in the form of Alfven and magnetosonic waves. Such anisotropic interstellar magnetic field power spectra are required in order to be in accord with the heating/cooling balance of the diffuse intercloud medium.  \\smallskip \\smallskip {\\it Acknowledgements} We gratefully acknowledge support by the Deutsche For\\-schungs\\-ge\\-meinschaft through Sonderforschungsbereich 191."
    },
    "0207/astro-ph0207542_arXiv.txt": {
        "abstract": "The most massive evolved stars (above 50\\,M$_{\\sun}$) undergo a phase of extreme mass loss in which their evolution is reversed from a redward to a blueward motion in the HRD. In this phase the stars are known as Luminous Blue Variables (LBVs) and they are located in the HRD close to the Humphreys-Davidson limit. It is far from understood what causes the strong mass loss or what triggers the so-called giant eruptions, active events in which in a short time a large amount of mass is ejected. Here I will present results from a larger project devoted to better understand LBVs through studying the LBV nebulae. These nebulae are formed as a consequence of the strong mass loss. The analysis concentrates on the morphology and kinematics of these nebulae. Of special concern was the frequently observed bipolar nature of the LBV nebulae. Bipolarity seems to be a general feature and strongly constrains models of the LBV phase and especially of the formation of the nebulae. In addition we found outflows from LBV nebulae, the first evidence for ongoing instabilities in the nebulae. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207297_arXiv.txt": {
        "abstract": "Stars account for only about 0.5\\% of the content of the Universe; the bulk of the Universe is optically dark.  The dark side of the Universe is comprised of:  at least $0.1\\%$ light neutrinos; $3.5\\% \\pm 1\\%$ baryons; $29\\% \\pm 4\\%$ cold dark matter; and $66\\% \\pm 6\\%$ dark energy.  Now that we have characterized the dark side of the Universe, the challenge is to understand it. The critical questions are:  (1) What form do the dark baryons take?  (2) What is (are) the constituent(s) of the cold dark matter? (3) What is the nature of the mysterious dark energy that is causing the Universe to speed up. ",
        "introduction": "The past five years have witnessed great progress in identifying the basic features of our Universe.  It is spatially flat and thus has the critical density ($\\rho_{\\rm crit} = 3H_0^2/8\\pi G \\approx 10^{-29}\\,{\\rm g\\,cm^{-3}}$).  The expansion is speeding up, not slowing down (i.e., $q_0 < 0$).  The mass/energy density is distributed as follows:  \\begin{itemize}  \\item Bright stars:  0.5\\% \\item Baryons (total):  $4\\%\\pm 1\\%$ \\item Nonbaryonic dark matter:  $29\\% \\pm 4\\%$ \\item Neutrinos:  at least 0.1\\% and possibly as large as 5\\% \\item Dark Energy:  $66\\% \\pm 6\\%$  \\end{itemize}  The evidence for flatness comes from measurements of the anisotropy of the cosmic microwave background (CMB) on angular scales of about 1 degree.  The position of the first acoustic peak at multipole number 200, as determined by the BOOMERanG, MAXIMA, DASI and CBI experiments, implies that $\\Omega_0 = 1\\pm 0.04$ (Sievers et al, 2002). This means that the curvature radius of the Universe is greater than about 5 times the Hubble radius since $R_{\\rm curv} = H_0^{-1}/|\\Omega_0 -1|^{1/2}$.  In the next sections I will discuss the evidence for the accounting of the various dark components.  I did not mention the cosmic microwave background, which today accounts for $\\Omega_{\\rm CMB} =  2.47h^{-2}\\times 10^{-5}$ (about 0.005\\%), or the relativistic neutrino backgrounds which account for $\\Omega_\\nu = 0.56h^{-2}\\times 10^{-5}$ per relativistic neutrino species.  Though unimportant to the energy budget today, at very early times relativistic particles were dominant.  The existence of (at least) three components that evolve differently with redshift divides the evolution of the Universe into (at least) three epochs:  1.  Early ($z> 10^4$ and $t< 10^4\\,$yrs) radiation-dominated era (photons, neutrinos, and a spectrum of relativistic particles that grows with temperature); 2. Matter-dominated era ($10^4 > z > 0.2$, $10\\,{\\rm Gyr} > t > 10^4\\,$yrs) during which cosmic structure grew; and 3. Dark-energy dominated era ($z< 0.2$, $t> 10\\,$Gyr) characterized by accelerated expansion and cessation of structure formation.  I note that the precision of the present accounting still allows for an unidentified component that contributes perhaps as much as 10\\% of the critical density.  Now that we can enumerate the components of the Universe, the task is to understand them.  Three critical questions arise:  \\begin{enumerate}  \\item What form do the dark baryons take?  \\item What is (are) the constituent(s) of the nonbaryonic dark matter?  \\item What is the nature of the dark energy?  \\end{enumerate} ",
        "conclusions": ""
    },
    "0207/astro-ph0207283_arXiv.txt": {
        "abstract": "{\\em XEUS}, the {\\em X-ray Evolving Universe Spectroscopy} mission, constitutes at present an ESA-ISAS initiative for the study of the evolution of the hot Universe in the post-Chandra/XMM-Newton era. The key science objectives of XEUS can be formulated as the:  \\noindent -- Search for the origin, and subsequent study of growth, of the first massive black holes in the early Universe.  \\noindent -- Assessment of the formation of the first gravitationally bound dark matter dominated systems, i.e. small groups of galaxies, and their evolution.  \\noindent -- Study of the evolution of metal synthesis up till the present epoch. Characterization of the true intergalactic medium.  To reach these ambitious science goals the two salient characteristics of the XEUS observatory entail:  1. Its effective spectroscopic grasp, combining a sensitive area $> 20$ m$^{2}$ below a photon energy of 2 keV with a spectral resolution better than 2 eV. This allows significant detection of the most prominent X-ray emission lines (e.g. O-VII, Si-XIII and Fe-XXV) in cosmologically distant sources against the sky background.  2. Its angular resolving power, between 2 and 5 arc seconds, to minimize source confusion as well as noise due to the galactic X-ray foreground emission.  To accommodate these instrument requirements a mission concept has been developed featuring an X-ray telescope of 50 meter focal length, comprising two laser-locked spacecraft, i.e. separate mirror and detector spacecraft's. The telescope is injected in a low earth orbit with an inclination commensurate with the ISS, a so-called fellow traveler orbit. At present an on-orbit growth of the mirror spacecraft is foreseen through a robotic upgrade with the aid of the ISS, raising the mirror diameter from 4.5 to 10 meter. The detector spacecraft, formation flying in a non-Keplerian orbit in tandem with the mirror spacecraft will be replaced at 5 year intervals after run-out of consumables with an associated upgrade of the focal plane package. ",
        "introduction": "At the end of the 20th century, the promise of high spatial and spectral resolution in X-rays has become a reality. The two major X-ray observatories nowadays operational, NASA's Chandra and ESA's XMM Newton, are providing a new, clear-focused, vision of the X-ray Universe, in a way that had not been possible in the first 40 years of X-ray astronomy. These two missions complement each other very well: Chandra has a $\\simless 0.5$ arcsec angular resolution and the capability of high spectral resolution on a variety of point sources with the High- and Low-Energy Transmission Grating Spectrometers, while XMM-Newton has a larger spectroscopic area and bandwidth (up to $\\sim$15 keV), and the capability of high-resolution spectroscopic observations on spatially extended sources.  Despite their superb capabilities, these two missions cannot be used for detailed studies of objects at very high redshifts ($z \\simmore 5$). At present, the Chandra and XMM-Newton deep surveys in very narrow pieces of the sky allow us to detect quasars up to redshifts of 6.28, (\\cite{jbleeker-D:bra02}), but we are not able to produce X-ray spectra of these objects at such distance.  Both Chandra and XMM-Newton are expected to be operational for the next ten years. This is the typical timescale for a new mission to be planned and developed, therefore this is the time to assess what is the future of X-ray astronomy, and to start planning for Chandra's and XMM-Newton's follow-up. ESA's response to this challenge has been cosmology, and the unique role that X-ray astronomy can play in studying the formation and evolution of the hot Universe. ",
        "conclusions": ""
    },
    "0207/astro-ph0207556_arXiv.txt": {
        "abstract": "We report here the successful commissioning of the  \\\\ PN.Spectrograph, the first special-purpose instrument for the \\\\ measurement of galaxy kinematics through the PN population. ",
        "introduction": "Planetary nebulae: to most of the participants at this conference, beautiful \\\\ objects displaying a range of structures, spherical, bipolar and rhomboidal,\\\\ riotous colours, and the complex spectral signature of ionised gases, shocks, and dust. To others, the extragalactic crowd, they are seen as featureless points of light with the simplest of all possible spectra, a solitary green emission line. ",
        "conclusions": ""
    },
    "0207/astro-ph0207410_arXiv.txt": {
        "abstract": "A numerical scheme that incorporates a thermal leakage injection model into a combined gas dynamics and cosmic ray (CR, hereafter) diffusion-convection code has been developed. The hydro/CR code can follow in a very cost-effective way the evolution of CR modified planar quasi-parallel shocks by adopting subzone shock-tracking and multi-level adaptive mesh refinement techniques. An additional conservative quantity, $S= P_g/\\rho^{\\gamma_g-1}$, is introduced to follow the adiabatic compression accurately in the precursor region, especially in front of strong, highly modified shocks. The ``thermal leakage'' injection model is based on the nonlinear interactions of the suprathermal particles with self-generated MHD waves in quasi-parallel shocks. The particle injection is followed numerically by filtering the diffusive flux of suprathermal particles across the shock to the upstream region according to a velocity-dependent transparency function that controls the fraction of leaking particles. This function is determined by a single parameter, $\\epsilon$, which should depend on the strength of postshock wave turbulence, but is modeled as a constant parameter in our simulations. We have studied CR injection and acceleration efficiencies during the evolution of CR modified planar shocks for a wide range of initial shock Mach numbers, $M_0$, assuming a Bohm-like diffusion coefficient. For expected values of $\\epsilon$ the injection process is very efficient when the subshock is strong, leading to fast and significant modification of the shock structure. As the CR pressure increases, the subshock weakens and the injection rate decreases accordingly, so that the subshock does not disappear. Although some fraction of the particles injected early in the evolution continue to be accelerated to ever higher energies, the postshock CR pressure reaches an approximate time-asymptotic value due to a balance between fresh injection/acceleration and advection/diffusion of the CR particles away from the shock. In the strong shock limit of $M_0 \\gsim 30$, the injection and acceleration processes are largely independent of the initial shock Mach number for a given $\\epsilon$, while they are sensitively dependent on $M_0$ for $M_0<30$. We conclude that the injection rates in strong parallel shocks are sufficient to lead to rapid nonlinear modifications to the shock structures and that self-consistent injection and time-dependent simulations are crucial to understanding the non-linear evolution of CR modified shocks. ",
        "introduction": "Two decades ago it was recognized that CRs are probably produced very efficiently via diffusive shock acceleration (DSA) in ubiquitous astrophysical shocks, \\citep[for early reviews see, \\eg][]{dru83,blaeic87,berzkry88}. After the initial successes of the simple concept that the particles can gain energy while temporarily trapped in the converging flows around a shock, it was quickly realized that the full DSA treatment requires one to consider the complex nonlinear interactions between energetic particles, resonantly scattering waves and the underlying plasma \\citep{madru01}. One of the important aspects of those interactions is the injection of suprathermal particles into the CR population at shocks. According to quasi-linear theory as well as plasma simulations of strong quasi-parallel shocks, the streaming motion of superthermal particles against the background plasma can induce wave generation leading to strong downstream MHD waves that scatter particles and inhibit the particles from leaking upstream \\citep[\\eg][]{Bell78,Quest88}. Particles in or close to the thermal population are especially restricted in this way. As a consequence only a small fraction of suprathermal particles can swim upstream against the wave-particle interactions in the plasma flow and be injected into the higher energy CR population to be further accelerated via the Fermi process. The injection process and its efficiency control the amplitude of the CR population and hence the degree of shock modification.  Recently, significant progress in understanding this injection process in parallel shocks has been made through self-consistent, analytic, nonlinear calculations by \\citet{malvol95} and \\citet{mal98}. The resulting theory has only one parameter; namely, the intensity of the downstream waves, and that is tightly restricted, both by the theory and by comparison with hybrid plasma simulations. By adopting Malkov's analytic solution, we have developed a numerical treatment of this injection model and incorporated it into a combined gas dynamics and CR diffusion-convection code \\citep{gies00}. According to the \\citet{gies00} simulations, the injection process seemed to be self-regulated in such a way that the injection rate reaches and stays at a nearly stable value after quick initial adjustment, but well before the CR shock reaches a steady state structure. \\citet{gies00} found about $10^{-3}$ of incoming thermal particles to be injected into the CRs, roughly independent of Mach numbers. However, due to severe computational requirements associated with the need to resolve structures down to the physical shock thickness, those simulations were carried out only until the characteristic maximum momentum of $(p_{\\rm max}/m_{\\rm p} c)\\sim 1$ was achieved. Since strong shocks were still evolving at the end of the simulations, the time-asymptotic limit could not be estimated for either the CR acceleration efficiency or the CR spectrum.  Unlike ordinary gas shocks, the CR shock is a collisionless structure, and includes a wide range of length scales associated not only with the dissipation into ``thermal plasma'', but also with the nonthermal particle diffusion process. Those are characterized by the so-called diffusion lengths, $D_{\\rm diff}(p) = \\kappa(p)/u$, where $\\kappa(p)$ is the spatial diffusion coefficient for CRs of momentum $p$, and $u$ is the characteristic flow velocity \\citep{kanjon91}. For strong scattering of suprathermal particles $D_{\\rm diff}$ may not greatly exceed the physical dissipative, or ``gas'' shock thickness . Accurate solutions to the CR diffusion-convection equation require a computational grid spacing significantly smaller than $D_{\\rm diff}$, typically, $\\Delta x \\sim 0.1 D_{\\rm diff}(p)$. On the other hand, for a realistic diffusion transport model with a steeply momentum-dependent diffusion coefficient, the highest energy, relativistic particles have diffusion lengths many orders of magnitude greater than those of the lowest energy particles.  To follow the acceleration of highly relativistic CRs from suprathermal energies, all those scales need to be resolved numerically. However, the diffusion and acceleration of the low energy particles are important only close to the shock owing to their small diffusion lengths. Elsewhere, they are effectively advected along with the underlying gas flow. At higher energies the needed resolution is less severe, and on scales larger than $\\sim D_{\\rm diff}(p_{\\rm max})$ the bulk flow and the nonthermal particles decouple, so resolution requirements are controlled by whatever factors are necessary to define the broader flow properties.  Thus it is necessary to resolve numerically the diffusion length of the particles only around the shock. So, in \\citet{kang01} we first implemented a {\\it shock tracking scheme} \\citep{levshy95} to locate the shock position exactly and then refine the grid resolution only around the shock by applying multi-levels of refined grids \\citep{berglev98}. The main properties of this code are: 1) the shock is tracked as an exact discontinuity, 2) a small region around the shock is refined with multi-level grids, 3) it is very cost-effective in terms of computational memory and time.  In the present contribution, we have studied the CR injection and acceleration during the evolution of modified planar shocks by implementing the numerical method for the thermal leakage injection model of \\citet{gies00} into an enhanced version of the CR/AMR hydrodynamics code, which we name CRASH (Cosmic Ray Amr SHock) code. With our new CRASH code we were able to calculate the CR injection and acceleration efficiencies with a Bohm-like diffusion coefficient for higher particle energies ($p/m_{\\rm p}c \\gg 1$) in shocks over a wide range of initial Mach numbers. In the following section the basic equations solved in the simulations are presented. We describe the thermal leakage injection model and its numerical implementation in our CR transport code in \\S 3. We then outline our numerical methods in the CR/AMR hydrodynamics code in \\S 4. In \\S 5 we present and discuss our simulation results, followed by a summary in \\S 6. Finally we discuss some test calculations in the Appendix. ",
        "conclusions": "In order to study cosmic-ray modified shocks we have developed a new numerical code, CRASH (Cosmic-Ray Amr SHock), by implementing a thermal leakage injection scheme introduced by \\citet{gies00} into a new hydro/CR code with Adaptive Mesh Refinement and Shock Tracking scheme by \\citet{kang01}. Our CR injection model is based on the ``thermal leakage'' process at quasi-parallel CR shocks. Injection is regulated by the convolution of the population in the high energy tail of the Maxwellian velocity distribution of the postshock gas and a transparency function, $\\tau_{\\rm esc}(p,u_d,\\epsilon)$, determined by the strength of downstream MHD waves, expressed by a parameter, $\\epsilon$. The injection rate in our model is then controlled largely by the subshock Mach number, $M_s$, since that determines the ratio of the postshock flow speed to the breadth of the thermal distribution. For weak shocks injection becomes more difficult as $M_s$ decreases, but is independent of $M_s$ for strong shocks, when the subshock compression asymptotes. With our CRASH code the CR injection and acceleration at astrophysical shocks can be simulated numerically even for strongly momentum dependent spatial diffusion coefficients.  Using these tools the time evolution of CR shocks has been followed with a Bohm type diffusion model. We started from the initial conditions for pure gasdynamic shocks of various initial Mach numbers ($M_0$) without any pre-existing CRs. In such simulations the CR injection rate is not treated as a fixed free parameter, as it has been done traditionally. Instead, once an ``intelligent estimate'' is made of the strength of the trapping wave field of the downstream MHD turbulence, the injection is followed naturally and self-consistently during nonlinear evolution of the flows. For strong shocks with initial strengths, $M_0\\gsim30$, a substantial portion of the particles in the tail of the Maxwellian distribution have velocities high enough to leak upstream against the wave-particle interactions, so the injection is efficient and fast. As the CR pressure increases at the subshock and the in-flowing plasma is compressed, however, the subshock slows down with respect to the far-upstream flow. This modification occurs rather promptly and before the postshock CR pressure reaches an approximate time-asymptotic value. Afterwards the evolution of the shock structure becomes secular. For strong shocks, the subshock speed in our simulations decreases by about 15-17 \\%, and the postshock CR pressure absorbs up to 60 \\% of the ram pressure of the initial shock, which corresponds to 90 \\% of the ram pressure of the upstream flow in the evolved subshock rest frame at the end of our simulations. Once the postshock CR pressure becomes constant, the shock structure evolves approximately in a ``self-similar'' way, because the scale length of shock broadening increases linearly with time.  The injection rate, defined as the fraction of the particles passed through the shock that are accelerated to form the CR population, becomes as high as $\\xi \\sim 0.01$ early in the evolution of strong shocks, independent of $M_0$. As the shock is modified, however, the subshock Mach number decreases down to $M_s\\sim 2.9 M_0^{0.13}$, and the injection rate reduces towards the limiting value corresponding to weak subshocks. Our approximate numerical estimate for the limiting value is $\\xi \\sim 10^{-3.4}$ for models with $\\epsilon=0.2$. Since the CR spectrum continues to extend to higher momenta, and since those highest energy particles diffuse rapidly away from the shock when $\\kappa \\propto p$, energy begins to be transported away from the shock transition. This allows the subshock to persist and the total compression to become large compared to steady energy-conserving gasdynamic shocks. Also the shock structure evolves towards the high compression ratios seen in kinetic equation spherical CR shocks and steady state planar shocks with energy escape through spatial or momentum boundaries.  Finally, the main conclusions of our time-dependent simulations of CR modified shocks based on Bohm-like diffusion and a physically-based thermal leakage CR injection are:  1. In the strong shock limit of initial Mach number $M_0 \\gsim 30$, significant physical processes such as injection and acceleration become largely independent of the initial shock Mach number, and only weakly depend on the postshock wave amplitude parameter, $\\epsilon$, for values considered here ($0.2\\le \\epsilon \\le0.3$). According to our thermal leakage model, the overall injection rate approaches $\\xi \\sim 10^{-3}$ and the fraction of upstream flow kinetic energy {\\it in the initial shock rest frame} that has been transferred to CRs is $\\Phi \\sim 0.6$  for strong shocks. The ratio of the postshock CR pressure to the {\\it instantaneous} ram pressure of the subshock with respect to the upstream plasma approaches $\\sim 90$\\%. On the other hand, the injection rate and acceleration efficiency are sensitively dependent on $M_0$ for low Mach number shocks ($M_0\\lsim 30$).  2. In our simulations, with no initial CRs around the shock, the thermal leakage injection is very efficient initially when the subshock is strong, but it becomes much less efficient as the subshock weakens due to nonlinear feedback from the CR pressure. For example, the time-averaged injection rate can be fitted as a power-law form, $\\xi \\sim 2.5\\times 10^{-3} ( \\tilde t)^{-0.4}$ for $1<\\tilde t< 200$ ($\\tilde t = t/t_0$) for strong shocks of $M_0\\gsim 30$ and the inverse wave-amplitude parameter $\\epsilon =0.2$.  3. Although some particles injected early in the shock evolution continue to be accelerated to ever higher energies, the immediate postshock CR pressure reaches a quasi-steady value when a balance between injection/acceleration and advection/diffusion is achieved. The region of quasi-steady postshock properties spreads for Bohm diffusion in an almost self-similar fashion because both diffusive and advection rates scale linearly with time. This expansion maintains the large precursor compression at values that can be large compared to what one would derive from Rankine-Hugoniot relations for steady gas shocks with a relativistic equation of state.  4. The sum of the compression through the precursor and across the subshock calculated near the terminal time, $\\tilde t=100$, for models with $\\epsilon=0.2$ can be approximated by $\\rho_2 / \\rho_0 \\sim 1.5 M_0^{0.6}$ for $M_0 < 80$, but this $\\rho_2/\\rho_0 - M_0$ relation flattens for $M_0\\ge 80$. Considering that $\\rho_2$ continues to increase until the terminal time, especially for higher Mach number shocks, and that the Mach number of the subshock with respect to the far upstream ($M_0^\\prime$) decreases for strong shocks with high initial Mach numbers ($M_0$), this is probably consistent with  the $\\rho_2/\\rho_0 \\propto M_0^{\\prime 3/4}$ scaling relation found for steady shocks and for spherical shocks by several authors \\citep{berz95, mal97, berzell99}. The large density compression is possible in the planar shocks without energy escape from the whole computed system, since the shock structure  spreads out far upstream and far downstream, due to strong momentum dependent diffusivity.  Although our simulations have been carried out only until the highest momentum is $p_{\\rm max} \\sim 40 m_{\\rm p}c$ due to severe computational requirements, these simulation results provide useful guidance for long-term evolution of CR modified shocks. In a future study, we will extend the integration time so that $p_{\\rm max}/m_{\\rm p}c\\gg1$ can be achieved and the free escape at the upper momentum boundary is applied in order to compare the time-dependent simulations with previous studies of steady-state shocks \\citep{elleich85,mal97,madru01}.    The algorithms applied in CRASH are already mostly well-documented and well-tested. The hydrodynamic scheme, for example, is a member of the well-known Godunov family and was verified by \\citet{levshy95}. Our Crank-Nicholson scheme for solution of the diffusion-convection equation was tested against previous, independent implementations and analytic test particle diffusive shock acceleration solutions in \\citet{kanjon91}, where we employed the PPM hydrodynamics algorithm. Other tests of nonlinear CR shock solutions using several previous hydrodynamical implementations were given for both diffusion-convection \\citep[\\eg][]{kjr92,kanjon95} and two-fluid \\citep[\\eg][]{jk90,kanjon95,kanjon97} models of CR transport. \\citet{kanjon95} using PPM hydrodynamics and \\citet{kanjon97} using TVD magnetohydrodynamics methods, for example, provided comparisons between the CR transport approach applied here with our earlier ``thermal leakage'' injection model and Monte Carlo, hybrid plasma simulations, and {\\it in situ} measurements of heliospheric shocks. While the newer injection scheme utilizes a significantly improved physical model to determine which particles are injected at a shock, both the new and the old underlying numerical schemes have much in common, as discussed in \\citet{gies00}. More recently \\cite{kang01} provided convergence tests of the combined shock tracking, AMR, convection diffusion code upon which CRASH is based, again using our previous thermal leakage injection model.  In order to extend the above test base to CRASH we present here several comparisons of solutions to results from these earlier works. In particular we have tested CRASH against piston-driven shock problems that were calculated with our PPM/CR code \\citep{jk90, kjr92}. First, using a two-fluid version of our CRASH code with the modified entropy equation, we calculated a shock driven by a constant speed, plane piston with $u_p=0.3$ moving into a medium of $\\rho=1$, $P_g=7 \\times 10^{-4}$, and $P_c= 3.5\\times 10^{-4}$. Closure parameters are assumed to be $\\gamma_c=5/3$ and $<\\kappa>=0.01$. There is no injection of new CRs in this case, only reacceleration of the previous CR population. Thus, it is a test for nonlinear modified shocks of the dynamical coupling terms included in equations (\\ref{mocon})-(\\ref{econ}). Fig. 13 shows the evolution of the shock at $t=3,~6,$ and 9, which can be compared with Fig. 1 of \\citet{jk90}. Agreement is excellent. The dashed horizontal line shows the steady state solution for $P_c$, which is fully consistent with the CRASH simulation. This test calculation demonstrates that the two fluid version of the CRASH code can accurately follow the evolution of a strongly CR modified shock. The dynamical coupling between the CRs and the gas is the same in the two-fluid model as in the diffusion-convection, kinetic equation model, except for the multi-scale character introduced by a momentum dependent diffusion coefficient into the kinetic equation model.  To extend the dynamical evaluation to include momentum dependent influences, we calculated another piston driven shock using the diffusion-convection version of CRASH. In this case the piston has velocity $u_p=1$ moving into a medium of $\\rho=1$, $P_g=P_c=0.001$. Two simulations were done, one with and one without the modified entropy equation (\\ref{scon}). The injection process was turned off and only the base-grid was used (\\ie $l_{\\rm g,max}=0$). The diffusion coefficient was given the form $\\kappa=p^{0.5}$. Fig. 14 shows the evolution of the shock at $t=10,~30,~50,~100$ and 150, which was calculated with the modified entropy equation. The last of these compares directly with the dotted curves in Fig. 3 of \\citet{kjr92}. Some details of the early evolution differ, which cause the exact locations and peaks of the transient density spike to differ somewhat, but the overall structures and the shock transitions at $t = 150$ are in very good agreement, including the compression, as well as the gas and CR pressures. The noted small early evolution differences among different codes have been seen before. They reflect the sensitivity to numerical details of the very rapid changes taking place in the shocks at the start of such simulations.  Finally, we demonstrate the benefits of including the modified entropy equation in treatments of strongly modified CR shocks. Fig. 15 shows from two simulations the subshocks and precursor structure for the piston-driven shock shown in Fig. 14 at $t=20$. The two sets of curves were calculated with the modified entropy condition enforced (S code: solid line) and without it (E code: dashed line). The lower left panel shows the specific entropy, $\\log{P_g/\\rho^{\\gamma}}$. To the right of the subshock, where the flow is adiabatic, this quantity should be a constant. However, because the flow there is highly supersonic, the total energy is dominated by the kinetic energy. Consequently, errors in computing the gas pressures used in equations (\\ref{mocon}) and (\\ref{econ}), as extracted from a combination of the total energy and momentum quantities lead to significant errors in the entropy of the preshock gas. That in turn, leads to errors in the gas subshock and CR acceleration behaviors. Enforcing entropy conservation in smooth flows eliminates the problem, as illustrated in the figure."
    },
    "0207/hep-ph0207125_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.1in}} We propose that cold dark matter is made of Kaluza-Klein particles and explore avenues for its detection.  The lightest Kaluza-Klein state is an excellent dark matter candidate if standard model particles propagate in extra dimensions and Kaluza-Klein parity is conserved. We consider Kaluza-Klein gauge bosons.  In sharp contrast to the case of supersymmetric dark matter, these annihilate to hard positrons, neutrinos and photons with unsuppressed rates.  Direct detection signals are also promising.  These conclusions are generic to bosonic dark matter candidates. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207376_arXiv.txt": {
        "abstract": "We used ultraviolet spectra from HST/STIS ($R=30,000$), together with optical spectra from Keck/HIRES ($R=45,000$), to study the three MgII-selected absorption systems at $z=0.9254$, $0.9276$, and $0.9342$ toward the quasar PG~$1206+459$. A multi-phase gaseous structure, with low-ionization components produced in small condensations and high-ionization ones in diffuse clouds, is indicated in all three systems. Each system is likely to represent a different galaxy with absorption due to some combination of interstellar medium, coronal gas, halo gas, and high-velocity clouds.  Even with the improved sensitivity of HST/COS, we will only be able to obtain high-resolution ultraviolet spectra of the brightest quasars in the sky.  A larger telescope with ultraviolet coverage will enable quasar absorption line studies of hundreds of galaxies, including a wide range of galaxy types and environments at low and intermediate redshifts. ",
        "introduction": "The lightbeams from distant quasars pass through various galaxies and probe physical properties of gaseous structures in the galaxies. Therefore, studying absorption features in the background quasar spectrum is a unique method for tracing the cosmic evolution of the universe.  Strong MgII absorbers are almost always associated with luminous galaxies ($> 0.05L^*$).  Thus, using this tool we can study the predecessors of the giant spiral and elliptical galaxies that we see in the nearby universe.  Three MgII systems (A, B, and C), at redshifts $z=0.9254$, $0.9276$, and $0.9342$, are found along the line of sight toward the quasar PG~$1206+459$. In a previous study, Churchill \\& Charlton (1999) found the three systems to be multi-phase absorbers with MgII clouds embedded in extended, high-ionization gas that gives rise to CIV, NV, and OVI. Their analysis was based upon the combination of low-resolution data from HST/FOS and high-resolution data from Keck/HIRES.  However, many important issues remained unresolved, such as (1) whether CIV, NV and OVI arise in the same layer of gas and whether their profiles are smooth or have sub-structure; (2) whether the high-ionization phase ``envelops'' the low-ionization phase or whether it is offset in velocity; (3) whether the majority of SiIV in the $z=0.9276$ system arises in a single MgII cloud that is similar to a Milky Way high-velocity cloud.  In May 2001, we obtained a stunning high-resolution ($R=30,000$) HST/STIS spectrum, covering Lya and high-ionization transitions SiIV, CIV, and NV, for the three systems. Through photoionization modeling of the various chemical transitions in this spectrum and in the earlier Keck/HIRES spectrum, the metallicities, abundance patterns, and ionization states of absorbing gas clouds have been constrained (for details on the modeling technique, see Ding et al. 2002). The results are presented in the following sections along with the physical interpretations of individual phases of gas. ",
        "conclusions": ""
    },
    "0207/astro-ph0207140_arXiv.txt": {
        "abstract": "The method of constrained randomisation, which was originally developed in the field of time series analysis for testing for nonlinearities, is extended to the case of three-dimensional point distributions as they are typical in the analysis of the large scale structure of galaxy distributions in the universe.\\\\ With this technique it is possible to generate for a given data set so-called surrogate data sets which have the same linear properties as the original data whereas higher order or nonlinear correlations are not preserved. The analysis of the original and surrogate data sets with measures, which are sensitive to nonlinearities, yields valuable information about the existence of nonlinear correlations in the data. On the other hand one can test whether given statistical measures are able to account for higher order or nonlinear correlations by applying them to original and surrogate data sets.\\\\ We demonstrate how to generate surrogate data sets from a given point distribution, which have the same linear properties (power spectrum) as well as the same density amplitude distribution but different morphological features.\\\\ We propose weighted scaling indices, which measure the local scaling properties of a point set, as a nonlinear statistical measure to quantify local morphological elements in large scale structure. Using surrogates is is shown that the data sets with the same 2-point correlation functions have slightly different void probability functions and especially a different set of weighted scaling indices.\\\\ Thus a refined analysis of the large scale structure becomes possible by calculating local scaling properties whereby the method of constrained randomisation yields a vital tool for testing the performance of statistical measures in terms of sensitivity to different topological features and discriminative power.\\\\ Keywords: cosmology: theory - large-scale structure of Universe - methods: numerical ",
        "introduction": "One of the important issues in cosmology today is characterising the nature of the large scale structure in the spatial distribution of galaxies as revealed by observations. Statistical measures provide important tools for the quantitative characterisation of the morphology of the galaxy distribution and for the comparison of the various cosmological models with observations. Among the first and still most frequently used measures are the 2-point correlation function (e.g. Peebles 1980 and references therein; Norberg et al. 2001) and the power spectrum (e.g. Szalay et al. 2001; Tegmark et al. 2001; Schuecker et al. 2001) which have the advantage of being directly related to simulations for different cosmological models. However, they are linear measures which cannot provide any information about higher order or nonlinear correlations in the data set. Nowadays the large surveys like the SDSS (York et al 2000) or 2dF (Colless et al. 2001) yield excellent observations from galaxy distributions consisting of up to one million galaxies with which it becomes possible to identify higher order correlations. Therefore it is necessary to develop statistical descriptors, which go beyond the 2-point correlation function.\\\\ Many measures which go beyond the 2-point correlation function have already been studied in detail. The correlation analysis of the data sets has very early been extended to higher order correlation functions (e.g. 3-point correlation function (Groth \\& Peebles 1977),  4-point correlation function (Fry \\& Peebles 1978), up to 8-point correlation function (Meiskin,  Szapudi and \\& Szalay 1992)) and are now applied to the newest available data sets (Szapudi et al. 2002). Analysis in the Fourier space have involved the calculation of eigenvectors of the sample correlation matrix (e.g. Vogeley et al. 1996) and of the bispectrum (Mataresse, Verde \\& Heavens 1997; Verde et al. 1998; Scoccimarro et al. 2001). More recently, also the correlations between Fourier phases have been quantified by calculating entropies (Chiang \\& Coles 2000, Chiang 2001), which measure the amount of non-gaussian signatures in the spatial patterns of a density field. Other measures have been developed in order to characterise the topology of the large scale structure. Among the first measures of this kind introduced in cosmology has been the void probability function (e.g. White 1979; Ghigna et al. 1994), which can be expressed by a sum over all n-point correlation functions. Another well-known measure is the genus curve of the density contrast (Weinberg, Gott \\& Mellott 1987), which has only recently been applied to the 2dF galaxy redshift survey data set (Hoyle, Vogeley \\& Gott 2002). Both the void probability function and the genus curve can be regarded as special cases of the Minkowsky functionals which also have extensively been used in the analysis of the galaxy distributions (e.g. Mecke et al. 1994; Kerscher et al. 1997; Bharadwaj et al. 2000). The concepts derived in the field of non-linear dynamics have been applied to large scale structure analysis by calculation e.g. the multifractal dimension spectrum (e.g. Borgani 1995 and references therein; Pan \\& Coles 2000). One common feature of all these measures is that they analyse the data set as a whole and therefore focus on the {\\it global} aspects of matter distribution.\\\\ In the field of image analysis various statistical methods for the morphological and textural description of given structures have been developed, too (for an overview see e.g. Tuceryan \\& Jain 1993 and references therein). It has been shown that in the context of (human) texture analysis it is crucial to consider both global and {\\it local} aspects of given structures in order to perform an effective structure characterisation (Sagi \\& Julesz 1985; Jain \\& Farrokhnia 1991) leading e.g. to texture detection and discrimination. Furthermore it has been pointed out (Julesz 1981, 1991) that nonlinear and local data processing steps play a crucial role in the detection and discrimination of textural features. It has been shown that nonlinear local filters (so-called scaling indices) which measure the local scaling properties of point sets are well suited to accomplish feature and texture detections tasks in image processing (R\\\"ath \\& Morfill 1997; Jamitzky et al. 2001). The general approach for estimating these measures, which is closely related to the formalism of the multifractal dimension spectrum, makes them ideal candidates for describing the local structural features in galaxy distributions, too. In this paper we propose a modified version of the scaling index formalism ('weighted scaling indices') as a local nonlinear statistical measure for analysing the large scale structure in the universe.\\\\ For the assessment of the different statistical measures it is of vital interest to have detailed knowledge about the performance of the different measures in terms of sensitivity to certain morphological features or in terms of discrimination power. In the analysis of nonlinear time series (Theiler et al. 1992; Schreiber \\& Schmitz 1996; Schreiber \\& Schmitz 1997; Schreiber 1998) the technique of constrained randomisation, that allows a test for weak nonlinearities in time series, has been developed. Applying this method to a given data set one obtains an ensemble of randomised versions of the original data set (so-called surrogate data), in which some previously defined statistical constraints are maintained while all other properties are subject to randomisation. Using a different reasoning, one can also use this method in order to test whether given statistical measures are able to account for higher order or nonlinear correlations or special morphological features in the data applying the measures to be tested to both the original data and the surrogates and comparing their discriminative power. In this work we extend known techniques for generating surrogates to the case of three-dimensional point distributions as they are typical in the analysis of the large scale structure. We calculate several linear and nonlinear measures for the data and surrogates and evaluate them in terms of sensitivity and discriminative power.\\\\ The outline of the paper is as follows: In the next Section the properties of the simulated data set are briefly described. In Section 3 we introduce the statistical measures we used in our study. Whilst the well-known measures used for references are only briefly reviewed, the phase entropy and the concept of weighted scaling indices are described in more detail. In Section 4 the results of our calculations are shown. Section 5 contains the main conclusions and gives an outlook for future work. ",
        "conclusions": "We adapted and used the method of surrogate data to analyse three-dimensional point distributions. It could be shown that with the help of the method of iteratively refined surrogates it is possible to generate data sets which have the same power spectrum and amplitude distribution in configuration space but differ significantly with respect to their topological structure. The existence of these topological differences points to nonlinear processes in the early evolution of the universe and is likely to be important cosmologically. Hence nonlinear measures need to be developed to quantify them - after which the consequences for the different models have to be discussed. Amongst the standard measures the void probability function gave relatively small differences between the original and the surrogate data sets, while the 2-point correlation function and power spectrum were the same (by construction).\\\\ These results show that linear global measures like the 2-point correlation function and power spectrum are only of limited usefulness for the characterisation of the morphological content of given point distribution and that their discriminative power is, therefore, also limited. This is mainly due to the fact that these second order statistical measures are 'blind' to the distributions of Fourier phases, which are responsible for the fine details of cosmic structures. We further analysed the distribution of the phases by calculating the phase entropy and found that the surrogates cannot be told apart from the original data set using this measures. Therefore, a more sophisticated analysis of the obviously inherent correlation in the distribution of the phases is required.\\\\ We showed that the development of nonlinear morphological descriptors, which are based on the analysis of the local scaling behaviour of the mass distribution, can offer new possibilities to refine our statistical methods so that previously ignored subtle but important features can be both detected and quantitatively characterised. Using such a measure (weighted scaling indices) a clear distinction based on the different topological features between surrogates and the original data set is possible. In the context of evaluating different statistical measures used in the analysis of large scale structure the method of constrained randomisation represents a vital tool with which the quality of the newly developed measures can be tested systematically. Thus a better quantitative characterisation of the spatial patterns in the galaxy distribution becomes possible, improving the interpretation and our outstanding of the large scale structure in the universe."
    },
    "0207/astro-ph0207189_arXiv.txt": {
        "abstract": "The Sloan Digital Sky Survey (SDSS) is making a multi-colour, three dimensional map of the nearby Universe. The survey is in two parts. The first part is imaging one quarter of the sky in five colours from the near ultraviolet to the near infrared. In this imaging survey we expect to detect around 50 million galaxies to a magnitude limit $g \\sim 23$. The second part of the survey, taking place concurrently with the imaging, is obtaining spectra for up to 1 million galaxies and 100,000 quasars. From these spectra we obtain redshifts and hence distances, in order to map out the three-dimensional distribution of galaxies and quasars in the Universe. These observations will be used to constrain models of cosmology and of galaxy formation and evolution.  This article describes the goals and methods used by the SDSS, the current status of the survey, and highlights some exciting discoveries made from data obtained in the first two years of survey operations. ",
        "introduction": "What is the Universe made of? How did the Universe begin? How will it end?  These are some of the fundamental questions which can be addressed by studying the large scale distribution of galaxies in the Universe. It is widely believed that the galaxies we see today formed at the sites of tiny ($\\sim$ 1 part in $10^5$) density fluctuations in the early Universe. The form of these density fluctuations (which, if Gaussian in nature, may be fully described by their power spectrum) are predicted by cosmological models, and depend on such parameters as the mean matter density $\\Omega_m$, the fraction of baryonic matter $\\Omega_b/\\Omega_m$ and any contribution to the cosmological density from vacuum energy, also known as the cosmological constant $\\Omega_\\Lambda$. On large scales, the clustering of galaxies can be predicted from the primordial density fluctuations using linear perturbation theory. By measuring this large-scale clustering, we can thus obtain important constraints on cosmological models. By studying the intrinsic properties of the galaxies themselves, such as luminosity, colour and morphology, we can test theories for how galaxies are born and evolve.  In order to measure the clustering of galaxies reliably, it is important to use a systematic and well-defined catalogue. Systematic surveys date back to that of Messier, published in three parts in the 1770s and 1780s, although it was not realized at the time that some of the nebulae catalogued by Messier were other galaxies outside our own Milky Way. It was only in 1923, by a careful measurement of the distance to the Andromeda Nebula (M31 in Messier's catalogue), that Edwin Hubble proved definitively that M31 was a large galaxy separate from our own. Hubble later discovered the expansion of the Universe, and found that the recession velocity of a galaxy is in direct proportion to its distance from us, the Hubble law. Since then, a number of galaxy surveys have been published, starting with that of Shapley and Ames in 1932 \\cite{sa1932}, and including more recently the APM Galaxy Survey \\cite{msel90}, which contains positions and magnitudes for about three million galaxies. Most of these surveys are based on photographic plates, and there is concern that such surveys could be missing a substantial fraction of low surface-brightness galaxies, eg.~\\cite{disney76}. There is also apprehension that uncertainties in the photometric calibration of these surveys could lead to spurious measurement of galaxy clustering on large scales \\cite{fhs92,mes96}.  Smaller surveys have been made using charge coupled device (CCD) detectors. These solid state devices, unlike photographic plates, have a linear response to light and $\\sim 50$ times higher quantum efficiency, but the limited size of these detectors has before now precluded the construction of wide-area galaxy surveys.  In order to map out the three-dimensional distribution of galaxies, as opposed to just their two-dimensional projection on the celestial sphere, one needs the distance to each galaxy. This may be obtained by measuring the spectrum of light emitted by a galaxy. The Doppler shift in features towards the red end of the spectrum (the redshift) may be used to infer a galaxy's recession velocity and hence its distance from the Hubble relation. Until recently, galaxy spectra were painstakingly measured one-by-one, and it is only in the last few years that optical fibre multiplexing has been used to measure redshifts for many thousands of galaxies.  The largest redshift survey to date is the nearly-completed Two Degree Field (2dF) Galaxy Redshift Survey carried out on the Anglo-Australian Telescope \\cite{colless2001}. While containing redshifts for more than 200,000 galaxies, the 2dF survey is based on the photographic APM Galaxy Survey, with the potential problems mentioned above.  The Sloan Digital Sky Survey (SDSS) collaboration was therefore formed in 1988 with the aim of constructing a definitive map of the local universe, incorporating CCD imaging in several passbands over a large area of sky, and measurement of redshifts for around one million galaxies. In order to complete such an ambitious project over a reasonable timescale, it was decided to build a dedicated 2.5-metre telescope equipped with a large CCD array imaging camera and multi-fibre spectrographs. The survey itself began in April 2000, and observations are scheduled to finish in June 2005. In this article I review some important aspects of the survey, including an overview of survey operations (\\S\\ref{sec:overview}), a description of the preliminary public data release (\\S\\ref{sec:edr}), and a selection of some early science results (\\S\\ref{sec:science}). ",
        "conclusions": "The Sloan Digital Sky Survey is now fully operational and is producing high quality data at a prodigious rate. We have imaged \\area\\ deg$^2$ of sky in five colours and have obtained more than \\nspec\\ spectra. Much exciting science has already come out of just a small fraction of the final dataset and we look forward to many more exciting discoveries in the coming years.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is {\\tt http://www.sdss.org/}.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Princeton University, the United States Naval Observatory, and the University of Washington.  It is a pleasure to thank SDSS colleagues for supplying some of the figures. I would particularly like to thank Donald York for his careful reading of the manuscript.  \\clearpage"
    },
    "0207/astro-ph0207230_arXiv.txt": {
        "abstract": "We examine the reliability of the merger trees generated for the Monte-Carlo modeling of galaxy formation. In particular we focus on the cold gas fraction predicted from merger trees with different assumptions on the progenitor distribution function, the timestep, and the mass resolution.  We show that the cold gas fraction is sensitive to the accuracy of the merger trees at small-mass scales of progenitors at high redshifts. One can reproduce the Press--Schechter prediction to a reasonable degree by adopting a fairly large number of redshift bins, $N_{\\rm step}\\sim 1000$, in generating merger trees, which is a factor of ten larger than the canonical value used in previous literature. ",
        "introduction": "Understanding the formation and evolution of galaxies is a fundamental step in linking the initial condition of the universe and the cosmological observational data.  Recent systematic studies of high-redshift objects, such as quasars and Lyman-break galaxies, should provide important clues to the early universe, although their proper interpretation is often not so straightforward, mainly because those objects certainly do evolve in time.  A theoretical study of galaxy evolution, especially its spectroscopic evolution, from a cosmological context, was begun by \\citet{tinsley80} and followed by many authors (e.g., \\cite{Bruzual83}; \\cite{AY86}; \\cite{GR87}; \\cite{CB91}; \\cite{BC93}; \\cite{KA97}). These studies are based on a so-called `one-zone' model which assumes that a galaxy does not interact with other galaxies. It is now fairly established, however, that structures in the universe have built up hierarchically from small to large scales as in a cold dark matter (CDM) model. This means that a galaxy interacts and sometimes merges with other galaxies even if it was an isolated system at birth. The predictions in the one-zone model therefore may be significantly different from what happened to galaxies in a hierarchical universe.   White and Frenk (1991) developed a detailed analytic formalism to describe the formation and evolution of galaxies while taking account of the hierarchical merging of dark-matter halos, gas cooling, star formation, and supernova feedback.  Subsequent numerical approaches in modeling hierarchical merging of dark halos employ two somewhat different algorithms; one is called the `block model' in which a random-Gaussian density fluctuation field is generated by dividing a hypothetical rectangular box recursively (\\cite{CK88}; \\cite{Cole91}; \\cite{Cole94}). While this algorithm is simple and straightforward, the resulting halo masses are necessarily binned in discrete steps of a factor of two. The other generates a realization of halo merger trees according to a probability distribution function predicted by the extended Press--Schechter theory (\\cite{Bower91}; \\cite{Bond91}; \\cite{KW93}; \\cite{SK99}; \\cite{SL99}).  The latter is widely used in studying the cosmological evolution of galaxies in a hierarchical universe (\\cite{KWG93}; \\cite{baugh98}; \\cite{SP99}; \\cite{Cole00}; \\cite{nagashima}). Throughout the present paper, we call the latter method the Monte-Carlo modeling of merger histories (simply, the Monte-Carlo modeling), while it is usually referred to as a semi-analytic model of galaxy formation (SAM).  The most important ingredient in Monte-Carlo modeling is the conditional joint-probability distribution function of a set of \\textit{progenitor} halos of mass $M_2^{j}$ at a redshift of $z_2$, which is a part of a \\textit{parent} halo of mass $M_1$ at $z_1$, conceptually written as \\begin{eqnarray} \\label{eq:jointprob} {\\rm Prob}(M_2^1, M_2^2, \\cdots, M_2^N, z_2 | M_1, z_1) dM_2^1 dM_2^2 \\cdots dM_2^N \\cr \\qquad (N=1, \\cdots , \\infty). \\end{eqnarray} Unfortunately only an analytical expression for the conditional one-point probability distribution function, Prob($M_2^i$, $z_2 |M_1$, $z_1$), is known based on the extended Press--Schechter theory (for the special case of the Poisson initial power spectra, see a different approach by \\cite{SL99}); one thus needs to employ an additional \\textit{assumption} in generating realizations of merger trees of halos in general (e.g., \\cite{KW93}; \\cite{SK99}). Furthermore, any numerical procedure to generate them necessarily involves several \\textit{ad hoc} parameters due to the limitation of the available computation resources including the finite timestep of computation, the minimum mass of halos to be included in merger trees, and the maximum number of progenitors for each halo at each step.  The purpose of the paper is to perform a systematic investigation of possible artificial effects of the above-mentioned problems on merger tree realizations, and to re-examine the validity of the Monte-Carlo modeling. In particular, we focus on the extent to which the resulting merger trees reproduce the conditional one-point probability distribution function predicted by the extended Press--Schechter theory, which directly changes the fraction of cold gas.  Exactly for this reason, we adopt a conventional $\\Lambda$CDM model with the cosmological parameters $\\Omega_{0}=0.3$, $\\lambda_{0}=0.7$, $h=0.7$, $\\sigma_{8}=1.0$, and $\\Omega_{\\mathrm{B}}=0.015h^{-2}$ (e.g., Kitayama, Suto 1997; Kitayama et al.\\ 1998), and neglect star formation and a feedback effect for definiteness. ",
        "conclusions": "We attempted several convergence tests of the merger trees generated with the Monte-Carlo method. While this method provides a useful tool for modeling galaxy formation in a complementary manner to more intensive cosmological simulations with ad hoc recipes of galaxy formation (e.g., Cen, Ostriker 1992; Weinberg et al.\\ 1997; Yoshikawa et al.\\ 2001), the lack of an explicit expression for the joint distribution function of progenitors [equation~(\\ref{eq:jointprob})] requires one to put an additional assumption in practice. We confirmed that a repeated use of the \\textit{mass-weighted} conditional probability [equation~(\\ref{eq:eps-num})] reasonably reproduces the progenitor distribution predicted in the extended Press--Schechter theory if one adopts fairly small timesteps in redshift, $N_{\\rm step} \\sim 1000$, a factor of ten larger than a typical value used in previous work. We note, however, that one can alternatively achieve a similar result by fine-tuning the timestep as a function of $M_{1}$ (e.g., \\cite{SK99}) instead of equation~(\\ref{eq:zbin}), as we adopted here.  One may avoid the above problem also by using merger trees generated via $N$-body simulations (\\cite{gif99a}; Somerville et al.\\ 2001). In fact, they claim that the agreement between the $N$-body simulations and the Monte-Carlo method is good. Benson et al.\\ (2001) compared the SPH simulations and the Monte-Carlo modeling, and concluded that both agree with each other on the cold gas mass fraction and mass function of the halos.  While this comparison is encouraging, it is not yet clear if the lack of the joint distribution function of progenitors [equation~(\\ref{eq:jointprob})] in the Monte-Carlo modeling may not be essential. Thus, further detailed studies are definitely important to test the reliability of {\\it both\\/} $N$-body and the Monte-Carlo modeling in generating merger tree realizations.  \\vspace*{0.5cm}  We thank Kazuhiro Shimasaku and Tomonori Totani for discussions and suggestions in the early phase of this work.  This research was supported in part by the Grant-in-Aid from Monbu-Kagakusho, Japan (07CE2002, 12304009, 12640231).  T.K. gratefully acknowledges support from Research Fellowships of the Japan Society for the Promotion of Science for Young Scientists (7202).  \\onecolumn"
    },
    "0207/astro-ph0207006_arXiv.txt": {
        "abstract": "We present an re-analysis of the longest timescale gravitational microlensing event discovered to date: MACHO-99-BLG-22/OGLE-1999-BUL-32, which was discovered by both the MACHO and OGLE microlensing alert systems. Our analysis of this microlensing parallax event includes a likelihood analysis of the lens position based upon a standard model of the Galactic velocity distribution, and this implies that the lens could be a black hole of $\\sim 100\\msun$ at a distance of a few hundred parsecs in the Galactic disk or a massive stellar remnant (black hole or neutron star) in the Galactic bulge. Our new analysis includes data from the MACHO, GMAN, and MPS collaborations in addition to the OGLE data used in a previous analysis by Mao et al (2002). The crucial feature that distinguishes our analysis from that of Mao et al is an accurate constraint on the direction of lens motion and an analysis of the implications of this direction. ",
        "introduction": "\\label{intro}  The abundance of long timescale microlensing events towards the Galactic bulge has remained a puzzle \\citep{hangould-mspec,bennett-parbh} since shortly after the first systematic analyses of Galactic bulge microlensing events were reported \\citep{ogle-tau,macho-bulge45}. A long time scale event can be caused by a massive lens, a slow relative transverse velocity between the lens and source, or by large source star-lens and lens-observer separations. Thus, the excess of long timescale events could be caused by errors in our assumptions regarding the phase space distribution of the source stars or the lenses, or they could be caused by an unexpectedly large population of massive lenses. However, the phase space distribution of stellar mass objects in the Galactic disk and bulge is tightly constrained by observations, whereas little is known about the mass function of massive stellar remnants. Thus, it is possible that the excess of long timescale events is due to a population of black holes and neutron stars that is larger than expected.  This possibility has recently received observational support from the MACHO \\citep{bennett-parbh} and OGLE \\citep{ogle-99b22} collaborations which have identified individual candidate black hole microlensing events based upon the detection of the microlensing parallax effect \\citep{refsdal-par,gould-par1,macho-par1}. This refers to the detection of light curve features due to the Earth's accelerating motion around the Sun, and it is often detectable in long timescale microlensing events. To date, some 12 clear cases of this effect have been reported by the MACHO \\citep{macho-par1,bennett-parbh,becker-thesis}, OGLE \\citep{mao-par,ogle2000bul43,ogle-par2,ogle-99b22}, MOA \\citep{moa-par} and PLANET \\citep{planet-er2000b5} collaborations. These microlensing parallax events are useful because they allow one to learn more about the nature of the lens than can be learned for most microlensing events. For most microlensing events, the only measurable parameter which constrains interesting properties of the lens is the Einstein diameter crossing time, $\\that$,  which depends on the lens mass ($M$), distance ($D_\\ell$), and the magnitude of the transverse velocity ($\\vperp = |\\vperpbold |$). It is given by \\begin{equation} \\that = {2 R_E\\over \\vperp} = {4\\over \\vperp c} \\sqrt{GM D_\\ell (D_s - D_\\ell )\\over D_s} \\ , \\label{eq-that} \\end{equation} where $D_s$ refers to the distance to the source (typically $\\sim 8\\,$kpc for a bulge source), and $R_E$ is the radius of the Einstein Ring. In a microlensing parallax event, it is also possible to measure $\\vp$, which is the lens star's transverse speed projected to the Solar position, given by \\begin{equation} \\vpbold = \\vperpbold D_s/(D_s-D_\\ell) \\ . \\label{eq-vp} \\end{equation} The measurement of both $\\that$ and $\\vp$ give us two measurements for three unknowns ($M$, $\\vperp$, and $D_\\ell$), which allows us to solve for the mass as a function of distance: \\begin{equation} M = {\\vp^2 \\that^2 c^2 \\over 16 G} {D_s-D_\\ell \\over D_\\ell D_s} = {\\vp^2 \\that^2 c^2 \\over 16 G} {1-x \\over x D_s} \\ , \\label{eq-m} \\end{equation} where $x=D_\\ell/D_s$. For the MACHO-99-BLG-22/OGLE-1999-BUL-32 microlensing event, \\citep{ogle-99b22} have used Eq.~\\ref{eq-m} to argue that the lens mass is well above the measured mass of neutron stars ($\\sim 1.4\\msun$) for plausible values of $x$.  \\citet{bennett-parbh} have also made use of the measured direction of the projected velocity, $\\vpbold$, in a likelihood analysis employing the velocity and density distributions of a standard Galactic model. This allows a likelihood estimate for the lens distance, $x$, which can be converted to an estimate of the lens mass using Eq.~\\ref{eq-m}. This analysis has revealed two other candidate black hole lenses, MACHO-96-BLG-5 and MACHO-98-BLG-6, as well as several other events with possible neutron star lenses. In this paper, we add additional data to the analysis of the MACHO-99-BLG-22/OGLE-1999-BUL-32 microlensing event, and show that this results in tighter constraints on the microlensing parallax parameters. We then apply the likelihood analysis developed in \\citet{macho-par1} and \\citet{bennett-parbh} to estimate the likely distance and mass of the lens. ",
        "conclusions": "\\label{sec-con}  We have presented a microlensing parallax analysis of the MACHO-99-BLG-22/OGLE-1999-BUL-32 microlensing event using the publicly available OGLE and MACHO data along with MPS and GMAN data. The projected velocity indicated by the microlensing parallax fit has a direction that is nearly opposite of the direction of disk rotation. This suggests that the lens is either very close to us in the Galactic disk or in the Galactic bulge. The microlensing parallax mass-distance relationship, \\ref{eq-m}, indicates that the lens must be very massive if it is nearby, with a mass of order $\\sim 100\\msun$. If the lens is in the Galactic bulge, however, its mass is much smaller, $\\sim 4\\msun$. In both cases, the lens is likely to be a black hole. These conclusions have been quantified with a likelihood analysis similar to those presented by \\citet{macho-par1} and \\citet{bennett-parbh}. We should note that this analysis does depend on the assumption that stellar mass black holes are not born with a large velocity ``kick\\rlap.\" This assumption is supported by the available evidence \\citep{bh-kick}, but if this assumption is wrong, then the MACHO-99-BLG-22/OGLE-1999-BUL-32 could be a $\\sim 10\\msun$ black hole half way to the Galactic with a rotation speed of only about $60\\kms$.  We have not carried a Bayesian likelihood analysis with an assumed mass distribution prior as has been advocated by \\citet{agol} and also done by \\citet{bennett-parbh}. Such an analysis is of little use for this event because the lens mass is very likely to be $> 1.5\\msun$. Since a main sequence lens of this mass is excluded, we would have to provide a prior distribution of stellar remnants, but there is virtually nothing known about mass function of stellar remnants above $1.5\\msun$.  Our analysis agrees with the previous analysis of \\citet{ogle-99b22} and indicates that the MACHO-99-BLG-22/OGLE-1999-BUL-32 lens is very likely to be a black hole, but we show that a nearby, massive black hole of $\\sim 100\\msun$ and a low mass black hole of $\\sim 4\\msun$ in the Galactic bulge are two distinct possibilities that can explain the observed microlensing parallax parameters. When this event is combined with the two other candidate black hole microlenses presented by the MACHO Collaboration \\citep{bennett-parbh}, it appears that black holes may comprise a significant fraction of the mass of the Galactic disk and bulge, perhaps as large as $\\sim 10$\\%. Future observations with large ground based interferometers \\citep{vlti} should be able to accurately determine the mass of future black hole microlenses and determine the black hole mass fraction."
    },
    "0207/astro-ph0207320_arXiv.txt": {
        "abstract": "We present ground-based optical observations of \\grb{} starting 1.6 hours after the burst, as well as subsequent Very Large Array (VLA) and {\\it Hubble Space Telescope} (HST) observations.  The optical afterglow of \\grb{} is one of the faintest afterglows detected to date, and it exhibits a relatively rapid decay, $F_\\nu\\propto t^{-1.60\\pm 0.04}$, followed by further steepening.  In addition, a weak radio source was found coincident with the optical afterglow. The HST observations reveal that a positionally coincident host galaxy must be the faintest host to date, $R\\gtrsim 29.5$ mag.  The afterglow observations can be explained by several models requiring little or no extinction within the host galaxy, $A_V^{\\rm host}\\approx 0-0.9$ mag. These observations have significant implications for the statistics of the so-called dark bursts (bursts for which no optical afterglow is detected), which are usually attributed to dust extinction within the host galaxy.  The faintness and relatively rapid decay of the afterglow of \\grb{}, combined with the low inferred extinction indicate that some dark bursts are intrinsically dim and not dust obscured.  Thus, the diversity in the underlying properties of optical afterglows must be observationally determined before substantive inferences can be drawn from the statistics of dark bursts. ",
        "introduction": "\\label{sec:intro}  One of the main observational results stemming from five years of $\\gamma$-ray burst (GRB) follow-ups at optical wavelengths is that about $60\\%$ of well-localized GRBs lack a detected optical afterglow, (``dark bursts''; Taylor et al. 2000; Fynbo et al. 2001; Reichart \\& Yost 2001; Lazzati, Covino, \\& Ghisellini 2002)\\nocite{tbf+00,fjg+01,ry01,lcg02}.  In some cases, a non-detection of the optical afterglow could simply be due to a failure to image quickly and/or deeply enough.  However, there are two GRBs for which there is strong evidence that the optical emission should have been detected, based on an extrapolation of the radio and X-ray emission (Djorgovski et al. 2001a; Piro et al. 2002)\\nocite{dfk+01a,pfg+02}.  One interpretation in these two cases is that the optical light was extinguished by dust, either within the immediate environment of the burst or elsewhere along the line of sight (e.g.~Groot et al. 1998)\\nocite{ggv+98}.  An alternative explanation is a high redshift, leading to absorption of the optical light in the Ly$\\alpha$ forest.  However, the redshifts of the underlying host galaxies of these GRBs are of order unity (Djorgovski et al. 2001a; Piro et al. 2002)\\nocite{dfk+01a,pfg+02}.  Several authors have recently argued that a large fraction of the dark bursts are due to dust extinction within the local environment of the bursts (e.g.~Reichart \\& Yost 2001; Lazzati et al. 2002; Reichart \\& Price 2002)\\nocite{ry01,lcg02,rp02}, but other scenarios have also been suggested (e.g.~Lazzati et al. 2002).  Moreover, it has been noted that regardless of the location of extinction within the host galaxy, the fraction of dark bursts is a useful upper limit on the fraction of obscured star formation (Kulkarni et al. 2000; Djorgovski et al. 2001b; Ramirez-Ruiz, Trentham, \\& Blain 2002; Reichart \\& Price 2002)\\nocite{kbb+00,dkb+01,rtb02,rp02}.  However, from an observational point of view, we must have a clear understanding of the diversity of afterglow properties before extracting astrophysically interesting inferences from dark bursts. For example, afterglows which are faint or fade rapidly (relative to the detected population) would certainly bias the determination of the fraction of truly obscured bursts.  In this vein, Fynbo et al. (2001), noting the faint optical afterglow of GRB\\,000630, argued that some dark bursts are due to a failure to image deeply and/or quickly enough, rather than dust extinction.  Here we present optical and radio observations of \\grb{}, an afterglow that would have been classified dark had it not been for rapid and deep searches.  Furthermore, \\grb{} is an example of an afterglow, which is dim due to the combination of intrinsic faintness and a relatively fast decline, and not strong extinction. ",
        "conclusions": "\\label{sec:conc}  Regardless of the specific model for the afterglow emission, the main conclusion of \\S\\ref{sec:model} is that the optical afterglow of \\grb{} suffered little or no dust extinction.  Still, this afterglow would have been missed by typical searches undertaken even as early as 12 hours after the GRB event.  As shown in Fig.~\\ref{fig:rlims}, about $70\\%$ of the searches conducted to date would have failed to detect an optical afterglow like that of \\grb{}.  This is simply because the afterglow of \\grb{} was faint and exhibited relatively rapid decay.  From Fig.~\\ref{fig:alpha_f0} we note that \\grb{} is one of the faintest afterglows detected to date (normalized to $t=1$ day), and while it is not an excessively rapid fader, it is in the top $30\\%$ in this category.  Thus, the afterglow of \\grb{}, along with that of GRB\\,000630 (Fynbo et al. 2001; Fig.~\\ref{fig:alpha_f0}), indicates that there is a wide diversity in the brightness and decay rates of optical afterglows.  In fact, the brightness distribution spans a factor of about 400, while the decay index varies by more than a factor of three.  Coupled with the low dust extinction in the afterglow of \\grb{}, this indicates that some dark bursts may simply be dim, and not dust obscured.  Given this wide diversity in the brightness of optical afterglows, it is important to establish directly that an afterglow is dust obscured.  This has only been done in a few cases (\\S\\ref{sec:intro}). Therefore, while {\\it statistical} analyses (e.g.~Reichart \\& Yost 2001) point to extinction as the underlying reason for some fraction of dark bursts, and may even account for an afterglow like that of \\grb{}, it is clear that observationally the issue of dark bursts is not settled, and the observational biases have not been traced fully.  Since progress in our understanding of dark bursts will benefit from observations, we need consistent, rapid follow-up of a large number of bursts to constrain the underlying distribution, as well as complementary techniques which can directly measure material along the line of sight.  This includes X-ray observations which allow us to measure the column density to the burst (Galama \\& Wijers 2001)\\nocite{gw01}, and thus infer the type of environment, and potential extinction level.  Along the same line, radio observations allow us to infer the synchrotron self-absorption frequency, which is sensitive to the ambient density (e.g.~Sari \\& Esin 2001)\\nocite{se01}; the detection of radio emission, as in the case of \\grb{}, implies a density $n\\lesssim 10^2$ cm$^{-3}$.  Finally, prompt optical observations, as we have carried out in this case, may uncover a larger fraction of the dim optical afterglows, and provide a better constraint on the fraction of truly obscured bursts."
    },
    "0207/astro-ph0207116_arXiv.txt": {
        "abstract": "We report proper motion dispersions for stars in the direction of two fields of the Galactic bulge, using HST/WFPC2 images taken six years apart.  Our two fields  are Baade's Window $(l,b)=(1.13^\\circ,-3.77^\\circ)$ and Sgr I $(l,b)=(1.25^\\circ ,-2.65^\\circ)$. Our proper motion dispersions are in good agreement with prior ground- and space-based proper motion studies in bulge fields, but in contrast to some prior studies, we do not exclude any subset of stars from our studies.  In Baade's Window, we find the $l$ and $b$ proper motion dispersions are 2.9 and 2.5 mas/yr, while in Sgr I, they are 3.3 and 2.7 mas/yr, respectively.  For the first time, we can clearly separate the foreground disk stars out from the bulge because of their large mean apparent proper motion.  The population with non-disk kinematics (which we conclude to be the bulge) has an old main sequence turnoff point, similar to those found in old, metal rich bulge globular clusters while those stars selected to have disk kinematics lie on a fully populated main sequence. Separating main sequence stars by luminosity, we find strong evidence that the bulge population is rotating, largely explaining observations of proper motion anisotropy in bulge fields.  Because we have isolated such a pure sample of stars in the bulge, we have one of the clearest demonstrations that the old stellar population of the inner bulge/bar is in fact rotating. ",
        "introduction": "The bulge of our Galaxy is interesting for a number of reasons. It is the nearest galactic bulge to the sun, and represents a unique place to study the stellar populations and stellar dynamics of such objects in detail. Such analysis can provide important information for our understanding of how bulges formed their stellar populations, what gravitational potential they sit in, and how they came to have the structure they do.  In the case of our Galactic bulge, stellar motions are of great interest both in testing the hypothesis that our bulge is in fact a bar (e.g.\\ \\citealt{zhaoetal96}) and in the modeling of microlenses, the overwhelming majority of which are seen in the direction of the bulge \\citep{alcock00}.  There are several complications in trying to build an understanding of our bulge. Close to the galactic plane, extinction by foreground dust is large and uneven (even on very small scales) and effectively limits optical work to a few ``windows'' of low and relatively uniform extinction.  More troubling is the difficulty that the stellar population seen in these directions samples everything along the line of sight. Simple exponential models for the disk, for example, predict that in Baade's Window at least half of the stars visible are actually disk stars, and not bulge stars. Employing photometry alone, it is not possible to effectively sort the populations; bulge and disk populations overlap in color, especially near the turnoff \\citep{holtz98} greatly complicating the use of HST-derived CMDs for age determination.  Blue stragglers extending brighter than the turnoff in an old population overlap with the main sequence locus of a young population.  Among the first stellar populations imaged in 1994 with the repaired WPFC2 on board HST were fields near two of the low extinction regions originally identified by Walter Baade: what is now known as Baade's Window $(l,b)=(1.13^\\circ,-3.77^\\circ)$,  and Sgr I $(l,b)=(1.25^\\circ ,-2.65^\\circ)$.   A discussion of the photometry and luminosity function of the Baade's Window field is given in \\citet{holtz98}.  We noticed that the fields were ripe for a revisit with the aim of measuring proper motions and we proposed successfully (GO-8250).  Although HST has been used to measure proper motions of field stars in the rough vicinity of the Galactic Center, these two fields are of special interest because they lie well within the {\\sl COBE} bulge \\citep{dwek95} and, in the case of Baade's Window, abundances are measured \\citep{mcw94} and many other studies have been done.  Combining proper motions with precision photometry might make it possible to separate the observed populations based on their kinematics.  The only previous study of bulge proper motions in Baade's Window was based on photographic plates \\citep{spaen92} and reported motions only of stars thought to be candidate red giants in the bulge, by \\citet{arp65}. By systematically excluding the bluer disk stars, and by only measuring a few hundred of the brightest giants, this study, while pioneering, leaves much of the problem ripe for inquiry.  Hence our decision to obtain second epoch images of both bulge fields, using WFPC2 on board HST. ",
        "conclusions": "With HST photometry and proper motions determined with high enough precision, it is possible to separate the disk and bulge populations by their kinematics alone.  The long standing question regarding the nature of the blue main sequence extension in the bulge field population is settled: the great majority of those stars evidently belong to the foreground disk.  When these stars are excluded, the old turnoff population in the bulge remains, and no measureable population of blue stragglers or intermediate age stars is present. This was first demonstrated for the Baade's Window field by \\cite{ortolani95}. \\cite{feltzgil00} came to the same conclusion based on counts of stars brighter and fainter than the turnoff point in their WFPC2 data. However, a proper accounting for the foreground disk has been a persistant issue, and our application of the kinematic data strengthens greatly the conclusion that the bulge is dominated by old stars.  We take our study a step further, and find direct evidence for the rotation of the bulge population. The observed proper motion anisotropy of bulge stars is largely caused by the line-of-sight gradient of the rotation of the bulge; when this is removed a nearly isotropic velocity distribution of the bulge stars results. The velocity dispersion declines from Sgr~I to Baade's Window.  When appropriate samples in Baade's Window are compared, our proper motion dispersions agree with those found by \\citet{spaen92} and by \\cite{feltz01}.  Our results compare well with predictions of the \\citet{zhao96} and \\citet{zhaoetal96} bulge model, but since our proper motion samples are more than an order of magnitude larger than existing radial-velocity and proper-motion samples of bulge stars, and extend well below the turnoff, the time is ripe for more involved modelling. Extension of this work to further bulge fields (e.g. \\citealt{zocc01}), and combination of these results with spectroscopy (for radial velocities, metallicity, and improved distance estimates) should open the way for a new chapter in our study of the Galactic bulge."
    },
    "0207/astro-ph0207266_arXiv.txt": {
        "abstract": "We have carried out the wind analysis of six A-type supergiants in NGC 300. The derived Wind Momentum-Luminosity Relationship is compared with that of Galactic and M31 blue supergiants and with theoretical models. ",
        "introduction": "The theory of radiatively driven winds predicts a relationship between the modified wind momentum and a power of stellar luminosity: $$\\log{(\\dot{M}\\,v_\\infty\\,R)} = x\\log{L} + const$$  \\noindent wher $\\dot{M}$ is the mass-loss rate, $v_\\infty$ the wind terminal velocity and $R$ the stellar radius.  Empirical verifications of this Wind Momentum-Luminosity Relationship (WLR) have so far been carried out for O stars in the Milky Way and the Magellanic Clouds by Puls et al.~(1996), and for Galactic and M31 B and A supergiants by Kudritzki et al.~(1999). The brightest A-type supergiants in galaxies are extremely bright, with $M_V$ up to $-10$, making the WLR a potential extragalactic distance indicator. However, extensive work still needs to be carried out on the empirical calibration of the relation and on the theoretical modeling of the effects of stellar metallicity and spectral type. ",
        "conclusions": ""
    },
    "0207/astro-ph0207050_arXiv.txt": {
        "abstract": "Thanks to its unprecedented spatial resolution, the Hubble Space Telescope has ended a 20-year long stalemate by detecting the dynamical signature of nuclear supermassive black holes (SBHs) in a sizeable number of nearby galaxies. These detections have revealed the existence of a symbiotic relationship between SBHs and their hosts, changing the way we view SBH and galaxy formation.  In this contribution I review which are the most pressing outstanding issues in SBH research, and what are the technological requirements needed to address them. ",
        "introduction": "The study of supermassive black holes is one of the areas of modern astrophysics which has benefited most from the launch of HST. After two decades of tantalizing but inconclusive ground-based studies, the HST/FOS observations of M87 (Harms et al.  1994) and NGC 4261 (Ferrarese et al.  1996) provided the first firm measurements of SBH masses in galactic nuclei.  In the years that followed, FOS and STIS data lead to detections in ten additional galaxies (Bower et al. 1998; van der Marel \\& van den Bosch 1998; Ferrarese \\& Ford 1999; Emsellem et al.  1999; Cretton \\& van den Bosh 1999; Verdoes Kleijn et al.  2000; Gebhardt et al.  2000a; Joseph et al.  2001; Barth et al. 2001; Sarzi et al.  2001).  The superiority of HST over ground based facilities in this field is easily understood.  Only dynamical evidence, either from gas or stellar kinematics, can yield compelling proof of the existence of SBHs.  With rare exceptions (e.g.  M31, M87), ground based telescopes lack the spatial resolution necessary to resolve the SBH ``sphere of influence'', i.e.  the region of space within which the SBH gravitational influence dominates that of the surrounding stars:  \\begin{equation}r_h = G\\mh/\\sigma^2 \\sim 11.2(\\mh/10^8~{\\rm M_{\\odot}}) / (\\sigma /200~ {\\rm km ~s^{-1}})^2 {\\rm ~pc},\\end{equation}  \\noindent with $\\sigma$ the stellar velocity dispersion and $\\mh$ the SBH mass.  Resolving $r_h$ is a necessary condition for a SBH detection to be made; not meeting it leads to spurious detections and biased masses (Merritt \\& Ferrarese 2001a).  With over a dozen secure measurements, it has become possible to search for correlations between $\\mh$ and the overall properties of the host galaxies.  The first relation to emerge was one between $\\mh$ and the blue luminosity $L_B$ of the surrounding bulge (Kormendy \\& Richstone 1995).  A much tighter correlation was subsequently discovered between $\\mh$ and the bulge stellar velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al.  2000b):  \\begin{equation}\\mh = \\beta {\\left({\\sigma} \\over {200 {\\rm ~km~s^{-1}}}\\right)}^\\alpha.\\end{equation}  \\noindent with $\\alpha = 4.58 \\pm 0.52$ and $\\beta = (1.66 \\pm 0.32) \\times 10^8$ \\msun~ (Ferrarese 2002a).  More recently, evidence has emerged that a fundamental relation might exist between $\\mh$ and the mass $M_{DM}$ of the dark matter halos in which the SBHs presumably formed (Ferrarese 2002b):  \\begin{equation}{{\\mh} \\over {10^8~{\\rm M_{\\odot}}}} \\sim 0.10 {\\left({M_{DM}} \\over {10^{12}~{\\rm  {\\rm M_{\\odot}}}}\\right)}^{1.65}\\end{equation}  The above relations have proven invaluable in the study of SBH demographics (Merritt \\& Ferrarese 2001b; Ferrarese 2002a; Yu \\& Tremaine 2002) and have generated intense activity on the theoretical front (Haehnelt, Natarajan \\& Rees 1998; Silk \\& Rees 1998; Cattaneo, Haehnelt \\& Rees 1999; Adams, Graff \\& Richstone 2000; Monaco et al. 2000; Haehnelt \\& Kauffmann 2000; Wyithe \\& Loeb 2002).  At the same time, new questions have arisen, and with them the need for further observational constraints.  In this contribution, I will address three such questions, and identify the technological requirements necessary to answer them.  \\begin{itemize}  \\item {\\it What are the characteristics of the $\\ms$ relation?  Does the slope and/or normalization of the relation depend on Hubble type, environment, and/or redshift?} As theoretical models are refined, tighter observational constraints will be required.  Most of the SBHs detected to date are in the $10^8 \\lae \\mh \\lae 10^9$ range.  The $\\ms$ relation is not sampled below $10^6$ \\msun, and badly sampled for $10^6 \\lae \\mh \\lae 10^7$ \\msun.  There are few spiral galaxies represented, all of which are early type, and only two galaxies well beyond 30 Mpc.  \\item {\\it Are binary supermassive black holes long lived?} The existence and lifespan of binary SBHs can have dramatic consequences, from shaping the morphology and dynamics of the resulting galaxy (Milosavljevic \\& Merritt 2001; Milosavljevic et al.  2002; Ravindranath et al.  2002; Yu 2002) to destroying nuclear dark matter halo cusps (Merritt et al.  2002).  \\item {\\it How small can nuclear BHs be?  Are there nuclear BHs in globular clusters?} There is no dynamical evidence for ``intermediate mass'' black holes (IBHs) in the $\\mh \\sim 10^2 - 10^6$ \\msun~range, although their existence in the off-nuclear regions of some starburst galaxies is supported by energetic arguments (Fabbiano et al.  2001; Matsumoto et al.  2001).  There is also no dynamical evidence that BHs are formed in the nuclei of globular clusters (van der Marel et al. 2000; Gebhardt et al.  2000).  However, whether such black holes exist is critical for our understanding of how SBHs form.  In ``top-down'' self-regulating models that trace the formation of SBHs to the very early stages of galaxy formation, there is a natural lower limit of $\\sim 10^6$ \\msun~to $\\mh$ (e.g.  Loeb 1993; Silk \\& Rees 1998; Haehnelt, Natarajan \\& Rees 1998).  On the other hand, in ``bottom-up'' models nuclear SBHs are formed by the merging of IBHs. These are deposited at the galactic center as the globular clusters in which they originally formed spiral in due to dynamical friction (Portegies Zwart \\& McMillan 2002; Ebisuzaki et al.  2001).  In the latter scenario, no physical reason would prevent the formation of SBHs with $\\mh \\lae 10^6$ \\msun.  \\end{itemize} ",
        "conclusions": ""
    },
    "0207/astro-ph0207099_arXiv.txt": {
        "abstract": "We describe a near-infrared imaging survey of Globule 2 in the Coalsack. This Bok globule is the highest density region of this southern hemisphere molecular cloud and is the most likely location for young stars in this complex.  The survey is complete for $K$ $<$ 14.0, $H$ $<$ 14.5, and $J$ $<$ 15.5, several magnitudes more sensitive than previous observations of this globule.  From the large number of background stars, we derive an accurate near-infrared extinction law for the cloud. Our result, $E_{J-H}/E_{H-K}$ = $2.08 \\pm 0.03$, is significantly steeper than results for other southern clouds. We use the $J-H$/$H-K$ color-color diagram to identify two potential young stars with $K$ $<$ 14.0 in the region. We apply $H$-band star counts to derive the density profile of the Coalsack Globule 2 and use a polytropic model to describe the internal structure of this small cloud. For a gas temperature T $\\sim$ 15 K, this globule is moderately unstable. ",
        "introduction": "The Coalsack is a conspicuous, nearby \\citep[d $\\sim$ 180 pc;][]{fra89} dark cloud located close to the Galactic plane in the southern Milky Way (l $\\sim$ 303$^o$, b  $\\sim$ 0$^o$). \\citet{nym89} estimated a total mass of $\\sim$ 3550 M$_{\\odot}$ from $^{12}$CO data covering an area of $\\sim$ 15 deg$^2$ in the direction of this complex. Star counts toward the same area indicate an average optical extinction of $A_V \\sim$ 5 mag \\citep{gre88,cam99}. However, the extinction is $A_V \\sim$ 20 mag or larger in several small dark globules \\citep{tap73,har86,bou95a}.  This structure is similar to other well-known clouds such as Taurus, Chamaeleon I, and Lupus that are actively forming stars \\citep{ceba84,bou98,argo99,cam99}.  Despite this similarity, previous searches of the Coalsack have failed to detect signs of recent star-formation \\citep{wea74,sch77,rei81,nym89}.  Recently \\citet{kat99} have observed the Coalsack in the $^{13}$CO and C$^{18}$O J$=$1--0 lines. The $^{13}$CO data show a massive cloudlet ($\\sim$ 200 M$_{\\odot}$) located at the western edge of the complex, coinciding with the region of the largest optical extinction \\citep[see][]{cam99}. In addition, these authors identified five C$^{18}$O cores \\citep[see also][]{vil94}, corresponding to the positions of well-known optically dark globules \\citep{tap73, har86, bou95a}. These cores have high column densities, typical of star-forming cores, and masses between 4--10 M$_{\\odot}$. However, there are no {\\it IRAS} sources, and thus no evidence of recent star-formation, in these cores \\citep{kat99}.  Tapia's Globule 2 ($\\alpha$ $=$ 12$^h$ 31.5$^m$, $\\delta$ $=$ $-$63$^o$ 44.5$'$; 2000.0) is the densest (n(H$_2$) $\\sim$ 4 $\\times$ 10$^3$ cm$^{-3}$) and most massive ($\\sim$ 10 M$_{\\odot}$) of the Kato et al. cores \\citep[see also][]{bou95b}.  It is an obvious, roughly circular patch \\citep[$\\sim$ 6$'$ radius;][]{bou95a} of extinction on shallow optical and near-infrared (near-IR) surveys. \\citet{jon80} observed an area of $\\sim$ 850 arcmin$^2$ to $K=$ 9.5; \\citet{jon84} scanned an area of 2$'$ $\\times$ 2$'$, 70\\% complete to $K=$ 13.7, with the central 1$'$ $\\times$ 1$'$ region 70\\% complete to $K=$ 14.7. Neither of these surveys detected candidate near-IR excess stars. At $K=$ 9.5, \\citet{jon80, jon84} could detect $\\sim$ 1 M$_{\\odot}$ main sequence stars at the distance of the Coalsack.  In the central 1$'$ $\\times$ 1$'$ region, 70\\% complete to $K=$ 14.7, this detection limit is $\\sim$ 0.4 M$_{\\odot}$, assuming A$_K$ $\\sim$ 2 \\citep[see][]{hemc93,del00}.  The Coalsack Globule 2 has also been observed by the {\\it Midcourse Space Experiment}  \\citep[MSX;][]{pri01}. Fits format images and a source catalog in four mid-infrared bands, A(8.28 $\\mu$m), C(12.13 $\\mu$m), D(14.65 $\\mu$m), and E(21.34 $\\mu$m),  are available through the NASA/IPAC IR Science Archive (IRSA)\\footnote{http://irsa.ipac.caltech.edu.}. Nine mid-infrared sources lie within \\hbox{$\\sim$ 11$'$} from the optical center of the globule. One of these is located $\\sim$ 2.8$'$ SW from the center of the most massive Kato et al. C$^{18}$O cores. This different is comparable to the 2$'$ resolution of the molecular line maps.  Here, we describe a near-IR imaging survey of $\\sim$ 15$'$$\\times$15$'$ centered on Globule 2. The observations are uniformly complete to $K=$ 14.0, $H=$ 14.5, and $J=$ 15.5. If $A_K \\sim$ 0--2 mag, these data are sensitive to $\\sim$ 0.2--0.5 M$_{\\odot}$ main sequence stars \\citep[see][]{hemc93,del00}. We detect $\\sim$ 6500 sources to the detection limits of $K=$ 16.5, $H=$ 17.0, and $J=$ 18.0.  We restrict our analysis to $K=$ 14.0 ($\\sim$ 2500 sources) where the photometric errors are relatively small ($<$ 0.08 mag).  In \\S 2 we describe the observations and data reduction. In \\S 3 we derive a reliable extinction curve for background stars in the Globule 2 region and use this result in \\S 4 to search for potential young stellar objects in our survey region based on their locations in the $J-H$/$H-K$ diagram. We identify only two objects with $K$ $<$ 14.0 and near-IR excesses characteristic of young stellar objects. In \\S 5 we use $H$-band stars counts to derive the density profile of this small cloud and compare it with the predictions of polytropic models. Assuming T $\\sim$ 15 K, Globule 2 is moderately unstable.  We conclude with a brief summary in \\S 6. ",
        "conclusions": "Our near-IR survey of Globule 2, the highest density and more massive core in the Coalsack complex, leads to three primary conclusions.  The slope of the near-IR extinction law in the cloud, $E_{J-H}/E_{H-K}$ = 2.08 $\\pm$ 0.03, is much steeper than in $\\rho$ Ophiuchi or Chamaeleon I.  If real, the trend of increasing slope with decreasing star-formation activity in these three clouds suggests changes in grain chemistry with extinction or star formation activity.  Based on near-IR excess emission, we detect two pre-main sequence candidates with $K < 14$ in the vicinity Globule 2.  If confirmed as pre-main sequence stars with optical or near-IR spectroscopy, these are the first pre-main sequence stars discovered in the region.  The globule contains no known IRAS sources \\citep{nym89,bou95a} or other pre-main sequence candidates. Our candidates are too faint for detection with IRAS.  The low success rate in finding pre-main sequence stars in the Globule 2 region illustrates the low activity of the Coalsack as stellar nursery, in agreement with previous investigations \\citep{wea74,sch77,rei81,nym89}.  We use $H$-band star counts to derive the density profile of this Bok globule. For a `typical' gas temperature T $\\sim$ 15 K, model fits suggest this small cloud is moderately unstable, with a Bonnor critical parameter $\\xi_{\\rm max}$ $=$ 7.0 $\\pm$ 0.3.  The mass derived from these models, {$M = 4.5 M_{\\odot}$}, agrees with estimates derived from the CO column density."
    },
    "0207/hep-ph0207035_arXiv.txt": {
        "abstract": "The 3--dimensional (3--D) calculation of the atmospheric neutrino flux by means of the \\FLUKA{} Monte Carlo model is here described in all details, starting from the latest data on primary cosmic ray spectra. The importance of a 3--D calculation and of its consequences have been already debated in a previous paper. Here instead the focus is on the absolute flux. We stress the relevant aspects of the hadronic interaction model of \\FLUKA{} in the atmospheric neutrino flux calculation. This model is constructed and maintained so to provide a high degree of accuracy in the description of % particle production.  The accuracy achieved in the comparison with data from accelerators and cross checked with data on particle production in atmosphere certifies the reliability of shower calculation in atmosphere.  The results presented here can be already used for analysis by current experiments on atmospheric neutrinos. However they represent an intermediate step towards a final release, since this calculation does not yet include the bending of charged particles in atmosphere. On the other hand this last aspect, while requiring a considerable effort in a fully 3--D description of the Earth, if a high level of accuracy has to be maintained, does not affect in a significant way the analysis of atmospheric neutrino events. ",
        "introduction": "\\label{sec:intro} Reliable calculations of flux of secondary particles in atmosphere, produced by the interactions of primary cosmic rays, are essential for the correct interpretation of the large amount of experimental data produced by experiments in the field of astroparticle physics. The increasing accuracy of modern experiments demands also an improved quality of the calculation tools. The most important example in this field is the analysis of the experimental results on atmospheric neutrinos from Super--Kamiokande\\cite{superk}, MACRO\\cite{macro} and Soudan2\\cite{soudan2}, which gave the first robust evidence in favor of neutrino oscillations. The interpretation in terms of the mixing parameters is affected by different sources of systematic errors, and the theoretical uncertainties on neutrino fluxes and cross sections constitute a significant fraction of them. This has stimulated different efforts to improve the existing flux calculations. Among them, one is the calculation based on the \\FLUKA{} Monte Carlo code\\cite{fluka}. Such a work was indeed started before\\cite{taup97} the Super--Kamiokande results, in the framework of design work for the ICARUS experiment\\cite{icarus} and of the analysis of MACRO experiment at Gran Sasso. The main motivation to propose the \\FLUKA{} based calculation was the idea that, in order to accomplish the goals summarized before, the highest degree of detail and accuracy should be accomplished.  In particular the \\FLUKA{} code is known for the accuracy of its particle production model in hadronic interactions, which is extensively benchmarked against accelerator data. The first important achievement obtained using \\FLUKA{} concerned the relevance of 3--D geometry in the flux calculations and was presented in ref.\\cite{flukanu}. However, at the time of that work it was not yet possible to make statements on the absolute values of neutrino fluxes. This is instead the purpose of the present paper, which is intended to provide a complete reference for the \\FLUKA{} neutrino flux.  In the following, we shall first summarize the main ingredients of the simulation, discussing with some detail the primary cosmic ray spectrum and the geometrical set--up. Then, in a next section, we shall discuss the physics models of \\FLUKA{} presenting a set of comparisons between data and predictions to demonstrate the validity of the models themselves. The final results on the flux are presented and discussed, also by means of a comparison to other calculation results. Neutrino fluxes are calculated for 3 relevant experimental sites (Kamiokande, Gran Sasso and Soudan) and representative tables are given here for the Kamiokande site. More detailed tables for all the 3 sites are available on the web\\cite{flukatab}. The energy region considered in this work (0.1$\\div$200 GeV) covers in practice the production of the event topologies of of sub-GeV and multi-GeV events as detected in Super--Kamiokande.  In the conclusions we shall debate the level of systematic error in  the predictions and mention the items that will require further improvements. ",
        "conclusions": "\\label{sec:concl} The first phase of atmospheric neutrino flux calculation using the \\FLUKA{} code is here concluded. We think that these results, apart from the question of the 3-D geometry, assessed already in \\cite{flukanu}, represent the first systematic attempt to explore the impact of a refined hadronic interaction model, which is capable at the same time to reproduce a wide range of accelerator and cosmic ray data. This work has stimulated a serious debate inside the scientific community about the validity of the previous ``traditional'' calculations, and we consider the recent convergence towards a lower normalization of neutrino fluxes (for the same or similar primary spectrum) as an important achievement and recognition of the importance of using accurate interaction models for cosmic ray calculations. Along this line other attempts have been followed, like the works of ref.\\cite{Tserk,Fior,Wentz,Liu,Plyaskin}. Some of these attempts are, in our opinion, biased again by the choice of interaction models which are not precise. In particular we refer to those which are based upon the use of the old GHEISHA model\\cite{gheisha}, a parametrized code which fails in giving proof of reliability in reproducing particle production properties, as recently shown for instance in the framework of ALICE experiment at LHC\\cite{alice}. The calculation of ref.\\cite{Tserk}, which is based on GEANT-FLUKA, cannot be compared to ours, since the FLUKA package contained in GEANT-3 is an old and incomplete version of the present FLUKA. In particular it does not contain the fundamental PEANUT section and the high energy part (above 5 GeV) is now considered obsolete. The work of ref.\\cite{Plyaskin} is instead originally biased by a technical error in the normalization, as recently communicated by the author\\cite{plya2}.  As mentioned above, these results cannot yet be considered as completely final, since we are aware that a fully certified calculation must include the geomagnetic field also during shower development. However this has to be done using the most accurate description of this field, avoiding approximations, and this will be the object of our next development. Furthermore, the calculation of neutrino fluxes is also going to be extended at higher energies, as soon as it will be released a next extension of the \\FLUKA{} model in order to deal with nucleus--nucleus interactions up to extreme high energies\\cite{flukaV}.  The importance of reducing as much as possible the theoretical uncertainties in the calculation of these fluxes may have limited impact in the analysis of present experimental results concerning neutrino oscillations where normalization is left as a free parameter so that results are not dependent on it. The matter can be different in the framework of a 3--flavor scenario, where sub-GeV electron neutrinos acquire some weight: there in fact, one expects to see the effects of interference terms involving $\\theta_{12}$, if the LMA solution for solar neutrino turns out to be the right one, as confirmed by the recent SNO results\\cite{sno}. From the experimental point of view, the ICARUS experiment could be the one who can investigate with low or negligible systematic error the sector of low energy electron neutrinos in the atmospheric flux."
    },
    "0207/astro-ph0207670_arXiv.txt": {
        "abstract": "s{XENON is a novel liquid xenon experiment concept for a sensitive dark matter search using a 1-tonne active target, distributed in an array of ten independent time projection chambers. The design relies on the simultaneous detection of ionization and scintillation signals in liquid xenon, with the goal of extracting as much information as possible on an event-by-event basis, while maintaining  most of the target active. XENON is expected to have effective and redundant background identification and discrimination power, higher than 99.5\\%, and to achieve a very low threshold, on the order of 4~keV visible recoil energy. Based on this expectation and the 1-tonne mass of active xenon, we project a sensitivity of 0.0001~events/kg/day, after 3~yr operation in an appropriate underground location. The XENON experiment has been recently proposed  to the National Science Foundation (NSF) for an initial development phase leading to the development of the 100 kg unit module.} ",
        "introduction": "Substantial astronomical evidence shows that at least 90\\% of the mass in the universe is dark, and that most of it is non-baryonic in nature (see e.g. reviews \\cite{trimble87,primack88,tremaine92,jungman96}).  Dark matter plays a central role in current structure formation theories, and its microscopic properties have a significant impact on the spatial distribution of mass, galaxies and clusters. Unraveling the nature of dark matter is therefore of critical importance. Several lines of arguments indicate that the dark matter consists of Weakly Interacting Massive Particles (WIMPs), a well-motivated example of which is the neutralino, the lightest supersymmetric particle. Direct detection, via elastic scattering of a WIMP on a suitable target, offers the hope of studying the dark matter properties in detail, and shedding light on particle physics beyond the Standard Model.   In spite of the experimental challenges, a number of efforts worldwide are actively pursuing to directly detect WIMPs with a variety of targets and approaches. One approach is to decrease the radioactive background to extreme low levels, using a high purity Germanium target and detector, with careful selection of surrounding materials \\cite{klapdor01,HDMS99,GENIUS98}. A second approach, followed by the DAMA \\cite{DAMA00} and the UKDM NAIAD \\cite{spooner00} groups, has been to use large NaI scintillators with pulse shape background discrimination. The third experimental approach relies on more powerful discrimination methods, using various schemes to extract as much information as possible from the target-detector. To this class belong the cryogenic detectors based on the simultaneous measurement of ionization and phonons in crystals of Ge or Si, as used by the CDMS experiment \\cite{abusaidi00} and the EDELWEISS experiment \\cite{benoit01}, or phonons and scintillation light in CaWO$_4$ crystals as used by the CRESST experiment~\\cite{bravin99}. % Experiments based on the simultaneous detection of ionization and scintillation light in liquid xenon (LXe) belong to the same class and offer an equally promising approach to direct detection of WIMPS in large scale targets. Two experiments, ZEPLIN~II\\cite{HWang:00:ZEPLINII} and III\\cite{ZEPLINAXE}, with 30~kg and 6~kg of Xe mass, respectively, are currently being developed as part of the UKDM LXe program. Scale-up to the 1-tonne level is in the planning or proposal stage~\\cite{ZEPLINAXE,ZEPLINIV}.  The XENON project, recently proposed to NSF for an initial development phase, is an alternative concept for a 1-tonne LXe experiment, to be located in the National Underground Science Laboratory (NUSL), under discussion in the US. The goal of the XENON experiment is to achieve a factor of 30 higher sensitivity than that projected for CDMS~II~\\cite{CDMS_Soudan} in the US and other experiment in Europe (e.g. EDELWEISS). This sensitivity increase is needed to probe the lowest SUSY predictions for the neutralino. With 1-tonne target mass, a visible energy threshold of 4 keV and a background discrimination factor much better than 99.5\\%, XENON projected sensitivity is 0.0001~events/kg/day after 3~yr operation. ",
        "conclusions": ""
    },
    "0207/astro-ph0207393_arXiv.txt": {
        "abstract": "{ Recent progress in the theory of solar and stellar dynamos is reviewed. Particular emphasis is placed on the mean-field theory which tries to describe the collective behavior of the magnetic field. In order to understand solar and stellar activity, a quantitatively reliable theory is necessary. Much of the new developments center around magnetic helicity conservation which is seen to be important in numerical simulations. Only a dynamical, explicitly time dependent theory of $\\alpha$-quenching is able to describe this behavior correctly. ",
        "introduction": "Starspot activity is presumably driven by some kind of dynamo process. Many stars show magnetic field patterns extending over scales of up to $30^\\circ$ in diameter. The commonly used tool to model such magnetic activity is the mean-field dynamo. Although mean-field theory has been used over several decades there have recently been substantial developments concerning the basic nonlinearity of dynamo theory. It is the purpose of this review to highlight these recent developments in the light of applications to stars. ",
        "conclusions": "Magnetic helicity seems to play a much more prominent role than what has been anticipated until recently. It has become clear that $\\alpha$ must satisfy an explicitly time-dependent equation. The dynamical $\\alpha$-quenching theory has significant predictive power: it describes the different quenching behaviors for helical and nonhelical fields, the value of the magnetic Reynolds number is explicitly incorporated, and the magnetic helicity equation is satisfied exactly at all times. So far, no departures between this theory and the simulations have been found. A major restriction of the theory in its present form is however the inability to handle cases with spatially nonuniform $\\alpha$-effect."
    },
    "0207/astro-ph0207446_arXiv.txt": {
        "abstract": "We present optical broadband ($B$ and $R$) observations of the Seyfert 1 nucleus NGC~3516, obtained at Wise Observatory from March 1997 to March 2002, contemporaneously with  X-ray 2--10 keV measurements with \\rxte. With these data we increase the temporal baseline of this dataset to 5 years, more than triple to the coverage we have previously presented for this object. Analysis of the new data does not confirm the 100-day lag of X-ray behind optical variations, tentatively reported in our previous work. Indeed, excluding the first year's data, which drive the previous result, there is no significant correlation at any lag between the X-ray and optical bands. We also find no correlation at any lag between optical flux and various X-ray hardness ratios. We conclude that the close relation observed between the bands during the first year of our program was either a fluke, or perhaps the result of the exceptionally bright state of NGC 3516 in 1997, to which it has yet to return. Reviewing the results of published joint X-ray and UV/optical Seyfert monitoring programs, we speculate  that there are at least two components or mechanisms contributing to the X-ray continuum emission up to 10 keV: a soft component that is correlated with UV/optical variations on timescales $\\gtorder 1$ day, and whose presence can be detected when the source is observed at low enough energies ($\\sim 1 $keV), is unabsorbed, or is in a sufficiently bright phase; and a hard component whose variations are uncorrelated with the UV/optical. ",
        "introduction": "The paradigm that active galactic nuclei (AGNs) are powered by accretion onto massive black holes (MBHs) has recently gained strong observational support, with the detection, in several AGNs, of X-ray emission lines that are thought to be broadened by relativistic effects near the MBH horizon (Nandra et al. 1997; Sako et al. 2002), the evidence for dormant black holes in many normal nearby galaxies (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000), and the estimates of MBH masses in several tens of AGNs via reverberation mapping (Kaspi et al. 2000). However, the detailed mechanisms by which accretion produces the observed spectral energy distributions, as well as other properties, of AGNs are unknown, and observations have placed few constraints on the many theoretical scenarios proposed.  It has been hoped that flux variations in different energy bands would provide clues toward understanding the AGN emission processes. In particular, a number of bright Seyfert-1 galaxies have been subject to contemporaneous X-ray and UV/optical monitoring aimed at detecting inter-band lags, which could etablish a relation between emission components, e.g., by identifying the primary and secondary (i.e., reprocessed) emissions (Done et al. 1990;  Clavel et al. 1992; Kaspi et al. 1996; Crenshaw et al. 1996; Warwick et al.  1996; Edelson et al. 1996; Nandra et al. 1998; Edelson et al.  2000; Peterson et al. 2000; Pounds et al. 2001; Turner et al. 2001; Collier et al. 2001; Shemmer et al. 2001). However, the results of these programs, which have searched for correlations and lags on timescales of hours to weeks, have not been conclusive. It is generally true that UV/optical variation amplitudes are much smaller than those in the X-rays, which could argue that the X-rays are the primary emission. Clear lags between X-ray and UV/optical variations have not been seen. In those cases where correlation at a lag between different X-ray bands has been detected (sometimes with debatable significance), the lag increased with band energy, indicating the X-rays are secondary (e.g. Chiang et al. 2000). In a variant on the idea of searching for correlations between fluxes at different bands, Nandra et al. (2000) found that the X-ray spectral index in NGC~7469 was correlated with UV flux at zero lag during a month-long campign on this Seyfert 1 galaxy. Papadakis, Nandra, \\& Kazanas (2001) have analyzed the cross-spectrum of variations in several X-ray bands in this object, and found  that harder X-rays are delayed with respect to soft ones, with the delay proportional to the Fourier period probed. Such behavior is common in Galactic black hole binaries, but several competing theoretical explanations exist for it.  The studies mentioned above have tended to be of limited duration - often just a few days (Peterson et al. 2000 being the main exception). A potential pitfall of short duration studies is that they may detect few or no large-amplitude variation events with which to search for inter-band correlations. Furthermore, the results of variability studies may depend on the timescale sampled, and different behavior may pertain to different sources. In 1997 we initiated a long-term X-ray/optical program to monitor continuously several Seyfert 1 galaxies, such that month- and year-long variation timescales can be properly probed, as well as shorter timescales. X-ray observations are obtained with the {\\it Rossi X-ray Timing Explorer} (\\rxte), and optical data are from the Wise Observatory 1m telescope.  In Maoz, Edelson, \\& Nandra (2000, hereafter Paper I), we presented the first 1.5 year of X-ray and optical data for NGC~3516. Paper I found that the low-frequency component of the X-ray variations appeared to mimic the optical variation during the first year, but with a lag of $\\sim 100$ days. However, this correlation ceased in the last 6 months of the data. Paper I found that the correlation was significant at the $\\sim$99\\% level, based on Monte Carlo simulations that assumed power-law fluctuation power spectra with slope --1.0 in the X-rays and --1.75 in the optical.  This was reasonable based on the best information available at the time (Edelson \\& Nandra 1999). More recent data suggest a steeper X-ray power spectrum slope, of --1.35  (Markowitz \\& Edelson 2002).  As discussed in Paper I, steeper slopes will yield lower significance levels, and revised simulations will be reported in a future paper (Edelson, Uttley, \\& Markowitz 2002). Although we proposed some physical explanations for the correlation, we cautioned that it was driven by a single variation ``event'' and could therefore be a statistical coincidence. Here, we revisit NGC~3516 after having accumulated 5 years of contemporaneous X-ray and optical data. ",
        "conclusions": "Much current thinking about the emission processes in AGNs centers around the notion that the X-rays arise from very close (within a few Schwarzschild radii) of a massive black hole. Support for this idea has come from the rapid variability that is observed in X-rays (implying small physical scales), as well as the detection in X-rays of a broad Fe K-shell emission line in many Seyfert 1s (e.g., Nandra et al. 1997). The emission line is thought to be gravitationally and Doppler broadened fluorescence of the inner parts of an accretion disk, after the disk is illuminated by the X-rays. More recently, such relativistic emission lines from the Ly$\\alpha$ transitions of several hydrogen-like ions may have been detected in XMM-{\\it Newton} data for two Seyfert galaxies (Branduardi-Raymont et al. 2001; Sako et al. 2002), though this claim has been contested using {\\it Chandra} data (Lee et al. 2001). The continuum-emission mechanism is not known, but most commonly it is assumed that the X-rays are optical/UV photons which have been upscattered by a population of hot electrons (e.g., Sunyaev \\& Truemper 1979). The acceleration mechanism and geometry of the X-ray source is not known. Neither is the source of seed photons, and despite some substantial problems it is still usually assumed that the optical/UV arises directly from an accretion disk (Shields 1978; Malkan 1983). It has also been hypothesized that X--rays illuminating the disk, or other optically thick gas, might be responsible for some or all of the optical/UV radiation, via reprocessing (Guilbert \\& Rees 1988; Clavel et al. 1992).  Variability data such as those we have presented above can provide constraints on possible models. In summary of the observational results, we have found a similarity between the optical and X-ray light curves during the first year of our program, when optically the source was particularly bright, and with the X-rays lagging  the optical variations by about 100 days. This correlation disappeared in the last 4 years of the data, during which we see no clear correspondence at any lag between the optical and the X-rays. Furthermore we do not find any clear trends when we examine X-ray softness ratios, rather than fluxes. The only positive signal we find are a rough trend for a softer spectrum in the 4-10 keV range when the source is optically brighter. (The relation between X-ray spectral slope and brightness in {\\it X-rays} will be examined in a separate paper on the X-ray properties of this object.)  Phenomenologically, the reality of any of these trends is debatable, and all of them may be chance coincidences. A more stringent test must await the results of continued monitoring, during which NGC~3516 may perhaps recover to the high optical brightness it attained between mid-1997 and mid-1998. The lack of any straightforward correlation between X-ray and optical fluxes, in its simplest interpretation, argues that there is no physical relation between the emission in the two bands, except perhaps that both ultimately derive their energy from the central black hole. If there is a connection between the emission mechanisms in these two wavelength regimes, at the very least it must be complex enough to wash out any evidence for it in the variability data.  Is NGC~3516 peculiar among AGNs in its lack of a clear correlation between X-ray and optical/UV fluxes? To address this, we critically review the results of previous campaigns on this and other Seyfert galaxies.\\\\ {\\bf NGC~4051} Done et al. (1990) monitored NGC~4051 for 2 days, and found no correspondence between the large-alplitude 2-10 keV variations seen with {\\it Ginga} and the constant (to $< 1\\%$) optical flux. Peterson et al. (2000) monitored this galaxy for 3 years with \\rxte~ at 2-10 keV and with ground-based optical spectroscopy. Typical sampling intervals were 1-2 weeks in both wavelength regimes. In the third year, the source went into an extremely low X-ray state. While confirming the lack of correlation found by Done et al. (1990) on short timescales, Peterson et al. (2000) found that the light curves are correlated at near-zero lag after smoothing on 30-day timescales. \\\\ {\\bf NGC~5548} Clavel et al. (1992) observed NGC~5548 simultaneously with {\\it Ginga} at 2-10 keV and with {\\it IUE} at 1350 \\AA\\ over a period spanning 51 days. The source brightness was lower than average both in UV and in X-rays. The authors claimed a significant zero-lag correlation, yet this was based on nine epochs, and basically one-half of an ``event'' in the light curves. Chiang et al. (2000) observed NGC~5548 for 2.8 days simultaneously with {\\it EUVE} (0.14-0.18 keV), {\\it ASCA} (0.5-1 keV), and \\rxte~ (2-20 keV), with 44 {\\it EUVE} epochs. They found a good correlation between the three bands, but as in the previous experiment on this object by Clavel et al. (1992), the correlation is dominated by a single ``step'' in the light curves. The connection of the extreme-UV with the UV range was previously given by Marshall et al. (1997) who compared {\\it EUVE} measurements to {\\it IUE} and {\\it HST} UV observations, but the correlation they claimed was based on 10 data points spaced over 10 days, and a low correlation coefficient.\\\\ {\\bf NGC~4151} Edelson et al. (1996) combined 14 epochs of {\\it Rosat} 1-2 keV data and four epochs of {\\it ASCA} (0.5-1 keV) data (Warwick et al.  1996), and compared them to {\\it IUE} ultraviolet (Crenshaw et al. 1996), and Wise Observatory optical (Kaspi et al. 1996) measurements of NGC~4151 which were comtemporaneous over 10 days. The source was near its peak historical brightness. In this case, the light curves at all bands showed zero-lag similarities on $\\sim 1$-day timescales. However, the X-ray light curves had an overall rising trend during the 10-day period, whereas a constant or falling trend was seen in the UV and optical light curves. Thus, the X-ray to UV/optical correspondence was far from perfect, and in some sense opposed.\\\\ {\\bf NGC~7469} Nandra et al. (1998) observed NGC~7469 for over a period of 30 days with 30 epochs (after averaging) and found that the \\rxte~ 2-10 keV and {\\it IUE} UV fluxes were poorly correlated. Nandra et al. (2000) then found in these data a better correlation of the UV flux with the X-ray slope, rather than X-ray flux. The object was close to its average brightness in X-rays and in UV.\\\\ {\\bf Akn 564} This narrow-line Seyfert 1 was monitored approximately daily for 50 days in the optical (Shemmer et al. 2001), in the UV with {\\it HST} (Collier et al. 2001), and in the X-rays with {\\it ASCA} at 0.7-1.3 keV (Turner et al. 2001) and with \\rxte~ at 2-10 keV (Pounds et al. 2001). Although variation amplitudes in the UV and optical were only of order a few percent, a  correlation at $<1$~day lag between UV and X-raya was reported by Shemmer et al. (2001), as well as a possible correlation between X-ray and optical, if only a particular segment of the optical light curve, surrounding a relatively large event, is used in the analysis. \\\\ {\\bf NGC~3516} Edelson et al. (2000) monitored NGC~3516 continuously for 3 days with \\rxte~ and {\\it ASCA} at 2-10 keV, and with {\\it HST} in the optical. They found no significant correlation between X-ray and optical variations. Those observations took place on days 917-920 (see Fig. 4) when the source optical brightness was average but the X-ray flux was relatively high. In the present work on this AGN, we find no correlation between optical and 2-10 keV variations on timescales of days to 5 years, except possibly a 100-day delayed correlation during 1997-1998, when the source was extremely bright. We do not see a correlation with X-ray slope analogous to that found in NGC~7469 by Nandra et al. (2000).  If we now attempt to synthesize the above results, the following picture emerges. There have been several cases of little or no correspondence between X-ray and UV/optical variability. There have also been several cases where a correlation has been claimed, but the result is not conclusive due to poor sampling, insufficient variability, or low significance. Perhaps the most convincing flux correlation has been seen by  Edelson et al. (1996) between soft (1-2 keV) X-rays and UV/optical flux in NGC~4151, yet, as mentioned above, the longterm trends in the two bands were opposed, and no rigorous simulations have been done to quantify the significance of the correlation. The correlation between UV flux and X-ray slope found in NGC~7469 by Nandra et al. (2000) also seems secure. These latter two results could arise if (but do not necessarily imply that) the UV/optical is better correlated with soft ($< 2$ keV) X-ray variations than with the hard X-rays. It is also important to note that, in terms of timescales, there has been no evidence in any Seyfert 1 of a relation between X-ray and UV/optical variations at short ($< 1$~day) timescales, and all claimed correlations have been on $\\ge 1$ day timescales. The fast variations therefore appear be associated mainly with the harder X-rays. This is supported by the finding by that the X-ray variation power density spectrum flattens with increasing energy  in NGC 7469 (Nandra \\& Papadakis 2001), Akn 564, and Ton S180 (Edelson et al. 2001).  Why, then, is a relation between optical flux and X-ray slope, such as seen in NGC~7469, not seen in NGC~3516 in the present work? It can be argued that, contrary to NGC~7469, NGC~3516 has strong and variable absorption in X-rays, and that this variable absorption decorrelates an intrinsic relation in NGC~3516 that is similar to the one in NGC~7469. Evidence for this can be seen in the fact that, in the present data for NGC~3516, the 4-7 keV and 7-10 keV light curves, as well as the full 2-10 keV light curve are more similar to each other than to the 2-4 keV light curve, where absorption will be strongest.  On the other hand, the X-ray absorption  in NGC~3516 is comparable to that in NGC~4151, where a UV-X-ray flux correlation is seen in a band that is only slightly softer than the \\rxte~ band. Perhaps this objection can be overcome by noting that the NGC~4151 correlation was seen when this source was exceptionally bright. If the source brightness is a factor, it can further be argued that a flux correlation was indeed seen in NGC~3516, but only during the first year of our program, when the source was exceptionally bright, as was the case in NGC~4151. Indeed, we note in NGC~3516 that the 2-4/4-7 keV and  4-7/7-10 keV ratios and the total X-ray counts all seem to track each other better during the first 700 days, when the source was brighter. Source brightness could conceivably affect the correlations by making visible the high-energy tail of the actually correlated emission at low energies, or by ionizing the absorbing gas, and thus reducing the decorrelating effect of the variable absorption. The clear change in X-ray spectral softness around day 1300 and the  accompanying optical dimming, while not necessarily connected, are at least  consistent with the expectation that a more photon-starved corona will produce a harder spectrum. Alternatively, rather than source brightness playing a role, it may be intrinsic differences between objects. For example, the fast, uncorrelated X-ray emission may be always dominant in NGC~3516, and hence swamp out the soft correlated emission in the light curves.   According to this picture, then, at least two components contribute to the X-ray continuum emission of Seyfert nuclei: a soft component which is temporally related to the UV/optical continuum, and which can be discerned in $> 1$~keV variability data only when the source is bright enough or relatively unabsorbed; and a fast/hard component that varies independently. We note that by ``components'' we do not necessarily mean emission from physically distinct regions (e.g., Shih, Iwasawa, \\& Fabian 2002).  Instead, the two components that contribute to the X-rays could arise from the same region via different mechanisms. For example, the physical conditions in the coronal regions could be affected autonomously by two processes, such as by time evolution (e.g. Poutanen and Fabian 1999) and by changes in seed photon input which cause an overall temperature change in all the coronal regions. One or the other of these mechanisms could dominate in a particular source at a particular time and on a particular timescale.   This empirical two-component picture and its underlying drivers could be tested via a monitoring program using X-ray observations with sufficient spectral resolution and signal-to-noise ratio to disentangle the continuum from the absorption, and measure the variability of the intrinsic flux and spectrum. Another path is to search for the the spectral-slope/optical-flux correlation in other objects, having either strong or weak X-ray absorption. We intend to do this in future papers for NGC~4151, NGC~5548, and PG0804+762.  Future programs can also test an alternative interpretation of all previous results, namely, that there is no real correlation of any sort between variations in optical/UV and X-ray bands. In that case, even the more convincing correlations seen, such as Edelson et al. (1996) and Nandra et al. (2000) are chance coincidences arising in the comparison of unrelated red-noise light curves. The simulations testing for this possibility can also be refined. For example, the simulations in Paper I assumed an X-ray power spectrum slope of $-1$, but  recent work indicates a slope of $-1.35$ or steeper over the time scales of interest (Markowitz \\& Edelson 2002), which would lead to a lowered significance level. The simulated light curves can include amplitude randomization (Timmer \\& Koenig 1995) as well as phase randomization. We do emphasize that all future claims of inter-band correlations would benefit by simulations demonstrating their significance."
    },
    "0207/astro-ph0207178_arXiv.txt": {
        "abstract": "We study the effect of imperfect subtraction of the Sunyaev-Zel'dovich effect (SZE) using a robust and non-parametric method to estimate the SZE residual in the Planck channels.  We include relativistic corrections to the SZE, and present a simple fitting formula for the SZE temperature dependence for the Planck channels. We show how the relativistic corrections constitute a serious problem for the estimation of the kinematic SZE component from Planck data, since the key channel to estimate the kinematic component of the SZE, at 217 GHz, will be contaminated by a non-negligible thermal SZE component. The imperfect subtraction of the SZE will have an effect on both the Planck cluster catalogue and the recovered CMB map. In the cluster catalogue, the relativistic corrections are not a major worry for the estimation of the total cluster flux of the thermal SZE component, however, they must be included in the SZE simulation when calculating the selection function and completeness level. The power spectrum of the residual at 353 GHz, where the intensity of the thermal SZE is maximum, does not contribute significantly to the power spectrum of the CMB. We calculate the non-Gaussian signal due to the SZE residual in the 353 GHz CMB map using a simple Gaussianity estimator, and this estimator detects a 4.25$\\sigma$ non-Gaussian signal at small scales, which could be mistaken for a primordial non-Gaussian signature. The other channels do not show any significant departure from Gaussianity with our estimator. ",
        "introduction": "The quality of the data from the Planck satellite will allow us to make precision cosmology, however, accurate parameter extraction will require precise modeling of the data.  The scientific possibilities with the new data will depend strongly on the ability to perform the component separation. In each of the Planck channels the data will be a mixture of galactic components (synchrotron, free-free and dust), extra-galactic components (unresolved galaxies and galaxy clusters), and instrumental noise.  Recently several algorithms have been proposed to perform such component separation. Most of these methods rely on {\\it a priori} knowledge of the frequency dependence of each individual component and their power spectrum (maximum entropy, \\cite{Hobson1998}; multi-frequency Wiener filter, \\cite{Tegmark1996,Bouchet1999}).  The advantage of these methods is that they can recover all the different components simultaneously, however, the drawback is that if some of the assumptions about the frequency dependence and/or power spectrum is wrong, then the final result will be biased.  Other methods are designed to recover just one of the components, and the most popular ones focus on the recovery of compact sources. Since the best resolution of Planck is 5 arcmin, the extra-galactic galaxies will appear as unresolved point sources with a shape matching the point spread function of the instrument. This fact can be used to define optimal filters which will increase the signal to noise ratio of the bright point sources, thus allowing the detection and removal of most of them \\cite{Tegmark1998,Sanz2001}. A similar technique can be applied to the detection of clusters if one assumes a circular shape. Since the typical diameter of a cluster is a few arcmin, most of the clusters will appear similar to the unresolved point sources, and only some will be resolved by the instrument.  The definition of the optimal filter is a bit more complicated in this case since the optimal scale of the filter will be different for each cluster, however, this problem can be partially solved by filtering the maps with different scales \\cite{Herranz2001}.  An alternative non-parametric approach to detect clusters in Planck data was recently proposed \\cite{Diego2002}.  In that method no assumptions are made about the specific frequency dependence of the different components (except for the SZE component), and also no assumption about the power spectrum of the components or scale (or symmetry) of the clusters. The only assumption is that the frequency dependence of the SZE (in the non-relativistic approximation) is known.  Even without many of the typical assumptions, the recovered SZE component is a good estimate of the total contribution of galaxy clusters to the different Planck channels, however, the SZE component recovered by this method (as well as by other methods) does not match the real SZE component perfectly in the simulations. Therefore, a residual will be left in the final CMB map due to the non-perfect subtraction of the SZE from the original data.  So far there have been no attempts to study how this SZE residual could affect the conclusions derived from the {\\it residual contaminated} CMB map.  One of the reasons is that the SZE residual depends on the method used to make the component separation, and different methods recover the SZE component with different {\\it quality factors} and residuals. Consequently, the residual SZE map (defined as the ``true'' minus the ``recovered'') is different for each method. In this work we will study the SZE residual from the non-parametric method proposed in \\cite{Diego2002}.  Since this method makes a minimum number of assumptions, the results are robust and less affected by systematic errors which could be introduced through wrong assumptions.  The conclusions of this work will therefore set an upper limit on the contributions of the SZE residual to the CMB. Any other method which makes more assumptions than our non-parametric method, should leave a smaller SZE residual (provided the assumptions are correct), and as a consequence the effect of the SZE residual for any other (suitable) method should be below the limits we present here. It is worth pointing out that since the non-parametric method adopted here is optimized for the detection of the SZE signal, then alternative methods making many more assumptions do not necessarily obtain a better reconstruction of the SZE signal. Therefore, although our approach provides an upper limit on the effects of the SZE residual, then that limit can be taken as a realistic estimate of the final contribution of the SZE residual to the CMB map.   In this work we will consider an additional source of systematic error which has not been considered previously, namely the relativistic corrections to the SZE.  Although these relativistic corrections are small for normal clusters ($ T \\approx $ few keV), they can be important for massive clusters ($ T \\approx 10$ keV). The cluster selection function of Planck (minimum mass detected as a function of redshift) rises very quickly from redshift 0 to redshift $\\approx$ 0.2, and above that redshift it is almost flat (see Fig. \\ref{fig_Selection_Function} below).  This means that above redshift $\\approx$ 0.2 Planck will only see massive clusters for which the relativistic corrections are important. In all component separation methods (including the non-parametric method considered in this work), the frequency dependence of the SZE component is assumed to follow the non-relativistic form, eq.~(\\ref{freqdep}). The validity of this approach is unknown, and it is therefore worth while exploring the effect of the relativistic corrections in the context of the component separation process.  The main effect of the relativistic corrections can be described effectively as a {\\it dilution} of the frequency dependence of the SZE.  At low frequencies ($\\nu < 217$ GHz), the relativistic correction lowers the absolute value of the SZE intensity change. The same thing happens at higher frequencies up to $\\nu \\approx 400$ GHz. Above this frequency, the intensity change is larger than the one given by the non-relativistic approach (see Fig. \\ref{fig_fx_Relat_NoRelat}), however, the dust contamination is very important at such large frequencies, and hence the effect of the relativistic corrections becomes negligible for $\\nu> 400$ GHz. \\begin{figure} \\begin{center} \\epsfxsize=8.cm \\begin{minipage}{\\epsfxsize}\\epsffile{fig1.eps} \\end{minipage} \\caption{\\label{fig_fx_Relat_NoRelat} Solid line is proportional to the non-relativistic intensity change which is assumed in all existing component separation algorithms. The dashed (dotted) line shows the frequency dependence of the SZE when the relativistic corrections are included for a cluster with $T = 10$ keV ($T = 20$ keV).} \\end{center} \\end{figure}  The non-relativistic form is systematically assumed in all component separation algorithms, however, the dilution effect due to the inclusion of the relativistic correction makes the form of the SZE for the hottest clusters different from the non-relativistic form.  This will lead to an additional error in the residual of the SZE signal. Our non-parametric method makes the (wrong) assumption that the relativistic corrections are negligible for all the clusters in our simulation.  However, we will test the method with SZE simulations where the relativistic corrections are incorporated. By doing that, we are simulating the realistic case in which the temperature of the cluster is not known, and therefore the non-relativistic form must be assumed in the algorithm to recover the SZE component. This wrong assumption in our method (and in all other methods) will add an additional error in the recovered SZE map.   The cluster abundance is a sensitive probe of cosmological parameters such as the matter density, $\\Omega_m$, and when the observations reach low statistical error it becomes important to control the systematic errors in the parameter extraction. The relativistic corrections to the SZE could add an additional error in the estimate of the cluster number counts, and we will below consider the importance hereof.   To study the effect of the SZE residual on the CMB we will focus on two aspects of the CMB, namely its power spectrum and its Gaussian nature.  The power spectrum depends strongly on the cosmological parameters, and a systematic error in the estimation of the power spectrum due to non-subtracted residuals could have important consequences for the best fitting cosmological model.  Gaussianity is a natural prediction of single field inflationary models, and a non-Gaussian signature in the CMB could have important consequences for such inflationary models.  The SZE signal is very non-Gaussian and so is the non-subtracted residual.  In previous works (see \\cite{Aghanim1999,Cooray:2001wa,rephaeli2001,Yoshida:2001vf} and references therein), the non-Gaussian signature of the SZE has been studied, but these works focus on the entire contribution of this component. Since a large part of the SZE signal will be removed in the component separation process, the non-Gaussian signature, if any, will be smaller than predicted in previous works, however, it could still be significant.  It is important to understand how this residual could leave a non-Gaussian imprint in the CMB map, such that an erroneous interpretation as of primordial nature of a non-Gaussian signature can be avoided.  In this work we will study the implications on Gaussianity studies of the SZE residual left after component separation.  The structure of the paper is the following. In section \\ref{Section_SZE} we will give a brief description of the SZE and the relativistic corrections. We also present a simple fitting formula to the temperature dependence of the SZE in the central frequencies of the Planck channels.  This fitting formula could be used in future works to include the relativistic corrections in the simulations.  In section \\ref{section_residual} we apply the non-parametric method to realistic Planck simulations and we recover the SZE component in two cases.  In the first case we simulate a population of clusters all with the same temperature.  This allows us to assume the real frequency dependence of the SZE in the non-parametric component separation method and see what the difference is with the case when the non-relativistic form is assumed in the component separation.  In the second test, we make realistic simulations of the clusters with the population of clusters having different temperatures. We include the relativistic correction in our simulation. Then we calculate the SZE residual assuming the non-relativistic approach for the frequency dependence of the SZE.  In section \\ref{section_EffectSZE} we discuss the effects of the imperfect SZE recovery on the Planck cluster catalogue focusing our attention on the systematic errors introduced in the recovered SZE map by the relativistic corrections and the problems this causes for the kinematic SZ component.  In section \\ref{section_EffectCMB} we discuss the effect of the SZE residual on the CMB map with emphasis on non-Gaussian signatures.  Finally we present our conclusions in section \\ref{section_Conclusions}. ",
        "conclusions": "In this paper we have studied the systematic errors introduced in the Planck cluster catalogue and in the CMB map due to the imperfect SZE subtraction caused by the non-relativistic assumption and to the intrinsic error in the component separation process.  We have used the non-parametric method proposed in \\cite{Diego2002} to perform the component separation method. This method is very robust in the sense that it makes a minimum number of assumptions. The drawback is that it is not the most precise in the determination of the SZE component. However, this allow us to put an upper limit on the systematic errors introduced by the imperfect SZE subtraction since any more sophisticated method in principle should produce smaller SZE residuals.   We have seen that the effect of the non-relativistic assumption in the frequency dependence of the SZE made in the component separation process is small when compared to the case where the real frequency dependence is assumed (between 4\\% and 8\\% relative difference for $T = 10$ keV clusters), however, the relativistic corrections should be considered in the simulations of the SZE in order to compute correctly the selection function and completeness level of the cluster catalogue.  Concerning the kinematic component of the SZE, relativistic corrections should be taken into account in order to recover this component. Otherwise errors as large as $50-100 \\%$ could be introduced in this component in the most relevant channel for its detection (217 GHz).  In the CMB, the SZE residual in the 353 GHz leaves a non-Gaussian signature at small scales which could be detected by some Gaussianity estimators like the MHW. This channel could thus be used to extract non-Gaussianity signatures from imperfect SZE subtraction from the other channels.  Our MHW Gaussianity estimator does not show significant deviations from Gaussianity due to imperfect SZE recovery in the other channels relevant to the CMB.  The SZE residual does not change significantly the power spectrum of the CMB in the 353 GHz channel. It is therefore safe to include the 353 GHz channel for the computation of the CMB power spectrum. This channel, together with the 217 GHz channel, are the most important ones which will contribute to the CMB map at the smallest scales ($\\approx 5$ arcmin)."
    },
    "0207/astro-ph0207452_arXiv.txt": {
        "abstract": "We re-analyse spectropolarimetric data of AB Dor taken in 1996 December using a surface imaging code that can model the magnetic field of the star as a non-potential current-carrying magnetic field. We find that a non-potential field needs to be introduced in order to fit the dataset at this epoch. This non-potential component takes the form of a strong unidirectional azimuthal field of a similar strength to the radial field. This azimuthal field is concentrated around the boundary of the dark polar spot recovered at the surface of the star using Doppler imaging. As polarization signatures from the center of starspots are suppressed, it is unclear whether or not this non-potential component genuinely represents electric current at the unspotted surface or whether it results from the preferred detection of horizontal field in starspot penumbrae. This model contains 20\\% more energy than the corresponding potential field model at the surface. This amount of free energy drops to under 1\\% about 1{\\,\\mbox{$\\mbox{R}_*$}}\\ above the photosphere.  We use these surface maps to model the coronal structure of the star. The mixed radial polarities at the pole in the surface maps support closed coronal loops in the high latitude regions, indicating that a component of the X-ray emission may originate in this area. Assuming that the field remains closed out to 5{\\,\\mbox{$\\mbox{R}_*$}}\\, we find stable surfaces where prominences may form out to the observed distances using this coronal model. ",
        "introduction": "Doppler imaging allows us to map surface brightness distributions across the surfaces of rapidly rotating stars (Vogt \\& Penrod 1983). This technique has proven to be an important tool when studying activity in rapidly rotating solar-type stars. Spot maps show that, while {\\em sunspots} tend to congregate between $\\pm 30${\\mbox{$^\\circ$}}\\ latitude, spot patterns on other cool stars are different. In rapid rotators, strong spot activity is often  concentrated in polar or high-latitude regions (Strassmeier 1996). This difference to the solar pattern has been the subject of much debate in recent years. The presence of strong flux at high latitudes  may be explained in terms of  increased Coriolis forces acting on flux tubes in the convective interior in rapidly rotating stars. These would lead to the emergence of flux tubes in high latitude regions ({Sch\\\"ussler} \\& Solanki 1992). An alternative  explanation for the polar spot phenomenon is that there are rings of alternating polarity at the pole (Schrijver \\& Title 2001). These are produced when increased flux is injected into solar-type models. Meridional flows transport this flux to the pole. As the flux injection rate is increased, flux cancellation occurs more slowly and results in strong flux gathering at the stellar pole.   In general, these dark starspots  are assumed to mark the regions where the strongest magnetic flux is concentrated. However, it is likely that smaller spots exist below our resolution limit that we cannot reconstruct. Semel (1989) proposed applying standard  Doppler imaging principles to circularly polarized spectra in order to detect  magnetic fields on the surfaces of rapidly rotating stars. This technique is called Zeeman Doppler imaging  (ZDI). Brown et al. (1991) first developed a code that employed maximum entropy techniques to enable the mapping of magnetic fields on the stellar photospheres. This method has since been applied to three rapidly rotating systems: the RS CVn binary, HR1099 (K1IV+G5V);  the K0 dwarfs; AB Dor and LQ Hya (Donati \\& Collier Cameron 1997, Donati 1999). The technique is not very sensitive to magnetic fields in spotted regions, but the resulting maps typically show the presence of strong magnetic field even in relatively unspotted parts of the photosphere. A puzzling feature in these maps is that the radial and azimuthal field components across the surface are of about the same strength. This is very different to the solar case, where strong horizontal field is only found in sunspot penumbrae.  A criticism of ZDI has been that no relationship is assumed between the three components of the field vector thus enabling the reconstruction of physically unrealistic magnetic field distribution patterns such as  monopoles at the stellar surface. Potential field configurations and their extrapolations have been used to model the global surface and coronal field of the Sun  (Altschuler \\& Newkirk 1969). We show how the surface and coronal field of AB Dor can be modeled using the magnetic field maps obtained using an advanced version of ZDI. Hussain, Jardine \\& Collier Cameron (2001) describe a technique that introduces a relationship between the different components of the surface magnetic field vector by assuming that the observed data can be fitted using a potential field distribution. In this paper we extend this method to non-potential fields, i.e. we assume that the surface field can be modeled using a potential field plus a non-potential field component that represents the presence of electric currents. This method is similar to that described by  Donati (2001).  The K0 dwarf, AB Doradus (HD 36705), is a relatively bright ($m_{V}\\approx 7$) example of a class of very active cool stars that are just evolving onto the main sequence. Its distance has been measured to be 15pc using HIPPARCOS and VLBI interferometry and its age is estimated to be between 20 to 30~Myr (Guirado et al. 1997, Collier Cameron \\& Foing 1997). It rotates at roughly 50 times the solar rate, $P_{\\mbox {rot}}=0.51479$~d (Pakull 1981). Its surface is covered with large cool starspots as indicated by rotational modulation of its photometric light curve (Rucinski 1985). AB Dor also displays strong X-ray variability on all timescales ({K\\\"urster} et al. 1997). Long-term photometry indicates that  the star was at its brightest in the first epoch of observation in 1978, it then decreased down to a minimum level in 1988 and has since been increasing steadily, most recently plateauing out to roughly the same $m_V$ level as 1978. This may be evidence of a solar-type activity cycle with the starspot coverage increasing and decreasing with a period of 22-23 years (Amado 2001).  Doppler images of this star have revealed the presence of a polar spot that has been consistently present since 1992 (Collier Cameron \\& Unruh 1994, Collier Cameron 1995, Donati \\& Collier Cameron 1997, Donati et al. 1999). The only image of AB Dor that does not display this polar feature was obtained using data taken in 1989 ({K\\\"urster}, Schmitt \\& Cutispoto 1994), the star was at its minimum brightness level at this epoch and hence at its most spotted (Amado 2001). Hence this feature may be associated in some way with a stellar activity cycle. This polar spot should be associated with strong magnetic field if the solar stellar analogy applies. Lower-latitude spots are also recovered in Doppler maps but are found to be much less stable.  ZDI studies of AB Dor carried out annually since 1995 reveal magnetic field maps with strong field in even relatively unspotted parts of the photosphere. The strong unidirectional azimuthal  flux recovered near the pole in maps obtained in successive years provides a puzzling non-solar pattern (Donati \\& Collier Cameron 1997, Donati et al. 1999, Hussain 2000). We re-analyse the best sampled dataset using a code that introduces a relationship between all three field components, and use this analysis to evaluate if this band of azimuthal flux is still present. These new surface maps are extrapolated to model the structure of the corona. This coronal topology is used to evaluate general properties of the stellar corona. The model is also used to compute sheet-like surfaces of stable mechanical equilibria taking the rotation of the star into account (Jardine et al. 2001). These are sites where the prominences are likely to form. We use these coronal magnetic field models to model the temperatures and densities expected in the corona of AB Dor in paper II (van Ballegooijen \\& Hussain 2002). ",
        "conclusions": "The magnetic field analysis of the stellar surface indicates that there is a strong non-potential field encircling the polar spot on AB Dor in 1996 December. This azimuthal field in located in the penumbra of the polar spot and its detection may be affected by the reduced visibility of the magnetic fields in the spotted regions. However, the observed azimuthal field appears to be real.  Schrijver \\& Title (2001) model the effect that the increased magnetic flux injected onto the surfaces of solar-type stars would have on subsequent flux patterns. They find that  concentric bands of radial field with opposite polarity form near the pole of the star. This is a consequence of meridional transport of flux and the increased amount of flux dissipating on longer timescales. While our maps do not show this pattern, if there is strong flux of opposite polarities in the polar spot, meridional fields would connect the alternating bands of radial field. As the star rotates differentially, this horizontal field would be sheared in the azimuthal direction. Therefore, this model can in principle explain the origin of the azimuthal field.  Alternatively, there is preliminary evidence associating this azimuthal feature with the epoch of observation. First results suggest that the 1995 December images show no strong non-potential component and that 1998 January images have an even stronger non-potential component than required by the 1996 December dataset. Could this non-potential component be related to the stellar activity cycle? Does it switch direction over the course of the cycle? These are questions that analysis of subsequent datasets will answer. If this non-potential component is indeed real it is possible to evaluate the amount of magnetic energy available to power flares and coronal mass ejections. The amount of free energy integrated over the entire coronal volume is 14\\% of the potential-field energy. Fig.~\\ref{fig:freeen} shows how the ratio of non-potential and potential magnetic energy densities depends on height. The amount of extra energy drops off with height from 20\\% at the surface to well under 1\\% at the source surface height (5{\\,\\mbox{$\\mbox{R}_*$}}).  Analysis of subsequent datasets will reveal how typical the features recovered in this paper are and if they evolve from year to year. The method presented here promises to be an important tool in probing the dynamo activity of rapidly rotating cool stars.   \\begin{figure} \\epsscale{1} \\plotone{f7.eps} \\caption{The ratio of non-potential to potential magnetic field energy as a function of height. As this figure shows, the amount of free energy is greatest at the surface and drops quickly, levelling out  to under 1\\% about 1{\\,\\mbox{$\\mbox{R}_*$}}\\ above the photosphere.} \\label{fig:freeen} \\end{figure}  When modeling the coronal structure of the star it is necessary to define the point beyond which the  field is open and radial. In the model presented in Fig.~\\ref{fig:prom}, this value has been set to 5{\\,\\mbox{$\\mbox{R}_*$}}\\ as suggested by the presence of prominences at these distances. This model would apply if these prominences fit the model of slingshot-type prominences. Observations suggest that prominences form out to 5{\\,\\mbox{$\\mbox{R}_*$}}, are stable over one rotation cycle but tend to have been ejected over a period of two days. This instability may reflect the inherent variability in the coronal field of AB Dor. Surface features recovered on AB Dor are, by contrast, stable on timescales of over a week.  If the source surface is at 1.7{\\,\\mbox{$\\mbox{R}_*$}}\\ as suggested by an analysis of the gas and magnetic pressures in  AB Dor's corona (see Paper II), the prominences are located beyond the source surface and their equilibrium and stability cannot be described in terms of potential field models. The magnetic field beyond the source surface is nearly radial with current sheets separating regions with opposite sign of $B_r$. According to this new model the prominences would be associated with these current sheets, but unlike in Fig.~\\ref{fig:prom} they would not be in a state if magnetostatic equilibrium. Instead, the prominence plasma would slowly flow outward along the current sheets with velocity significantly less than the rotational velocity. This new model may explain the observation that prominences occur over a  broad range of longitudes."
    },
    "0207/astro-ph0207387_arXiv.txt": {
        "abstract": "We investigate the radial velocity difference between the \\oiiiopt\\ and \\hb\\ lines for a sample of $\\approx$ 200 low redshift AGN. We identify seven objects showing an \\oiii\\ blueshift relative to H$\\beta$\\ with amplitude larger than 250 \\kms\\ (blue ``outliers''). These line shifts are found in sources where the broad high ionization lines (e.g. \\civ) also show a large systematic blueshift. Such blueshifts occur only in the population A region of the Eigenvector 1 parameter domain (that also contains NLSy1 sources). We suggest that \\oiiiopt\\ blueshifts are also associated with the high ionization outflow originating in these sources. This is a direct kinematic linkage between narrow and broad line emitting gas. ",
        "introduction": "Forbidden \\oiiiopt\\ emission  arises in the narrow line region (NLR) of Active Galactic Nuclei (AGN). This emission has now been partly resolved in the nearest AGN, where the geometry of the line emitting gas has been found to be far from spherically symmetric. This suggests that measures of integrated \\oiiiopt\\ emission may correlate with source orientation to the line of sight (Hes et al. 1993; Sulentic, Marziani \\& Dultzin-Hacyan 2000a, and references therein). Observations and theoretical models both (e.g., Steffen et al. 1997; Sulentic \\& Marziani 1999; Moiseev et al. 2001)  suggest a complex interplay between (where applicable) shocks driven by radio ejection and the ionizing continuum from the nucleus.  At the same time it is generally believed that radial velocity measures of the narrow emission lines (e.g. narrow H$\\beta$ and \\oiiiopt) provide a reliable measure of the systemic, or rest-frame, velocity. \\oiii\\ is preferred because it is not superimposed on a much stronger broad line component. Limited HI, CO and absorption line measures of the host galaxy rest frame suggest that \\oiiiopt\\ usually gives consistent results within 200 \\kms\\ (de Robertis 1985; Whittle 1985; Wilson \\& Heckman 1985; Condon et al. 1985; Stirpe 1990; Alloin et al. 1992; Evans et al. 2001). Several observations, however, indicate that the Narrow Line Seyfert 1 (NLSy1) prototype I Zw 1 shows an \\oiiiopt\\ blueshift of $\\Delta v_r = -$500 \\kms\\ relative to other rest frame measures (Boroson \\& Oke 1987).  This corresponds to a ~10 \\AA\\ shift which is larger than any conceivable measurement or calibration errors; see Marziani et al. 1996).  We report here a study of the velocity shift of \\oiii\\ relative to \\hb. Our aim was to identify objects with large radial velocity disagreement between \\oiii\\ and \\hb\\ and their relationship with the general population of AGN. ",
        "conclusions": "We find that the  \\hb\\ and \\oiii\\  lines provide measures of the radial velocity of AGN usually consistent within $\\pm$200 \\kms. We identified infrequent ($\\approx$ 5\\% in a sample of 200 sources)  AGN showing \\dvr $\\la$ -250 \\kms. They belong to the  extreme population A sources that also show a large broad line \\civ\\ blueshift. This kinematic coupling of the NLR and BLR HIL emission most likely involves a wind or outflow. Our analysis suggests that blue outliers are not peculiar objects, but rather AGN viewed of extreme L/M with a compact NLR. % A predominance of blueshifts in the sample indicates that \\oiii\\ peak velocity is affected by outflow motions occurring in the innermost NLR."
    },
    "0207/astro-ph0207664_arXiv.txt": {
        "abstract": "\\vspace{-0.3cm}\\\\ We present the results of a large library of cosmological $N$-body simulations, using power-law initial spectra. The nonlinear evolution of the matter power spectra is compared with the predictions of existing analytic scaling formulae based on the work of Hamilton et al.  The scaling approach has assumed that highly nonlinear structures obey `stable clustering' and are frozen in proper coordinates.  Our results show that, when transformed under the self-similarity scaling, the scale-free spectra define a nonlinear locus that is clearly shallower than would be required under stable clustering. Furthermore, the small-scale nonlinear power increases as both the power spectrum index $n$ and the density parameter $\\Omega$ decrease, and this evolution is not well accounted for by the previous scaling formulae.  This breakdown of stable clustering can be understood as resulting from the modification of dark-matter haloes by continuing mergers. These effects are naturally included in the analytic `halo model' for nonlinear structure; we use this approach to fit both our scale-free results and also our previous CDM data. This method is more accurate than the commonly-used Peacock--Dodds formula and should be applicable to more general power spectra.  Code to evaluate nonlinear power spectra using this method is available from {\\tt http://as1.chem.nottingham.ac.uk/$\\sim$res/software.html}.  Following publication, we will make the power-law simulation data publically available through the Virgo website {\\tt http://www.mpa-garching.mpg.de/Virgo/}. ",
        "introduction": "In the current cosmological paradigm, structures grow through the gravitational instability of collisionless dark matter fluctuations. This occurs in a hierarchical way, with small-scale perturbations collapsing first and large-scale perturbations later. One of the most direct manifestations of this nonlinear process is the evolution of the power spectrum of the mass, $P(k)$, where $k$ is the wavenumber of a given Fourier mode.  Understanding this evolution of the power spectrum is one of the key problems in structure formation, being directly related to the abundance and clustering of galaxy systems as a function of mass and redshift. If the processes that contribute to the evolution can be captured in an accurate analytic model, this opens the way to using observations of the nonlinear mass distribution (from large-scale galaxy clustering or weak gravitational lensing) in order to recover the primordial spectrum of fluctuations.  One of the most influential attempts at such an analytic description of clustering evolution was the `scaling ansatz' of Hamilton et al. (1991; HKLM), which is described in Section \\ref{scfmodels}.  This scaling procedure was generalized to models with $\\Omega\\ne1$ and given a more accurate $N$-body calibration by Peacock \\& Dodds (1996; PD96).  HKLM assumed that a nonlinear collapsed object would decouple from the global expansion of the Universe to form an isolated system in virial equilibrium -- the `stable clustering' hypothesis of Davis \\& Peebles (1977). This assumption has been widely adopted, and yet it appears somewhat inconsistent with hierarchical models -- in which objects are continuously accreting mass and growing through mergers. Indeed, the validity of stable clustering has been increasingly questioned in recent years (e.g. Yano \\& Gouda 2000; Caldwell et al. 2001).  One of our aims in this paper is thus to establish whether stable clustering is relevant for understanding the small-scale evolution of the power spectrum.  We therefore explore the gravitational instability of dark matter fluctuations through a series of large $N$-body simulations of clustering from power-law initial conditions, with \\begin{equation} P(k) \\propto k^n. \\end{equation} We consider both $\\Omega=1$ models, in which the evolution can obey a similarity solution, and also low-density models with and without a cosmological constant.  We demonstrate that the resolution of the simulations is sufficient to measure the power well into the regime at which the HKLM procedure predicts a well-defined slope for the power spectrum determined by stable clustering.  In practice, we find that the power spectra are generally shallower than would be required for clustering to be stable on small scales.  Furthermore, as both $n$ and $\\Omega$ decrease, the amplitude of the small-scale spectrum increases in a manner that is not well described by any of the previous fitting formulae.  In light of these results, a new method for predicting nonlinear spectra is proposed.  This method is based on the `halo model' (e.g. Seljak 2000; Peacock \\& Smith 2000), which does not assume stable clustering. This allows us to fit our data and also the cold dark matter (CDM) data of Jenkins et al. (1998; J98) with a high degree of accuracy.  The paper is structured as follows. In Section \\ref{scfmodels} we provide a brief overview of the theoretical understanding of nonlinear evolution. In particular, a description of the stable clustering hypothesis, the nonlinear HKLM scaling relations and the halo model are given, as these ideas are central to this paper. We also discuss the scale-free models and their self-similarity properties.  In Section \\ref{scfsimulations} we describe the numerical simulations and we provide a visual comparison of the growth of structure in the different scale-free models. In Section \\ref{scfmeasuring} we describe an improved method for measuring power spectra and in Section \\ref{scfresults} we present the power spectra data and contrast them with the current nonlinear fitting formulae. In Section \\ref{scfhalo} we describe a new approach to fitting power spectra and its generalization to CDM, and then compare our new globally optimized formula with the results from Section \\ref{scfresults} and also the CDM data. Finally, in Section \\ref{scfdiscussion} we draw our conclusions and discuss our findings in a wider context.  \\vfill ",
        "conclusions": "\\label{scfdiscussion}  In this paper we have presented a set of high-resolution, $256^3$ particle, scale-free $N$-body simulations, designed to investigate self-similar gravitational clustering and in particular the effects of nonlinear evolution. We have also performed a further series of numerical simulations, with the same resolution, to explore how the evolution of clustering depends upon the background density of the universe. Together, these simulations represent the best calculations that exist to date for the set of models explored, with a factor 512 improvement in mass resolution over the ground-breaking work of Efstathiou et al. \\shortcite{Efstathiouetal1988}.  We verified that the final output power spectra were robust by considering grid and glass particle loads. However, at early times the problem of discreteness correction is simpler to handle if a glass start is applied; we have described a detailed method for correcting the clustering signal in this case.  We have implemented the power spectrum estimation technique of J98, which allowed us to probe high spatial frequencies without aliasing effects or errors due to mass assignment to the Fourier mesh.  The simulation results may be summarized as follows:  \\begin{enumerate}  \\item Scale-free simulations with $0<n<-2$ show self-similarity under the scaling $k_{0}(a)\\propto a^{-2/(n+3)}$. This conclusion is in agreement with the results of Efstathiou \\shortcite{Efstathiouetal1988} and Jain \\& Bertschinger \\shortcite{JainBertschinger1998}.  \\item In the quasi-linear regime, the power spectrum is characterized by a steep power law. The exact slope depends upon the spectral index $n$ of the input spectrum and the value of $\\Omega$, the slope steepening as $n$ becomes more negative and as $\\Omega$ is reduced.  \\item The observed nonlinear asymptote of the Einstein--de Sitter simulations was found to be inconsistent with the $\\Delta_{\\NL}^2\\propto [\\Delta_{\\L}^2]^{3/2}$ prediction of stable clustering.  A shallower slope with $\\Delta_{\\NL}^2\\propto [\\Delta_{\\L}^2]^{1}$ is preferred. This result makes sense in terms of the halo model: calculations using the extended Press-Schechter apparatus show that haloes will tend to merge with systems of similar mass to their own \\cite{LaceyCole1993}. Mergers of this kind will disrupt the virial equilibrium of the system, violating the basic assumption that underlies stable clustering. However, if this process were rare then stable clustering could be upheld in a statistical sense.  \\item The nonlinear fitting formulae of PD96 and JMW95 failed to reproduce the $n=-2$ results and were only marginally successful at reproducing the steeper spectra. The low-density power-law data were poorly fit by PD96.  \\item For the $\\Omega<1$ simulations, it is interesting to consider how the nonlinear slope changes with density. In the nonlinear limit equation (\\ref{app-halo1}) (appendix C) becomes \\be \\Delta^2(k)\\propto k^{\\;[3(f_1-1)+\\gamma_n]}\\ .\\ee For a given $n$, $f_1$ increases as $\\Omega$ decreases, and so the power-law slope steepens. This result supports the idea that small scale clustering is more closely related to the emergence of the internal density structure through the continual accretion and merger of haloes. The reasoning is as follows: for a low-density universe mergers are less frequent and so haloes have more time to virialize. This means that stable clustering may be considered to be a better approximation for these systems. From the arguments in Section 1 and 5, this would then be manifest as a steepening of the nonlinear slope.  \\end{enumerate}  In the second part of this paper, we proposed an improved fitting function for mass power spectra to replace the much-used PD96 formula. We have adopted a new approach to fitting power spectra, based upon a fusion of the halo model and a HKLM scaling.  The method was generalized to fit more realistic curved spectra, by introducing two new parameters, $n_{\\rm eff}$ the effective spectral index on the nonlinear scale, and the spectral curvature, $C$. We found that the halo model as previously envisaged in the literature fails to approach linear theory on large scales for $n\\ge 0$.  We have argued that this should be cured by changing the self-halo power from $n=0$ to $n=2$ on large enough scales, and we have shown empirically that this approach allows an accurate description of a very wide range of power spectrum data.  Our new fitting formula reproduced the scale-free power spectrum data and also the CDM results of J98 with an rms error better than 7\\%. This is to be preferred to the widely-used PD96 prescription, and should be useful for a variety of cosmological investigations.  In particular, our preliminary investigations show that the present formalism should cope naturally with spectra containing a realistic degree of baryonic features (e.g. Meiksin, White \\& Peacock 1999).  The halo model provides a novel way to view structure formation, and has yielded useful insights into the origin of nonlinear aspects of galaxy clustering.  This work has concentrated on the low-order statistics of the density field, but it is also possible to consider higher-order statistics such as the bispectrum.  This three-point function in Fourier space probes the shapes of large-scale structures that are generated by gravitational clustering. No shape information is included in the current formalism, so it will be interesting to see how well the model can account for higher-order statistics.  Initial results in this direction (Scoccimarro et al. 2001) seem to be promising.  In general, the important question is the extent to which the halo model can encapsulate the phase information in the density field, since fields with identical power spectra can possess completely different real-space distributions (e.g. Chiang \\& Coles 2000).  The halo model will inevitably fail to encompass these details of the density field in full, although it may still offer useful insights.  However, at the two-point level, we have shown that the model is far more than an educational device, and it can be used as a tool for a high-precision description of the evolution of the dark-matter power spectrum."
    },
    "0207/astro-ph0207044_arXiv.txt": {
        "abstract": "We simulate the formation of a realistic disk galaxy within the hierarchical scenario of structure formation and study its internal properties to the present epoch. We use a set of smoothed particle hydrodynamic (SPH) simulations with a high dynamical range and force resolution that include cooling, star formation, supernovae (SN) feedback and a redshift dependent UV background. We compare results from a $\\Lambda$CDM simulation with a $\\Lambda$WDM (2keV) simulation that forms significantly less small scale structure.  We show how high mass and force resolution in both the gas and dark matter components play an important role in solving the angular momentum catastrophe claimed from previous simulations of galaxy formation within the hierarchical framework.  Hence, a large disk forms without the need of strong energy injection, the z = 0 galaxies lie close to the I--band Tully--Fisher relation, and the stellar material in the disk component has a final specific angular momentum equal to 40\\% and 90\\% of that of the dark halo in the $\\Lambda$CDM and $\\Lambda$WDM models respectively.  If rescaled to the Milky Way, the $\\Lambda$CDM galaxy has an overabundance of satellites, with a total mass in the stellar halo 40\\% that in the bulge$+$disk system. The $\\Lambda$WDM galaxy has a drastically reduced satellite population and a negligible stellar spheroidal component. Encounters with satellites play only a minor role in disturbing the disk.  Satellites possess a variety of star formation histories linked to mergers and pericentric passages along their orbit around the primary galaxy. In both cosmologies, the galactic halo retains most of the baryons accreted and builds up a hot gas phase with a substantial X--ray emission.  Therefore, while we have been successful in creating a realistic stellar disk in a massive galaxy within the $\\Lambda$CDM scenario, energy injection emerges as necessary ingredient to reduce the baryon fraction in galactic halos, independent of the cosmology adopted. ",
        "introduction": "\\label{sec:intro}  Within the current paradigm of galaxy formation, realistic disks form only if gas retains most of its angular momentum gained by torques from nearby structures while it cools at the center of cold dark matter (CDM) halos (White \\& Rees 1978, Fall \\& Efstathiou 1980, Fall 1983, Mo, Mao \\& White, 1998).  First simulations of galaxy formation that included star formation (Lake \\& Carlberg 1988, Katz 1992) provided strong evidence that hierarchical models do create rotationally supported stellar systems.  However, simulations of galaxy formation in a full cosmological context have not yet been able to form realistic disk galaxies: dynamical friction suffered by dense gaseous lumps and subsequent catastrophic angular momentum loss caused typical disk scale lengths to come short of those observed (Navarro \\& White 1994).  In addition, high resolution simulations of dark matter halos in CDM models have far more substructure (\"satellites\") than observations; encounters with these satellites could destroy the stellar disk component of spiral galaxies (Moore et al. 1999).    If these dark subhalos were the main problem, any mechanism capable of reducing the lumpiness of galaxy assembly would lead to realistic disks. However resolution simulations (N $\\sim 10^4$) also suffer from excessive artificial viscosity that would overestimate the angular momentum transfer, as was early recognized in previous works (e.g (Sommer--Larsen, Gelato \\& Vedel 1999).  Poorly resolved disk will also suffer from two--body effects which will lead to substantial disk angular momentum loss even after its eventual formation (Mayer et al. 2001a).   Together, the failure to build galactic disks similar to those observed and the overabundance of substructure posed a formidable challenge to the $\\Lambda$CDM paradigm, suggesting that a fundamental ingredient is missing from our understanding of the formation of galactic and subgalactic structures. As possible solutions, strong SN feedback (e.g Thacker \\& Couchman 2000 (TC) and 2001 (TC01)), an external UV background produced by QSOs and massive stars (Quinn, Katz, \\& Efstathiou 1996, Benson et al. 2002, Somerville 2002) and alternatives to $\\Lambda$CDM (Spergel \\& Steinhardt 2000, among others) have been advocated.  Among alternative dark matter models, Warm Dark Matter (WDM) (Pagels \\& Primack 1982) emerged as a good candidate as its lack of power at small scales could prove beneficial in reducing the number of subhalos and allowing a smoother accretion of gas, forming larger disks (Bode, Ostriker \\& Turok 2001, Col{\\' i}n, Avila-Reese, \\& Valenzuela 2000). Some of these solutions have been explored in recent studies (e.g. TC01, Sommer--Larsen \\& Dolgov 2001, Sommer--Larsen, Gotz \\& Portinari 2003, Abadi et al 2003).  These most recent works have been successful in forming a rotationally supported stellar component although in many cases a massive stellar spheroid is formed as well. In runs that use less than a few thousand particles to represent the virialized part of a galaxy dark matter halo (e.g Eke, Navarro \\& Steinmetz 2001) the rotating stellar component does not allow to distinguish a disk.  In TC01, where $\\sim$ 1.5 10$^4$ dark matter (DM) particles with the halo virial radius where used, the ``bulge'' component is still almost twice as massive as the thin disk.  The bulge component decreases to only 50\\% of the stellar disk in the simulation presented by Abadi et al (2003) that had 3.6 10$^4$ DM particles within the virial radius of their simulated galaxy.  While these results are encouraging, previous simulations do not clarify if stronger feedback was necessary, or just sufficient to avoid systematic loss of angular momentum during the formation of galactic disks. Runs with larger particle numbers also had larger disks, suggesting that resolution effects still play a major role.  Important differences in the results might also arise from the different merging histories of each individual halo simulated, the spin of their halos (e.g TC01 simulated a halo with a larger than average spin which in turn makes it easier to form larger disks) and the different SN feedback recipes used in various works.   We have performed a new set of smoothed particle hydrodynamic (SPH) simulations that use a high number of particles and high force resolution to follow the formation and internal structure of the stellar component of galaxies.  These simulations implement cooling, star formation, supernova feedback and a UV background.  In this paper, we describe the formation and evolution of a realistic disk dominated galaxies formed in a cosmological numerical simulation and compare a high-resolution spiral galaxy formed in 2 different ($\\Lambda$CDM and $\\Lambda$WDM) flat cosmologies. Rather than simulate a large sample of galaxies, our strategy is to achieve the highest possible resolution with the available computing resources to avoid unwanted numerical systematics. Our aim is to find a good, clear cut case where a realistic disk forms in a fully cosmological environment rather than exploring the full range of galactic morphologies. We also evaluate the various numerical and physical processes responsible for the angular momentum loss of the gaseous and stellar components.  We describe the initial conditions in section 2, we present our results and discuss their implications in sections 3 and 4 respectively. ",
        "conclusions": "We simulated the formation of a disk galaxy in a cosmological context to the present time with a high force and mass resolution.  Without resorting to strong feedback recipes, we successfully formed an object with bulge-to-disk mass ratios, scale lengths, current star formation, disk age and dynamical properties representative of those observed in present day massive spirals. Our results show that the large angular momentum loss reported in some early works was in at least in part due to insufficient mass or force resolution. High resolution plays a significant role in simulating correctly the formation of disk galaxies in a cosmological context, independently of the SN feedback recipe adopted.  We estimate that at a resolution of $\\sim 10^5$ particles within $R_{vir}$ for each of the DM, gas and stellar components numerical effects still account for 10\\% of the disk angular momentum loss during its formation and evolution.     WDM improves on the standard$\\Lambda$CDM model on a few aspects and leads to testable differences.  Namely it reduces the number of satellites and the galaxy forms with a smaller bulge and almost no stellar halo component. Its disk is more gas rich at the present time and has colder kinematics (lower $v_{rot}/\\sigma$). However, WDM fails to solve the issue for which it was introduced: WDM does not make substantially more extended disks.  The presence of a substantial fraction of hot gas in the halo, independent of the cosmology adopted, results in an X--ray emission higher than estimates for galaxies of similar v$_{c}$ to those in our study. Numerical and semi-analytical results (e.g Borgani et al. 2002, Babul et al. 2002) show that significant energy injection leading to an entropy floor of 50--100 keV/cm$^2$, is needed on the mass scale of small groups (just a factor of ten higher than the mass of the galaxy in our study) to reduce the fraction of hot baryons and reconcile models with the observed X--ray Luminosity--Temperature relation. However, more numerical studies including cooling and detailed star formation are needed.  There is little evidence for SN energy injection of this magnitude in dwarf galaxies (Martin 1999), but it is possible that a combination of SN and AGN activity linked to bulge formation (Binney, Gerhard \\& Silk 2001) could simultaneously reduce the baryon fraction in the halo and quench early star formation in the disk and the satellites, making a large fraction of them ``dark''. If energy injection is required regardless of cosmology, then the need to invoke non standard dark matter models will be substantially weakened."
    },
    "0207/astro-ph0207272_arXiv.txt": {
        "abstract": "Core collapse supernovae (SN) are the final stages of stellar evolution in massive stars during which the central region collapses, forms a neutron star (NS), and the outer layers are ejected. Recent explosion scenarios assumed that the ejection is due to energy deposition by neutrinos into the envelope but detailed models do not produce powerful explosions. There is new and mounting evidence for an asphericity and, in particular, for axial symmetry in several supernovae which may be hard to reconcile within the spherical picture. This evidence includes the observed high polarization and its variation with time, pulsar kicks, high velocity iron-group and intermediate-mass elements material observed in remnants,  direct observations of the debris of SN1987A etc. Some of the new evidence is discussed in more detail. To be in agreement with the observations, any successful mechanism must invoke some sort of axial symmetry for the explosion. We consider jet-induced/dominated explosions of core collapse supernovae. Our study is based on detailed 3-D hydrodynamical and radiation transport models. We find that the observations can be explained by low velocity, massive jets which stall well within the SN envelope. Such outflows may be produced by  MHD-mechanisms, convective dominated accretion disks on the central object or asymmetric neutrino emissions. Asymmetric density/chemical distributions and, for SN2002ap, off-center energy depositions have been identified as crucial for the interpretation of the polarization. ",
        "introduction": "\\vskip -0.2cm There is a general agreement that the explosion of a massive star is caused by the collapse of its central parts into a neutron star or, for massive progenitors, into a black hole. The mechanism of the energy deposition into the envelope is still debated.  The process likely involves the bounce and the formation of the prompt shock (e.g. Van Riper 1978), radiation of the energy in the form of neutrinos (e.g. Bowers \\& Wilson 1982), and the interaction of the neutrino with the material of the envelope and  various types of convective motions ( e.g. Herant et al. 1994, Burrows et al. 1995, M\\\"uller \\& Janka 1997, Fryer \\& Waren 2002), rotation (e.g. LeBlanc \\& Wilson 1970, M\\\"onchmeyer et al. 1991) and magnetic fields (e.g. LeBlanc \\& Wilson 1970, Bisnovati-Kogan 1971, Symbalisty 1984). Currently, the most favored mechanism invokes the neutrino deposition but it is neither clear whether it can provide powerful explosions, can account for large scale asymmetries or for large kick velocities as observed in neutron stars (see below). Likely, a quantitative model of the core collapse must eventually include all the elements mentioned above (H\\\"oflich et al. 2001). In this paper we study the effects and observational consequences of an asymmetric, jet-like deposition of energy inside the envelope of SN. ",
        "conclusions": ""
    },
    "0207/astro-ph0207102_arXiv.txt": {
        "abstract": "I present a {\\it XMM-Newton} observation of the accretion driven millisecond X-ray pulsar SAX J1808.4--3658 during its 2000 outburst. The source was conclusively detected, albeit at a level of only $\\sim2\\times10^{32}$ \\Lunit. The source spectrum could be fitted with a power-law model (with a photon index of $\\sim$2.2), a neutron star atmosphere model (with a temperature of $\\sim$0.2 keV), or with a combination of a thermal (either a black-body or an atmosphere model) and a power-law component.  During a {\\it XMM-Newton} observation taken approximately one year later, the source was in quiescence and its luminosity was a factor of $\\sim$4 lower. It is possible that the source spectrum during the 2000 outburst was softer than its quiescent 2001 spectrum, however, the statistics of the data do not allow to make a firm conclusion. The results obtained are discussed in the context of the 2000 outburst of SAX J1808.4--3658 and the quiescent properties of the source. ",
        "introduction": "}  The X-ray transient SAX J1808.4--3658 was discovered in September 1996 when it exhibited a weak outburst lasting only a few weeks \\citep{intzandetal1998,intzandetal2001}.  In April 1998 the source was found to be in outburst again \\citep{marshall1998} and it was discovered that the source exhibits coherent millisecond X-ray oscillations with a frequency of approximately 401 Hz \\citep{wvdk1998}. In early 2000, the source exhibited a third outburst during which it showed erratic luminosity behavior with luminosity swings of three orders of magnitude within a few days \\citep{wijnandsetal2001_rxte,wijnandsetal2002_bepposax}. This erratic behavior lasted for several months before the source returned to quiescence. Very recently, in October 2002, a fourth outburst of the source was detected \\citep{markwardtetal2002} during which its peak luminosity was very similar to that observed during the 1996 and 1998 outbursts.  In quiescence, SAX J1808.4--3658 has been observed on several occasions with the {\\it BeppoSAX} and {\\it ASCA} satellites \\citep{stellaetal2000,daw2000,wijnandsetal2002_bepposax}. The source was very dim in quiescence, with a luminosity close to or lower than $10^{32}$ \\Lunit. Due to the limited angular resolution of {\\it BeppoSAX}, doubts were raised as to whether the source detected by this satellite was truly SAX J1808.4--3658 or an unrelated field source \\citep{wijnandsetal2002_bepposax}. \\citet{campanaetal2002} reported on a quiescent observation of the source performed with {\\it XMM-Newton} which resolved this issue. They detected the source at a luminosity of $5\\times10^{31}$ \\Lunit~and found that the field around SAX J1808.4--3658 is rather crowded with weak sources. Two such sources are relatively close to SAX J1808.4--3658 and might have conceivably caused a systematic positional offset during the {\\it BeppoSAX} observations of SAX J1808.4--3658.   After it was found that the source had become active again in January 2000 \\citep{vanderklisetal2000}, a Director's Discretionary Time (DDT) observation request was submitted to the director of {\\it XMM-Newton} to study the outburst X-ray spectrum of SAX J1808.4--3658. This request was granted and on March 6, 2000, {\\it XMM-Newton} observed the source.  Due to the extreme variability of the source during its 2000 outburst \\citep{wijnandsetal2001_rxte}, the {\\it XMM-Newton} observation was performed during times when the source had very low luminosities (see \\S~\\ref{section:observation}). Because of this the DDT observation was considered to be in conflict with an approved Cycle 1 observation on this source with the purpose of studying the quiescent X-ray properties of SAX J1808.4--3658 \\citep[see][for the results of that observation]{campanaetal2002}. Consequently, the DDT observation was not made public until recently.  In this paper I discuss this DDT observation during which the source was conclusively detected at a luminosity of $\\sim2\\times10^{32}$ \\Lunit. ",
        "conclusions": "This paper discusses the {\\it XMM-Newton} observation of the accretion-driven millisecond X-ray pulsar SAX J1808.4--3658 performed during its 2000 outburst. Similar to the {\\it BeppoSAX} observations performed around the same time (one of those observations was on the same day as the {\\it XMM-Newton} one), the {\\it XMM-Newton} observation revealed only a weak X-ray source with a 0.5--10 keV luminosity of $\\sim2\\times10^{32}$ \\Lunit~at the position of SAX J1808.4--3658. Despite its weakness, the source was the brightest one on the central CCD proving that the {\\it BeppoSAX} source detected by \\citet{stellaetal2000} and \\citet{wijnandsetal2002_bepposax} is indeed SAX J1808.4--3658 and not an unrelated field source. The systematic off-set between the measured and the true position of SAX J1808.4--3658 in those {\\it BeppoSAX} observations can likely be explained by two nearby field sources. These sources might have also contaminated the flux assigned to SAX J1808.4--3658 \\citep[see also][]{campanaetal2002}.  \\citet{campanaetal2002} reported on a quiescent observation of SAX J1808.4--3658 performed with {\\it XMM-Newton} about a year after the end of the 2000 outburst (the spectral fit parameters reported by \\citet{campanaetal2002} are also listed in Tab.~\\ref{tab:spectrum} for comparison). They found that the source had a luminosity of only $\\sim5 \\times 10^{31}$ \\Lunit, about a factor of 4 lower than what was measured during the 2000 outburst observation. This demonstrates that at very low luminosities, SAX J1808.4--3658 can exhibit variability and indicates that the source probably displays variability in its quiescent state. However, caution is advised because the 2000 outburst observation was performed when the source exhibited violent behavior \\citep{wijnandsetal2001_rxte}. Therefore, it is still possible that in ``true'' quiescence the source will be observed consistently at the low level reported by \\citet{campanaetal2002}.  The spectrum of SAX J1808.4--3658 during the 2000 outburst appears softer than its quiescent spectrum: when the spectrum is fitted with a power-law model, the photon index is 2.2\\pp0.3 during the 2000 outburst vs. 1.5$^{+0.2}_{-0.3}$ in quiescence. The fact that the 2000 outburst spectrum can be accurately fit with a neutron star atmosphere model or with a combination of a thermal plus power-law component \\citep[in contrast with the quiescent spectrum;][]{campanaetal2002}, also suggests a difference in the source spectra between the two epochs. It is interesting to note that in the two-component model, the photon index during the 2000 outburst observation is very similar to that measured during the quiescent observation, suggesting that the shape of the power-law component might not have changed considerably between the different epochs (although the flux of this component was still a factor of $\\sim$3 higher during the 2000 outburst than in quiescence).  Due to the limited statistics of the data, the exact degree of spectral variability cannot be determined nor can the exact cause of these possible variations be established.   During the 2000 outburst, the source fluctuated in luminosity by over 3 orders of magnitude on timescales of days \\citep{wijnandsetal2001_rxte}.  However, the {\\it BeppoSAX} observations reported by \\citet{wijnandsetal2002_bepposax} provided only rough estimates of the flux of SAX J1808.4--3658 between 2000 March 5 and 8.  With the {\\it XMM-Newton} detection of the source on 2000 March 6, we now have a clearer picture of how dim the source could become during certain phases of its 2000 outburst.  As stated by \\citet{wijnandsetal2001_rxte}, the large luminosity variations observed during the 2000 outburst are difficult to understand as due to similarly dramatic variations in the mass accretion rate. It is more likely that only modest variations in the accretion rate can trigger transitions between two significantly different luminosity states. For example, centrifugal inhibition of accretion by the neutron star's magnetic field is expected below a certain critical accretion rate \\citep[the 'propeller regime';][]{is1975} and small but erratic variations in the accretion rate around this critical rate could in principle give rise to the enormous luminosity swings observed during the 2000 outburst.  When the source is in this propeller regime, accretion is inhibited, but it is evident from the large brightness fluctuations during the 2000 outburst, that a considerable amount of matter was still available in the accretion disk. According to \\citet{campanaetal2002}, a pure propeller contribution is ruled out in quiescence since this mechanism is expected to stop operating at luminosities below $10^{33}$ \\Lunit~because the source should turn on then as a radio pulsar. \\citet{stellaetal2000} and \\citet{campanaetal2002} suggested that a possible explanation for the quiescent flux is the emission from the shock front between the relativistic wind of the radio pulsar and the matter out-flowing from the companion star. Therefore, the flux during the 2000 outburst observation could be higher than that observed in quiescence because of the large amount of matter still present close to the neutron star.  \\citet{db2003} suggested an alternative explanation for the quiescent emission of SAX J1808.4--3658 in which this emission is produced by direct dipole radiation from the radio pulsar. The quiescent X-ray luminosity and spectral shape of SAX J1808.4--3658 are indeed consistent with those observed from several field millisecond radio pulsars \\citep[see][for a recent review]{bp2002}, but \\citet{grindlay2002} found that the millisecond pulsars in the globular cluster 47 Tuc have a predominantly soft spectral shape suggestive of a thermal origin in contrast to what is observed for SAX J1808.4--3658. Conclusive proof for an active radio pulsar in SAX J1808.4--3658 would come from the detection of radio pulsations during quiescence, although such a detection might be inhibited by the ambient matter still present in the system \\citep[see also the discussion in][]{db2003}.  \\citet{campanaetal2002} noted that the quiescent properties of SAX J1808.4--3658 are remarkably different from those observed for other quiescent neutron star systems: (a) its 0.5--10 keV luminosity is the lowest observed so far for any neutron star system, and (b) its spectrum is dominated by a power-law component instead of a thermal component.  The quiescent emission of neutron star X-ray transients is most often explained by thermal emission from the neutron star surface releasing the heat deposited in the crust and core of the neutron star during outburst \\citep[see][ and references therein]{bbr1998}. In this model, the exact luminosities of the systems should depend on their time-averaged accretion rates \\citep{campanaetal1998,bbr1998}. Due to the low peak luminosity of SAX J1808.4--3658, \\citet{bbr1998} predicted that this source should be rather faint in quiescence. At first sight the low detected quiescent luminosity is consistent with this prediction, but no strong evidence could be found for a thermal component in the quiescent spectrum obtained. This indicates that the thermal luminosity of this source is very low, implying a rapidly cooling neutron star which would require enhanced core cooling processes to explain \\citep{campanaetal2002}.  The combination of the low time-averaged accretion rate and the possibility of rapid core cooling might also be able to explain why the quiescent spectra of SAX J1808.4--3658 is dominated by the power-law component.  If in the 'ordinary' quiescent systems the thermal component were to drop to low luminosity levels similar to those of SAX J1808.4--3658, then their quiescent spectra would be dominated by the power-law component (assuming that this component remains unchanged), just like in SAX J1808.4--3658. However, it remains to be determined if the power-law components in the different types of systems are due to the same mechanism, or if the quiescent spectrum of SAX J1808.4--3658 is due to some other process, likely related to the different configuration and/or magnetic field strength of the neutron star in SAX J1808.4--3658.  To get more insight into the nature of the unique quiescent properties of SAX J1808.4--3658, more neutron star transients must be detected in quiescence. Particularly interesting systems are the two other, very recently discovered, accretion driven millisecond X-ray pulsars XTE J1751--305 \\citep{markwardtswank2002} and XTE J0929--314 \\citep{rss2002}. It would be interesting to compare their quiescent properties with those of SAX J1808.4--3658 and correlate possible similarities and/or differences in quiescence with the outburst properties of those sources."
    },
    "0207/astro-ph0207428_arXiv.txt": {
        "abstract": "We use archival HST/WFPC2 V and I band images to show that the optical counterpart to the ultra-luminous x-ray source NGC~5204 X-1, reported by Roberts et~al., is composed of two sources separated by 0.5''. We have also identified a third source as a possible counterpart, which lies 0.8'' from the nominal X-ray position. PSF fitting photometry yields V-band magnitudes of 20.3, 22.0 and 22.4 for the three sources. The V$-$I band colours are $0.6$, $0.1$, and $-0.2$, respectively (i.e. the fainter sources are bluer). We find that all V$-$I colours and luminosities are consistent with those expected for young stellar clusters (age $<10$~Myr). ",
        "introduction": "EINSTEIN and ROSAT observations show that at high energies most nearby (spiral) galaxies are dominated by a relatively small number of discrete, compact, but highly luminous x-ray sources.  These so-called ultra-luminous x-ray sources (ULXs) have luminosities of $L_{\\rm x}\\sim10^{39}-10^{41}$~erg~s$^{-1}$, far in excess of the Eddington limit for spherical accretion onto a neutron star (Roberts and Warwick 2000).  Nearly 20\\% of ULXs appear to be associated with supernova remnants (SNR, Roberts et~al. 2002), while the rest are thought to be powered by accretion onto a compact object. One possibility is that ULXs represent the missing class of intermediate mass ($10^{2}-10^{5}$M$_{\\odot}$) black holes (Colbert and Mushotzky 1999). However, recent Chandra observations reveal a strong association between ULXs and active star-forming regions (Fabbiano et~al. 2001, Lira et~al. 2002), that are far too young to contain massive black holes (King et~al. 2001).  Subsequently, at least two alternative scenarios have been proposed.  In the first, ULXs are stellar mass black hole binaries undergoing a period of super-Eddington accretion (e.g. Watarai et~al. 2001, Begelman et~al. 2002).  In the second, the apparent super-Eddington luminosities are due to the beamed x-ray emission from an otherwise normal intermediate/high-mass x-ray binary (King et~al. 2001, Georganopoulos et~al. 2002). While both scenarios tie in neatly with the observed association between ULXs and active star forming regions (Long et~al. 2002, Terashima and Wilson 2002, Zezas et~al. 2002), the latter requires a significantly larger number of ULXs to be present.  Roberts et~al. (2001) identified the first optical counterpart to a ULX, NGC~5204~X-1, based on its Chandra position and optical multi-fibre spectroscopy.  The counterpart, located within 1'' of the nominal Chandra position, displays a blue featureless continuum and is surrounded by an ionized bubble of gas with a diameter of some 360~pc (Pakull and Mirioni 2002). Here, we present archival HST/WFPC2 V and I band images of this source, together with an improved analysis of the multi-fibre data. These data allow us to place more robust constraints on the nature of the counterpart, and thus obtain a deeper insight into the ULX phenomenon.  \\begin{figure*} \\epsfig{width=6.8in,file=MC645Lrv_fig1.eps} \\caption{Upper panel - HST WFPC2 V (left-hand side) and I (right-hand side) band images of NGC~5204 showing the location of the optical counterpart. Lower panel - enlarged V and I band images reveal that the optical counterpart is resolved into two components (HST-1 and HST-2) separated by $\\sim$0.47''. A third source (HST-3) is located 0.8'' east of the nominal Chandra position.  The 1.0'' radius ring shows the formal error on the ULX position. North is up.} \\end{figure*}  \\section[]{HST/WFPC2 images of NGC~5204~X-1}  We have obtained archival HST/WFPC2 V and I band (F606W and F814W) images of the dwarf magellanic type galaxy NGC~5204 (Fig~1 - upper panel). This data (now public) was obtained as part of an ongoing HST/WFPC2 snapshot survey of nearby dwarf galaxies (Seitzer, PI).  The HST/WFPC2 V and I band images (Fig~1, lower panel) show that the optical counterpart can be resolved into at least two sources, here designated HST-1 and HST-2, separated by 0.47''.  We also identify a third source designated HST-3, located 0.8'' to the east of the nominal Chandra position. This source was not detected in the fibre data of Roberts et~al. but due to its close proximity remains a possible counterpart to the ULX.  The observed separations correspond to distances of 10$-$16~pc at the distance of NGC~5204 (4.8~Mpc).  We have performed PSF fitting photometry on the V and I band images using HSTphot (Dolphin 2000).  Unfortunately, only one image was available in each filter, hampering cosmic ray rejection.  Given the difficulty of identifying real point sources, we decided against iterative adjustment of HSTphot's standard PSFs on the basis of the fit residuals.  The chip positions, RA, DEC, V and I band magnitudes (referenced to a 0.5'' aperture), colours, and distances from the nominal Chandra position ($\\Delta$), are quoted in Table~1 for each source.  \\begin{table*} \\centering \\caption{HSTphot PSF fitting photometry} \\begin{tabular}{@{}ccccccccrc@{}} Source & Chip no. & \\multicolumn{2}{c}{Chip position} & RA & DEC & V  &  I  & V$-$I & $\\Delta$\\\\ &   & X      & Y      &               &              &      &      &  & (arcsecs) \\\\ HST-1  & 2 & 273.75 & 343.96 & 13:29:38.57 & 58:25:05.78 & 20.3 & 19.7 & 0.6 & 0.34\\\\ HST-2  & 2 & 273.37 & 348.66 & 13:29:38.51 & 58:25:05.86 & 22.0 & 21.9 & 0.1 & 0.82\\\\ HST-3  & 2 & 274.44 & 332.56 & 13:29:38.71 & 58:25:05.62 & 22.4 & 22.6 & $-$0.2 & 0.80 \\\\ \\end{tabular} \\end{table*}  Our HSTphot photometry reveals that all sources are blue (i.e. $F_{\\lambda}$ decreases with increasing $\\lambda$), consistent with the finding of Roberts et.~al. for a single unresolved source. However the two fainter sources are significantly bluer than HST-1. The combined magnitudes of HST-1 and HST-2 are 20.1 in V and 19.6 in I, with a V$-$I colour of 0.5.  These magnitudes are referenced to a 0.5'' aperture, becoming 20.0 and 19.5 when corrected to infinite aperture. We do not include HST-3 in the combined magnitude estimate, as it does not fall within the same fibre in the data of Roberts et~al.  The best-fit Chandra position for the ULX (determined using WAVDETECT; Freeman et~al. 2002) -- $\\alpha=$13~29~38.61, $\\delta=$+58~25~05.7 -- is located 0.34'' east of HST-1.  The 1$\\sigma$ error associated with the Chandra reference frame is $\\pm 0.6$'' (Aldcroft et~al. 2000). However, errors of up to 1.0'' are not uncommon when linking the Chandra and HST reference frames (see e.g. Grindlay et~al. 2001).  Thus while HST-1 is the preferred location for the optical counterpart, neither HST-2 or HST-3 can be ruled out.  \\begin{figure} \\epsfig{width=3.4in,height=3.5in,file=MC645Lrv_fig2.eps} \\caption{The optical counterpart to NGC~5204 X-1 after removal of contaminating nightsky and host galaxy emission. The open triangles represent the derived V and I band fluxes for the combined sources as estimated from the HST/WFPC2 images.} \\end{figure}  \\section[]{Comparison with WHT/INTEGRAL multi-fibre spectroscopy}  We have reanalysed the optical multi-fibre data on NGC~5204~X-1 from Roberts et~al. (2001) with the aim of improving the photometric accuracy of their spectrum. This allows a direct comparison of their results with our new results from HSTphot. Our new analysis takes better account of fibre throughput corrections (using an improved throughput image, with repeatability errors of $<0.1$\\%) and contamination by the nightsky background. More importantly, we have also tried to correct for diffuse emission from within the host galaxy.  The reconstructed multi-fibre image of our source, shows that emission from the optical counterpart lies within a single central fibre and its surrounding nearest 6 neighbours. If we adopt the procedure followed by Roberts et~al. and correct for nightsky background only, we derive a V-band magnitude of 19.9. This is 0.2 magnitudes fainter than found by Roberts et~al.; this difference is due partly to our improved throughput correction and partly to the use of a median (instead of a mean) nightsky background.  As noted above, we have also tried to account for diffuse emission from the host galaxy, which shows a clear west-east gradient (see Fig~1). We do this by taking a median of the local background continuum over 18 fibres arranged in a hexagonal ring $\\sim3$ arcseconds from the source.  This avoids including source flux in the background fibres. As in Roberts et~al. the resultant spectrum (Fig~2) is blue and featureless at the resolution of our instrument (6~\\AA\\ per 2 pixel resolution element).  However, removal of local galaxy contributions results in a cleaner, fainter, and marginally bluer spectrum. The measured flux near 5500~\\AA\\ translates to a V-band magnitude of $20.4\\pm0.1$. This is 0.7~mag fainter than the magnitude quoted by Roberts et~al. and 0.4~mag fainter than that derived from the HST/WFPC2 V-band image.  The discrepancy between the WHT V-band flux estimate and that determined directly from the HST/WFPC2 V-band image most likely results from the intermittent presence of high cloud during some of the ground-based observations. This gives an indication of the accuracy of the absolute spectral flux calibration of the fibre data. ",
        "conclusions": "We have used archival HST/WFPC2 V and I band images to show that the previously reported optical counterpart to the ultra-luminous x-ray source NGC~5204 X-1, can be resolved into two sources separated by 0.47''. We have also identified a third possible counterpart 0.8'' east of the nominal Chandra position. All three sources appear blue (i.e. $F_{\\lambda}$ decreases with increasing $\\lambda$), with the fainter sources being significantly bluer. PSF fitting photometry with HSTphot yields source magnitudes of 20.3, 22.0 and 22.4 in V. Assuming that all three sources are co-eval, the V$-$I band colours suggest that NGC~5204~X-1 lies within a young ($<$10~Myr) stellar cluster.  This is consistent with the idea that ULXs are intermediate/high-mass x-ray binaries."
    },
    "0207/astro-ph0207334_arXiv.txt": {
        "abstract": "{We present the results of a detailed non-LTE analysis of the ultraviolet and optical spectrum of the O6.5\\,Iaf$^+$ star \\object{HD~153919} - the mass donor in the high-mass X-ray binary \\object{4U1700-37}. We find that the star has a luminosity log(L$_{\\ast}$/L$_{\\odot}$)=5.82$\\pm$0.07, T$_{\\rm eff}$=35,000$\\pm$1000~K, radius R$_{\\ast}$=21.9$^{+1.3}_{-0.5}$~R$_{\\odot}$, mass-loss rate \\.{M}=9.5$\\times$10$^{-6}$~M$_{\\odot}$~yr$^{-1}$, and a significant overabundance of nitrogen (and possibly carbon) relative to solar values.  Given the eclipsing nature of the system these results allow us to determine the most likely masses of both components of the binary via Monte Carlo simulations. These suggest a mass for \\object{HD~153919} of M$_{\\ast} = 58 \\pm 11$~M$_{\\odot}$ - implying the initial mass of the companion was rather high ($\\ga$60~M$_{\\odot}$). The most likely mass for the compact companion is found to be M$_x$ =2.44$\\pm$0.27~M$_{\\odot}$, with only 3.5~per cent of the trials resulting in a mass less than 2.0~M$_{\\odot}$ and none less than 1.65~M$_{\\odot}$. Such a value is significantly in excess of the upper observational limit to the masses of neutron stars of 1.45~M$_{\\odot}$ found by Thorsett \\& Chakrabarthy (\\cite{thorsett}), although a mass of 1.86~M$_{\\odot}$ has recently been reported for the \\object{Vela X-1} pulsar (Barziv et al. \\cite{barziv}).  Our observational data is inconsistent with the canonical neutron star mass and the lowest black hole mass observed ($\\ga$4.4~M$_{\\odot}$; \\object{Nova Vel}).  Significantly changing observational parameters {\\em can} force the compact object mass into either of these regimes but, given the strong proportionality between M$_{\\ast}$ and M$_x$, the O-star mass changes by factors of greater than 2, well beyond the limits determined from its evolutionary state and surface gravity.  The low mass of the compact object implies that it is difficult to form high mass black holes through both the Case A \\& B mass transfer channels and, if the compact object is a neutron star, would significantly constrain the high density nuclear equation of state. ",
        "introduction": "First detected by the {\\em Uhuru} satellite (Jones et al. \\cite{jones}) the eclipsing X-ray source \\object{4U~1700-37} was quickly associated with the luminous O6.5 Iaf$^{+}$ star \\object{HD~153919}, confirming \\object{4U~1700-37} as a high-mass X-ray binary (henceforth HMXB). HMXBs are systems composed of an OB star and a compact companion (neutron star or black hole), with the X-ray emission resulting from the accretion of material by the compact companion. In the subclass of supergiant HMXB systems material is accreted either via Roche lobe overflow or directly from the powerful stellar wind of the OB primary. Given that \\object{HD~153919} slightly underfills its Roche Lobe (e.g. Conti \\cite{conti}) mass transfer proceeds via the latter mechanism.  Although an orbital period of $\\sim$3.412~days (Jones et al. \\cite{jones}) was quickly identified for \\object{4U~1700-37}, extensive searches (e.g. Rubin et al. \\cite{rubin} and references therein) have failed to identify any other X-ray periodicities within the system that might correspond to the pulse period for a possible neutron star (although Konig \\& Maisack \\cite{konig} claim the presence of a 13.81~day period in CGRO BATSE \\& RXTE ASM datasets). Given the absence of any X-ray pulsations and the unusually hard nature of the spectrum various authors (e.g. Brown, Weingartner \\& Wijers \\cite{brown}) have suggested that the compact companion could be a low mass black hole rather than a neutron star. However, Reynolds et al.  (\\cite{reynolds}) point out that the 2-200~keV spectrum of \\object{4U~1700-37} differs from those commonly observed for black hole candidates such as \\object{Cygnus X-1}.  Given that the X-ray spectrum of \\object{4U~1700-37} is qualitatively similar to those of accreting neutron stars they suggest that the compact object is also a neutron star, and explain the lack of pulsations as due to either a weak magnetic field or an alignment of the magnetic field with the spin axis.  With a spectral type of O6.5 Iaf$^{+}$, \\object{HD~153919} is the hottest and potentially most massive mass donor of any of the the HMXB systems. As such, determination of its fundamental parameters (radius, temperature, mass and chemical composition) is of importance given that these will potentially provide valuable insights into the evolution of very massive stars and their ultimate fate.  In particular by determining the masses of the components of HMXB systems, limits to the progenitor masses for neutron stars and black holes in such systems can be found\\footnote{we note that Wellstein \\& Langer (\\cite{wellstein}) demonstrate that such limits derived from binary systems cannot be {\\em directly} applied to single stars} while the chemical composition of the mass donor can shed light on the pre-supernova (SN) mass transfer mechanisms.  The paper is ordered as follows. Section 2 describes the determination of the physical properties of \\object{HD~153919}; the complete dataset and the non-Local Thermal Equilibrium (non-LTE) code used to analyse it. Section 3 describes the Monte-Carlo technique used to determine the masses of both components of the binary system. In Sections 4 \\& 5 we discuss the implications of our results for the evolution of hot massive stars, limits for the progenitor masses of compact objects and the equation of state for nuclear matter.  Finally in Section 6 we summarize the main results of the paper. ",
        "conclusions": "We have performed a sophisticated NLTE analysis on the O6.5 Iaf+ star \\object{HD~153919}, the primary of the HMXB \\object{4U1700-37} and used the results to constrain the masses of both components of the system via a Monte Carlo simulation. Our NLTE model atmosphere analysis leads to parameters for \\object{HD~153919} of log(L$_{\\ast}$/L$_{\\odot}$)=5.82$\\pm$0.07, T$_{\\rm eff}$=35,000$\\pm$1000~K, R$_{\\ast}$=21.9$^{+1.3}_{-0.5}$~R$_{\\odot}$, \\.{M}=9.5$\\times$10$^{-6}$~M$_{\\odot}$~yr$^{-1}$, and an overabundance of nitrogen and possibly carbon over solar metallicities. Combined with the short orbital period of the system this implies a common envelope phase of pre-SN evolution - {\\em despite the mass of the SN progenitor apparently precluding a RSG phase} - leading to the formation of a close WC+O star binary, with the carbon enrichment of \\object{HD~153919} a result of the impact of the stellar wind of the WC star on the surface of the O star.  The Monte Carlo simulations result in masses for the O star and compact object of M$_{\\ast}$=58$\\pm$11~M$_{\\odot}$ and M$_x$=2.44$\\pm$0.27~ M$_{\\odot}$, with none of the 10$^6$ trials resulting in M$_x \\leq$1.65~M$_{\\odot}$, while only 3.5~per cent of the trials result in M$_x \\leq$2~M$_{\\odot}$. Given that no significant mass transfer via RLOF has occured this implies that the initial mass of the SN precursor must have been $\\ga$60~M$_{\\odot}$. Thus even very massive stars can effectively `melt down' to leave rather low mass post-SN remnants. Equally, the masses of both components imply that it is impossible for stars of $\\sim$60~M$_{\\odot}$ to leave 5-10 M$_{\\odot}$ remnants via Case A or B evolution, suggesting that most high mass black holes are instead formed via Case C mass transfer.  The mass of the compact object is found to lie in between the range of masses observed for neutron stars and black holes. Given that M$_{\\ast}$ and M$_x$ are strongly correlated, forcing consistency between M$_x$ and either type of object results in significant discrepancies between M$_{\\ast}$ and evolutionary predictions for the mass of \\object{HD~153919}.  In order to produce consistency with the range of masses observed for neutron stars the O star has to be significantly undermassive for its spectral type and luminosity class.  While theoretical evolutionary tracks for massive ($\\ga$60~M$_{\\odot}$) stars suggest that after a redwards excursion the star will evolve bluewards again with a substantially reduced mass, such a star would show significant chemical enrichment (and H depletion) which is not observed.  Equally, the surface gravity determined from modeling excludes such an anomolously low mass.  Forcing consistency between M$_x$ and the masses of known black hole candidates (M$_x \\ga 4.4$~M$_{\\odot}$) results in M$_{\\ast} \\ga$100~M$_{\\odot}$. Stars of such extreme masses are not expected to go through an O supergiant phase, rather evolving into H depleted WR stars via a H-rich pseudo WR phase, where their high mass loss rate simulates the spectrum of a WR star. Once again, the constraint implied by the surface gravity also appears to exclude this possibility.  We are therefore left with the conclusion that no solution is fully consistent with present expectations for stellar evolution and the chemical abundances and surface gravity of \\object{HD~153919}.  If the compact object has a mass consistent with the observed range of neutron star masses, the O star is significantly undermassive, while if it consistent with the lower limit to black hole masses the O star is overmassive by a similar (or larger) factor.  Finally if - as the Monte Carlo analysis implies - the probable mass of the O star is consistent with evolutionary predictions and the measured surface gravity, the mass of the compact object lies in between the two alternatives.  While our results do not allow us to distinguish between a massive neutron star or a low mass black hole, the existence of a neutron star of mass in the mass range $2.44\\pm 0.27M_{\\odot}$ would significantly constrain the high density nuclear equation of state and provide details about the QCD phase transition complementary to information about the temperature induced transition which will be obtained at RHIC and LHC.  Phenomena which might occur deep in the star such as the appearance of hyperons or a quark matter core would be strongly constrained as the existence of these phases might result in an equation of state too soft to support such a high mass star."
    },
    "0207/astro-ph0207645_arXiv.txt": {
        "abstract": "{\\small We have studied the dependence of the high frequency QPO ($\\sim$ 284 Hz) on the phase of the low frequency QPO ($\\sim$ 6 Hz) in the black hole X-ray binary XTE J1550-564 in the observations of the {\\it Rossi} X-ray Timing Explorer (RXTE) performed on MJD 51241. By selecting the local maxima and the local minima in the light curve on the 6 Hz QPO time scale, we have found the corresponding high frequency QPO frequencies are consistent within 1.5 $\\sigma$. However, the average central QPO frequency of the maxima and the minima is about 2.0 $\\sigma$ lower than the average high frequency QPO frequency obtained in the entire observation. This marginally suggests that the high frequency QPOs probably varies in frequency on short time scales. We briefly discuss these results and their consequences.} ",
        "introduction": "XTE J1550-564 was identified as a Galactic \"microquasar\" from both X-ray and radio observations. It outbursts have been observed by the {\\it Rossi} X-ray Timing Explorer (RXTE) for a few times. High frequency QPO with a frequency as high as 285 Hz was observed casually from this source in the outbursts in both 1998-1999 (Remillard et al. 1999; Homan et al. 2001) and 2001 (Miller et al. 2001). X-ray state transitions during the rise and the decay of its outbursts were also observed (see e.g. Homan et al. 2001). On the other hand, superluminal jet production has been observed in radio (Hannikainen et al. 2001).  The high frequency QPOs occur ocassionally in the frequency range between 102 Hz and 285 Hz in the 1998-1999 outburst (Remillard et al. 2001; Homan et al. 2001) and the 2000 outburst ( Miller et al. 2001). Miller et al. 2001  found that these high frequency QPOs are in fact two QPOs present simultaneuously, and with a rough frequency ratio of 2:3. Possible evidence for the harmonic origin of the high frequency QPOs has been found also in other microquasars such as GRO J1566-40 (Remillard et al. 2002; see also papers in this proceedings). If these QPOs are indeed harmonic-related features, the QPO width in frequency can put strong constrains on the range that the high frequency QPO possiblely moves. If the QPO frequency is related to the general relativistic frequencies at the inner accretion disk, then the range of the QPO frequency may put constrains on the movement of the inner disk radius. ",
        "conclusions": "The frequencies of the HFQPO in black hole X-ray binaries are thought to be related to some fundamental frequencies at the innermost orbit around the black hole, e.g. combinations of the azimuthal and radial coordinate frequencies in general relativity (Strohmayer 2001), diskoseismic oscillation frequency (Wagoner 1999) or resonance frequency between the Keplerian and radial coordinate frequencies (bramowicz \\& Kluzniak 2001).  If this is true, then the movement of the HFQPO frequencies probably suggests that the radius of the innermost orbit vary. An important consequence is that the maximal fundamental HFQPO frequency is the central HFQPO frequency plus its half width, about 25--40 Hz higher than the HFQPO central frequency which was used in constraining the mass and the spin of the central black hole."
    },
    "0207/astro-ph0207190_arXiv.txt": {
        "abstract": "Bursts of the count rate of extensive air showers (EAS) lead to the appearance of clusters in time series that represent EAS arrival times. We apply methods of nonlinear time series analysis to~20 EAS cluster events found in the data set obtained with the EAS--1000 prototype array. In particular, we use the Grassberger--Procaccia algorithm to compute the correlation dimension of the time series in the vicinity of the clusters. We find that four cluster events produce signs of chaotic dynamics in the corresponding time series. By applying a number of supplementary methods we assess that the nature of the observed phenomenon may indeed be chaotic and thus deterministic. We suggest a simple qualitative model that might explain an origin of EAS clusters in general and ``possibly chaotic'' clusters in particular. Finally, we compare our conclusions with the results of similar investigations performed by the EAS-TOP and LAAS groups. ",
        "introduction": "We have already studied the distribution of arrival times of extensive air showers (EAS) registered with the EAS--1000 prototype array both by methods of classical statistics \\cite{Dubna,Hamburg} and by methods of cluster analysis \\cite{IzvRAN01,clusters'02}. In particular, we presented 20 EAS cluster events---groups of consecutive showers that were registered in time intervals much shorter than expected ones. These groups represent bursts of the EAS count rate and may be considered as cosmic ray bursts, see \\cite{IzvRAN01,clusters'02}. This phenomenon has put forward at least two questions: one on the astrophysical nature of the process, another on its statistical properties. Namely, while the vast majority of sufficiently long samples satisfy the hypothesis for an exponential distribution of time delays between EAS arrival times \\cite{Dubna,Hamburg}, the $\\chi^2$-test performed for samples taken in the vicinity of some of the EAS clusters made us reject this hypothesis. Thus we have decided to apply methods of nonlinear time series analysis to samples that contain EAS clusters in order to clarify dynamical reasons of this situation. This approach has already proved to be a powerful tool of investigation in different fields of science including different fields of astrophysics in general~\\cite{CP87&GRL91} and cosmic ray physics in particular, see \\cite{Aglietta,Bergamasco92,Bergamasco94,Kudela,Japan:AP,Japan2001} and references therein.   Recall that the data set under consideration represents 203~days of regular operation of the array for the period August~30, 1997 to February~1, 1999. The total EAS number in the data set equals 1~668~489. The mean number of charged particles in a shower is of the order of~$1.2\\times10^5$. This corresponds to the energy of a primary particle $\\sim1$~PeV/nucleon. The mean interval between consecutive EAS arrival times is equal to 10.5~s. The discreteness in the moments of EAS registration approximately equals 0.055~s (one tic of the PC clock).  In~\\cite{nonlinear'02}, we have already discussed briefly the results of nonlinear analysis of one of the EAS clusters. In the present paper, we study this cluster in depth and briefly discuss three other clusters that demonstrate nonlinear features. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics[width=0.7\\hsize]% {981111_435-570_t40_NeB9.7e+04.eps} \\end{center} \\caption{Arrival directions (AD) of EAS that form the cluster registered on November~11, 1998 and have zenith angle $\\theta<40^\\circ$: $\\bullet$---AD of EAS with $N_e\\ge9.7\\times10^4$, $\\times$---AD of all other EAS. The curve shows the Galactic plane. Equatorial coordinates are used.} \\label{981111:AD} \\end{figure}  To get a deeper insight into physics of EAS clusters, let us consider the cluster registered on Nov.~11, 1998 from another point of view. A natural question that arises when one finds a burst of EAS count rate is whether primary particles that led to the appearance of the burst had had a common compact source. \\Fig{981111:AD} shows arrival directions of EAS that belong to the cluster and have zenith angle $\\theta<40^\\circ$. Obviously, one cannot give a positive answer to the above question. Still, this seems quite natural since the energy of primary particles in our case is too low to ``remember'' the source. Besides this, the majority of EAS in the cluster represent a normal count rate. Namely, since we register 5.7 EAS/min in average and the duration of the cluster adjusted to $P^*=742$~mm~Hg equals approximately 15~min then we may expect to register 85--86 EAS within this period. Naturally, arrival directions of these EAS should be distributed more or less isotropically. Thus, one may expect to observe a kind of a cluster in arrival directions for only remaining 50~EAS. Still, this does not happen. One can see doublets and triplets of showers with almost coincident arrival directions and a kind of grouping along the Galactic plane but no big groups of EAS are observed.  Another question is whether clusters consist of EAS that are more ``powerful'' than the others. Surprisingly enough, but the answer is negative. As we have already mentioned in Section~1, the mean value of electrons~$\\bar N_e^\\mathrm{tot}$ in EAS in our data set is of the order of $1.2\\times10^5$ particles. To compare, for the above cluster $\\bar N_e \\sim 9.7\\times10^4$, and only 14 EAS have~$N_e > \\bar N_e$, see \\Fig{981111:AD}. Two of these EAS have $N_e\\sim10^6$. On the other hand, at least 41 EAS have $N_e<1/2\\bar N_e$. (We remark that the geometry of the EAS--1000 prototype array does not allow one to obtain parameters of all registered EAS.)  A similar situation is observed for other EAS clusters, both PC and ``ordinary'' ones. Thus, we come to the following observations: (i)~EAS in a cluster do not have a joint source (arrival direction); (ii)~the majority of EAS in a cluster have $N_e<\\bar N_e^\\mathrm{tot}$. This allows us to suggest a simple conceptual model that is intended to explain qualitatively how EAS clusters might be produced. To do this, we need only to assume that a considerable part of EAS that constitute clusters are generated by charged particles.  As is well known, the interstellar space is filled with strong and highly inhomogeneous magnetic fields. Still, average magnetic fields in the heliosphere are not strong enough to influence the dynamics of particles with $E\\gtrsim10$~TeV/nucleon (see, e.g.,~\\cite{Burger}). On the other hand, a series of long-term investigations based on data obtained with Voyager~1, 2, and other space crafts have demonstrated that the large-scale magnetic field strength fluctuations frequently have large amplitudes and are intermittent, and that regions of relatively intense magnetic fields can have a radial extent of more than 10~AU, see~\\cite{Burlaga-IMF} and references therein.  \\begin{figure} \\begin{center} \\includegraphics[width=0.6\\hsize]{model.eps} \\end{center} \\caption{Model of EAS clusters appearance. Notation: LAP---a layer of accelerated particles; ML---magnetic lens.} \\label{981111:model} \\end{figure}   Basing on these results, let us consider the following\tsituation. Suppose there is an extended and sufficiently ``thick'' layer of accelerated particles (LAP) (possibly called a ``wave'') that moves through space towards Earth, see \\Fig{981111:model}. Normally, if it just passes through, an EAS array registers only a few particles originated from the LAP. Now suppose that the LAP meets a region of an extended strong and inhomogeneous magnetic field that works as a lens. If it happens that this ``magnetic lens'' declines particles in way that they get focused in a ``proper'' direction then an array may register an excess of EAS over the normal count rate, i.e., an EAS cluster. Notice that from an observer's point of view showers in the cluster may have very different arrival directions. In our opinion, the fact that the majority of EAS in the observed clusters are less ``powerful'' than an average shower witnesses in favour of this model since we do not need too strong magnetic fields. The duration of a cluster depends on the thickness of the LAP and the time during which the magnetic lens ``works.''  Now let us ask ourselves how does it happen that the time series that represent the majority of EAS clusters\t(as well as the other data) are stochastic but some clusters seem to demonstrate signs of chaotic dynamics. At the moment, we cannot give a definite answer but can only conjecture that different factors may be involved in this phenomenon. Among them, one can mention a possibly fractal nature of the interstellar medium (see, e.g.,~\\cite{Lagutin} and references therein), nonlinear mechanisms of particles acceleration, etc. A very interesting possibility for EAS time series to become chaotic suggests the fact that the large-scale fluctuations in the interplanetary magnetic field strength sometimes have fractal or multifractal structure~\\cite{Burlaga-fractals}. This phenomenon has been observed at different distances from Earth and for very different time scales. It is likely that in case that the ``magnetic lens'' discussed above has certain fractal or multifractal properties then it may lead to the appearance of chaotic dynamics in EAS time series.  Certainly, the presented model is only conceptual and oversimplified. Still we hope that it reflects the nature of processes that may lead to the appearance of EAS clusters, both ``ordinary'' and ``possibly chaotic'' ones.  Finally, it is interesting to compare our results with the conclusions of similar investigations performed by other research groups. In a considerable number of articles devoted to the nonlinear time series analysis, one can find a comprehensive investigation of EAS arrival times registered with the EAS-TOP array~\\cite{Aglietta}.  Basing on a detailed study of the available experimental data set and the results obtained with the underground muon monitor~\\cite{Bergamasco92} the authors of this work made a conclusion that cosmic ray signals are all colour random noise, independently of the nature of the secondary particle and of the primary parent particle, but an existence of deterministic chaotic effects in cosmic ray time series cannot be completely excluded. It was also demonstrated in one of the following articles that an impact of background noise brings additional difficulties to the problem of distinguishing between chaotic and stochastic dynamics~\\cite{Bergamasco94}. In our opinion, these conclusions as a whole do not contradict our results, especially in view of the fact that signs of chaotic behaviour have only been observed for about~0.1\\% of EAS in the whole data set.  Besides this, a whole series of investigations devoted to the nonlinear analysis of EAS time series are carried out in Japan beginning from early 1990s at the experimental arrays that now constitute the LAAS network, see, e.g., \\cite{Japan:AP,Japan2001} and references therein. The authors of these investigations presented several dozens of events that demonstrate chaotic dynamics. More than this, it was conjectured that the observed dynamics may be due not only to the chaotic structure of the medium through which particles have traversed but also to the nature of the primary particles~\\cite{Japan:particles}. Later on, there was suggested a model according to which chaotic events may be generated by cosmic rays that have a structure of a fractal wave arriving from a nonlinear accelerator like a supernova remnant~\\cite{Japan:FractalWave}. This model needs to be studied in details, but seems to be promising. Thus that the results obtained during our analysis do not contradict the conclusions of similar investigations performed at other EAS arrays.  The results presented above demonstrate that one can observe an unusual dynamics of EAS arrival times in the vicinity of certain clusters of EAS with the electron number of the order of~$10^5$. While the overwhelming majority of samples in our data set are unambiguously stochastic, a number of samples in the vicinity of four EAS clusters demonstrate signs of chaotic dynamics. Still it is rather difficult to make a final conclusion on the nature of this phenomenon: Does it represent deterministic chaos or a special type of a stochastic process? In our opinion, the majority of the tests performed witness in favour of the first of these two alternatives. Nevertheless we must mention that our analysis may somehow suffer of the fact that the phenomena discussed above are only observed at comparatively short time scales with short samples while time series analysis usually prefers longer samples. In this connection, we recall that our investigation of EAS clusters has revealed an existence of ``super clusters,\" i.e., clusters that have duration more than 30~min and consist of hundreds of EAS. Our future plans include an analysis of these events.  It seems to be necessary to continue the work in this area and to involve some other methods of nonlinear time series analysis. There are a number of other nonlinear tools that may help to make a more definite conclusion about the nature of the observed phenomenon. Among them, one can recall space-time--separation and recurrence plots and the Lyapunov exponents. There are also a number of other measures of nonlinearity besides the one used above~\\cite{Schreiber:PhysRep}. Last but not least, special signal filtering techniques may be used in future to reduce effects of background noise on the dynamics of PC samples.  Finally, as we have already mentioned above, perhaps the biggest puzzle in connection with signs of chaos in EAS time series is their astrophysical nature. It is likely that clusters that produce signs of chaotic dynamics in the corresponding time series are similar to the upper part of an iceberg in a sense that they do not present the complete process but only the most pronounced part of it. We point out that for all PC samples discussed above, the value of the correlation dimension is comparatively small ($D_2<3$). Since this value gives a lower estimate for the number of degrees of freedom in the underlying process, it is possible that the structure of this process is not too complicated. Still it seems to be a great challenge to work out a good model that could explain chaotic dynamics in EAS arrival times.   \\begin{ack} We gratefully acknowledge numerous useful discussions with A.~V.~Igoshin, A.~V.~Shirokov, and V.~P.~Sulakov who have helped us a lot with the experimental data set. We also thank Thomas Schreiber for a very helpful communication and anonymous referees for stimulating comments. This work was done with financial support of the Federal Scientific-Technical Program ``Research and design in the most important directions of science and techniques\" for 2002--2006, contract No.\\ 40.014.1.1.1110, and by Russian Foundation for Basic Research grant No.\\ 02-02-16081.  Only free, open source software was used for this investigation. \\end{ack}"
    },
    "0207/astro-ph0207473_arXiv.txt": {
        "abstract": "We present the discovery of OB type absorption lines superimposed to the emission line spectrum, and the first double-lined orbital elements for the massive Wolf-Rayet binary HDE~318016 (=WR~98), a spectroscopic binary in a circular orbit with a period of 47.825 days. The semiamplitudes of the orbital motion of the emission lines differ from line to line, indicating mass ratios between 1 and 1.7 for $\\mathcal{M}_{WR}/ \\mathcal{M}_{OB}$. ",
        "introduction": "HDE~318016 =WR~98 is one of the few stars with Wolf-Rayet spectrum showing both N and C emission lines. This star was found to be a single-lined binary with a period of 47.8 days with N and C emission lines moving in phase (Niemela 1991). Relevant parameters of WR~98 can be found in the recent VIIth Catalogue of galactic WR stars (van der Hucht  2001).  Here we present the results of a detailed radial velocity analysis of optical spectral lines of WR~98 showing it to be a double-lined binary.  We have obtained a total of 69 blue optical spectrograms of WR~98 at two different Observatories, namely CTIO in Chile (1980-84) and CASLEO (1997--2001) in Argentina. The observations at CTIO were performed with the Cassegrain IT spectrograph on the 1m Yale telescope, using photographic plates with fine grain emulsion as detector. The observations at CASLEO where obtained with the Cassegrain spectrographs with CCD detectors attached to the 2.15m telescope. ({\\footnotesize CASLEO is operated under agreement between CONICET, SeCyT, and the Universities of La Plata, C\\'ordoba and San Juan, Argentina.})  The photographic observations were digitized with a Grant engine at La Plata. The digital data were processed and measured with IRAF routines. ",
        "conclusions": ""
    },
    "0207/astro-ph0207535_arXiv.txt": {
        "abstract": "We present a simple phenomenological model of feedback in early-type galaxies that tracks the evolution of the interstellar medium gas mass, metallicity, and temperature.  Modeling the star formation rate as a Schmidt law with a temperature-dependent efficiency, we find that intermittent episodes of star formation are common in moderate-size ellipticals.  Our model is applicable in the case in which the thermalization time from SN is sufficiently long that spatial variations are relatively unimportant, an appropriate assumption for the empirical parameters adopted here, but one that can only be demonstrated conclusively though more detailed numerical studies. The departure from a standard scenario of passive evolution implies significantly younger luminosity-weighted ages for the stellar populations of low-mass galaxies at moderate redshifts, even though the more physically meaningful mass-weighted ages are changed only slightly.  Secondary bursts of star formation also lead to a natural explanation of the large scatter in the NUV-optical relation observed in clusters at moderate redshift and account for the population of E+A galaxies that display a spheroidal morphology.  As the late-time formation of stars in our model is due to the gradual cooling of the interstellar medium, which is heated to temperatures $\\sim$ 1 keV by the initial burst of supernovae, our conclusions do not rely on any environmental effects or external mechanisms.  Furthermore, a simple estimate of the X-ray emission from this supernova heated gas leads to an $L_X$ vs $L_B$ correlation that is in good agreement with observed values.  Thus feedback processes may be essential to understanding the observed properties of early-type galaxies from the optical to the X-ray. ",
        "introduction": "Of all the fields of astronomy, the study of galaxy formation may be the one in which the rift between theory and observation is the most difficult to bridge.  Theoretical astrophysicists, perhaps from a temperamental as well as a practical point of view, are best at predicting the most difficult component of the universe to measure: the overall distribution of invisible ``dark matter.'' While the evolution of this component is well understood, numerical and analytical models must stretch their predictive powers to the limit to superimpose on this history the messy heating, cooling, and enrichment processes that affect baryonic gas.  At the same time, observational astrophysicists, perhaps from a romantic as well as practical point of view, are happiest when looking at the most complicated thing they could ever measure: the stars in the night sky.  A proverbial tail that wags the dog, the magnitudes and colors of the stellar populations in galaxies are dependent not only on the history of the baryonic gas, but on the notoriously complicated process of star-formation.  As many possible environmental effects can have an unknown impact on the number and distribution of stars formed (See, eg. Kroupa 2001; Larson 1999; Scalo 1998) which is even more uncertain under primordial conditions (see, eg. Nakamura \\& Umemura 2001), predicting the optical properties of galaxies involved from ``first principles'' requires an enormous amount of extrapolation and simplifying assumptions.  This great rift calls for astronomers to work to meet each other from each side of the divide.  Thus simulations and semi-analytical models of galaxy formation are continuously striving to include all the important physical processes that affect forming galaxies and the surrounding intergalactic medium (e.g.  Kauffmann et al.\\ 1999; Baugh et al.\\ 1998; Somerville \\& Primack 1999).  From the observational point of view, astronomers must constantly develop more careful phenomenological models, which better express the features seen in simulations: making statements as to the ages of stars, their metallicities, and levels of dust extinction, rather than simply considering their infrared, optical, and ultraviolet colors and magnitudes.  Models of this sort include Tinsley (1980), Matteucci \\& Tornamb\\'e (1987), Tantalo et al.\\ (1998), and Ferreras \\& Silk (2000b).  From a theoretical point of view one of the most important issues to recently come to the fore is the role of supernova feedback in galaxy formation.  N-body simulations have shown that the baryonic components of galaxies that form in Cold Dark Matter (CDM) simulations systematically lose too much of their angular momentum to the dark matter halos in which they are contained (Navarro \\& Benz 1991).  As discussed in White (1994), the resolution of this problem is widely believed to be the inclusion of feedback, although the proper modeling of this process remains a subject of intense debate.  Thus since the first numerical model of thermal heating by Katz (1992), approaches have ranged from implementation of kinetic boosts (Milos \\& Hernquist 1994; Gerritsen 1997), a multiphase model of star formation and feedback (Yepes 1997; Hultman \\& Pharasyn 1999), a model of turbulent pressure support from feedback (Springel 2000), and a model in which energy persists in the interstellar medium for a period corresponding to the lifetime of stellar associations (Thacker \\& Couchman 2001). Finally, feedback processes are likely to play a more general role in impacting the intergalactic medium and conditions under which galaxy formation occurs (Mac~Low \\& Ferrara 1999; Martel \\& Shapiro 2001; Scannapieco \\& Broadhurst 2001; Scannapieco, Thacker, \\& Davis 2001).  Yet despite the many exploratory feedback models being examined theoretically, the role of feedback in empirical modeling has been more limited.  Ferreras \\& Silk (2000a; 2001) explored a simple model in which feedback was included phenomenologically using two free parameters: the star formation efficiency and the fraction of gas ejected in outflows. These values are dependent on the thermal and mechanical feedback from supernovae but they fail to capture their impact the state of the interstellar medium (ISM) itself, despite the close relationship between this state and star formation (McKee \\& Ostriker 1977).  Our primarily theoretical motivation in this paper, then, is to extend these models by tracking the temperature evolution of the ISM, taking into account the thermal contributions of supernovae and metallicity-dependent gas cooling.  The primarily observational motivation for this paper, on the other hand, is the measurement of a large scatter at the faint end of the near-ultraviolet (NUV) minus optical color-magnitude relation in cluster Abell~851 ($z=0.4$) observed by Ferreras \\& Silk (2000a) using passbands F300W and F702W of the WFPC2 on board the Hubble Space Telescope.  A simple analysis of this scatter shows that it can not be accounted for by old Horizontal Branch stars.  Moreover, significantly younger luminosity-weighted ages were found, although a degeneracy between the age of the young stellar component and its mass fraction prevented a detailed estimate of a more physically meaningful mass-weighted age. This result shows that the star formation history of early-type galaxies is much more complicated than the standard picture in which a single stellar population formed at a redshift $z\\gtrsim 3$ and then evolved passively.  The structure of this work is as follows.  In \\S2 we describe a simple thermal model of supernova feedback in elliptical galaxies and its impact on the gas, stellar, and metallicity evolution of these objects.  In \\S3 we describe the general features of our model, and the physics on which they depend.  Our evolutionary tracks are combined with population-synthesis models in \\S4 and used to study the color-magnitude relationship in ellipticals in both the optical and near-ultraviolet.  In \\S5 we examine the compatibility of our model with X-ray observations of early-type galaxies.  In \\S6 we discuss the implications of our modeling for future observational studies of feedback in ellipticals as well as how our model compares to other theoretical approaches, and conclusions are given in \\S7. ",
        "conclusions": "While stellar feedback in galaxy formation has been the subject of intense theoretical investigation, these studies have had little impact on the direct interpretation of the observed properties of elliptical galaxies.  Yet there are several observational clues that point to its importance. Ferreras \\& Silk (2000b) for example, found that the color-magnitude relation of early-type cluster galaxies can be easily understood in the context of a variable ejection efficiency of supernova material, which scales inversely with galaxy mass. Similarly, the large scatter in the faint end of the near-ultraviolet minus optical color-magnitude relation in early-type cluster galaxies hints at episodic star-formation in these objects, a natural consequence of stellar feedback (Ferreras \\& Silk 2000a).  Motivated by these observational and theoretical pointers, we have explored in this work a simple model of thermal feedback in elliptical galaxies.  Our model consists of only gas and stars in a single zone and is a natural extension of previous investigations (Tinsely 1980; Ferreras \\& Silk 2000b; Ferreras \\& Silk 2001) which includes the thermal state of the interstellar medium.  We account for heating by supernova feedback and line emission cooling, and their impact on the temperature and density of the ISM using a minimum of parameters, and relate these quantities to the overall star formation rate using a Schmidt law and a temperature-dependent efficiency.  This model is applicable in the case in which the time for SNe to percolate through the ISM is sufficiently long that spatial variations between regions are relatively unimportant.  While this is the case for the choice of parameters adopted in this paper, more detailed modeling is necessary to study the interplay between thermalization and substructure conclusively.  Our model leads to two important and independent results.  First the interplay between infall and variable SNe ejection efficiency provides a natural explanation of the optical and NIR color-magnitude relations in elliptical galaxies both locally and at moderate redshift ($z\\lesssim 1$).  Our model generates a mass-metallicity relation $Z\\propto M^{0.15-0.2}$, which accounts for the observed color range over three magnitudes in luminosity. Furthermore, the rapid cooling times for temperatures close to $10^4$~K result in short bursts of star formation, generating stellar populations that are very similar to the simple ones commonly used to describe ellipticals.  This implies no change in the slope or the scatter of the optical and NIR color-magnitude relation out to redshifts $z\\lesssim 1$ as observed by many authors (e.g.  Stanford et al.\\ 1998; Van~Dokkum et al.\\ 1998; Van~Dokkum et al.\\ 2000)  Secondly, we find that secondary peaks of late star-formation are ubiquitous in smaller systems in which the fraction of the SNe ejecta that escapes from the ISM, $\\bout$, is significant. In these cases the gas is heated to temperatures $\\propto (1-\\bout)$, leading to secondary bursts of star formation with a delay that scales as $\\propto (1-\\bout)^3$ due to the $T^3$ scaling of the cooling times of high temperature gas.  As such bursts occur within a Hubble time only in the smaller galaxies in which $\\bout$ is relatively large, this leads to a natural explanation of the large scatter in the NUV-optical relation as observed in clusters at moderate redshifts.  This can also account for the observed population of post-starburst E+A galaxies that display a spheroidal morphology (Ferreras \\& Silk 2000c). While the current claim for these post-starburst systems is the quenching of star formation while falling through the hot intracluster medium (Poggianti et al.\\ 1999) our model shows that late starbursts may arise in elliptical galaxies without resorting to environmental mechanisms.  In fact, we have found that no such oscillations arise in a broad class of theoretical models which study the interplay SNe and the further accretion of gas.  Given that the duration of these peaks of star formation is rather short ($\\sim 100$~Myr), the most promising observational approach is to examine the NUV properties of spheroidal galaxies.  A comprehensive study of the rest frame NUV$-$optical color-magnitude relation in ellipticals may be able to quantify both the number of that undergo late bursts and the mass fraction in young stars.  Furthermore, as the thermal mechanism studied this paper is independent of environment, comparisons of such colors between field, group, and cluster populations can help to differentiate between this feedback and other processes.  As the contribution in the NUV from evolved core helium-burning stars becomes significant with age, studies of galaxies at redshifts $z\\gtrsim 0.3$ may provide the cleanest samples for such comparisons.  At $z\\gtrsim 1$, late bursts can be significant even in high-mass galaxies, as shown in Figure~1.  This may result in the failure of any search for high-redshift ellipticals that is based on simple passively evolving models and does not invoke complicated scenarios of assembly (e.g. Zepf 1997).  Our model also underscores the bias intrinsic to observing luminosity weighted quantities. While the mass and luminosity weighted ages of higher mass ellipticals are quite similar, even the relatively small bursts of late star formation that arise in smaller objects cause large changes in their observed $V$ band luminosity-weighted ages. This effect is even more severe in the NUV, highlighting the importance accounting for the overall star-formation history when interpreting observed stellar ages.  One of the remarkable results of our model is the presence of an interstellar medium that has been heated by supernovae to temperatures around $1$~keV or higher. This hot gas exists at very low densities, with only a fraction leaving the galaxy, much in the same way as the hot gas of the solar corona is kept gravitating around the Sun.  A full 3D simulation would be needed in order to explore the real conditions under which this ``fountain effect'' can occur.  Finally our prediction of a hot corona of SNe heated gas has natural implications for the X-ray properties of elliptical galaxies.  By approximating the overall X-ray luminosity of each object as due to Bremsstrahlung emission from gas at the single temperature and density given by our model, we derive a $L_X$ vs $L_B$ correlation that is in good agreement with observed values.  Nevertheless our simple one-zone model does not account for environmental effects and ignores what is likely to be a significant contribution from unresolved X-ray binaries, resulting in a scatter that is smaller than observed.  Galaxy formation is one of the richest and most complex processes in all of astrophysics, forcing observers and theorists to approach it from completely different viewpoints.  And although both theoretical and observational progress has been clear and systematic, the rift between theory and observation in the field remains one of the most difficult to bridge.  In this work, we have attempted to draw the key theoretical issue of feedback into the direct interpretation of the observational properties of early-type galaxies.  While our approach has been explorational, the widespread agreement of our simple model with diverse observations suggests that thermal feedback processes are likely to be essential to fully understanding the optical, ultraviolet, and X-ray properties of early-type galaxies."
    },
    "0207/astro-ph0207253_arXiv.txt": {
        "abstract": "The diffuse extragalactic gamma-ray background (EGRB) above 100 MeV encodes unique information about high-energy processes in the universe.  Numerous sources for the EGRB have been proposed, but the two systems which are certain to make some contribution are active galaxies (blazars) as well as normal galaxies.  In this paper, we evaluate the contribution to the background from both sources. The active galaxy contribution arises from unresolved blazars. We compute this contribution using the Stecker-Salamon model. For normal galaxies, the emission is due to cosmic-ray interactions with diffuse gas.  Our key assumption here is that the cosmic-ray flux in a galaxy is proportional to the supernova rate and thus the massive star formation rate, quantified observationally by the cosmic star formation rate (CSFR). In addition, the existence of stars today requires a considerably higher ISM mass in the past. Using the CSFR to compute both these effects, we find that normal galaxies are responsible for a significant portion ($\\sim 1/3$) of the EGRB near 1 GeV, but make a smaller contribution at other energies. Finally, we present a ``minimal'' two-component model which includes contributions from both normal galaxies and blazars.  We show that the spectrum of the diffuse radiation is a key constraint on this model: while neither the blazar spectra, nor the galactic spectra, are separately optimal fits to the observed spectrum, the combined emission provides an excellent fit. We close by noting key observational tests of this two-component model, which can be probed by future gamma-ray observatories such as GLAST. ",
        "introduction": "All-sky $\\gamma$-ray observations by SAS 2 \\citep{fichtel,fichtel2} and most recently by EGRET \\citep{sreek_obs} have revealed the existence of an isotropic diffuse $\\gamma-$ray emission, presumably of extragalactic origin. This extragalactic gamma-ray background (EGRB) is well-described by a power-law energy spectrum with an index of $-2.10 \\pm 0.03$, while the extragalactic intensity for energies $> 100 \\ {\\rm MeV}$ has an all-sky average value $(1.45 \\pm 0.05) \\times 10^{-5} \\iun $ \\citep{sreek_obs,fichtel96}.  A variety of possible contributions to the EGRB have been proposed. There are, however, two classes of $\\gamma-$ray sources whose existence has been observationally established and thus guarantees that these make {\\em some} contribution to the EGRB: blazars and normal galaxies.  Blazars, which are active galactic nuclei with jets almost aligned with the line of sight \\citep{bk79,mgc92,dg95}, comprise the vast majority of the  identified  $\\gamma-$ray point sources detected by EGRET \\citep{egret3}. In addition, their energy spectra are power laws, with a distribution of indices peaked close to the EGRB energy spectrum index. It is therefore only logical to argue that a population of unresolved blazars with photon fluxes below EGRET sensitivity has to be the origin of a significant portion of the EGRB \\citep{ssm}. A large AGN contribution to the EGRB has been anticipated as early as the 1970's (e.g., \\citet{sww,bfht,pgfc}). Given the EGRET results on blazars, \\citet{ss94} and \\citet[henceforward SS96]{ss96} made a detailed calculation of the blazar contribution to the EGRB and indeed found it to be dominant, although the shape of the predicted blazar emission energy spectrum does not match the flatter spectrum of the latest EGRET EGRB determination \\citep{sreek_obs}.  While the EGRET catalog of point sources is dominated by blazars, the EGRET diffuse flux is dominated by emission from the Galactic plane. The latter is, for the most part, the result of the decay of neutral pions produced when cosmic rays interact with the interstellar medium. The superposition of this diffuse $\\gamma-$ray emission from all unresolved normal galaxies is the second guaranteed source of extragalactic background $\\gamma-$ray intensity. The contribution of normal galaxies to the EGRB was calculated by \\citet{sww} for the case of non-evolving galaxies and was found to be a few percent of the observed background. \\citet{lbp} extended this calculation to include galactic evolution, which they found to significantly enhance the predicted background. Their results spanned a range $\\Phi(> 100 \\ {\\rm MeV}) = 0.3 - 7 \\times 10^{-6} \\ \\iun$ which comes much closer to the observed level. \\citet{schramm} and \\citet{prantzos} inferred a similarly large result from spallogenic element abundances.   {\\it We define the sum of the $\\gamma-$ray emission from all unresolved blazars and from all unresolved normal galaxies to be the guaranteed EGRB.} If the intensity level of the guaranteed EGRB can be confidently estimated, then by comparison to the observed EGRB one can constrain the observationally allowed contributions from any other hypothesized sources.  In this paper we present a new calculation of the contribution of normal galaxies to the EGRB. We use observational estimates of the cosmic star formation rate (CSFR), which have recently become available, to model the evolution of normal galaxy $\\gamma-$ray emission. To the latter, we then add the blazar component of the spectrum as given by SS96. Our results are computed for the currently favored $\\Omega_\\Lambda=0.7$, $\\OmegaM=0.3$ cosmology. The resulting minimal two-component model will prove to be an excellent fit to the observed EGRET EGRB spectrum for energies up to 15 GeV, where $\\gamma-$ray extinction is not important. ",
        "conclusions": "The minimal 2-component model of the EGRB can be tested and improved in various ways when observations from future $\\gamma$ - ray telescopes such as GLAST become available. On the one hand, the improved effective collecting area of GLAST as compared with that of EGRET, especially at high energies, will allow GLAST to accurately measure the shape of the EGRB for energies up to 1 TeV. This will test whether the spectrum turns over at energies higher than $\\sim 15$ GeV due to electron-pair production via interactions with the IR, UV and optical backgrounds \\citep{mph96,ss98}. Any model in which the EGRB is assumed to be of cosmic origin should exhibit this behavior. If the GLAST data do not confirm that prediction, all extragalactic models are ruled out unless they can predict compensatingly steepening spectra with increasing redshift.  On the other hand, the improved point source sensitivity of GLAST will allow it to resolve a higher number of blazars ($\\sim $ 100 more than EGRET, Stecker \\& Salamon 1999), and therefore the blazar contribution to the EGRB will be reduced by about a factor of 2. If unresolved blazars are the only constituent of the EGRB, the {\\it fractional change} of the EGRB  will be the same as the fractional change of the background blazar emission. If, however, there is a second component in the EGRB (in our case, that of normal galaxies), the fractional change of the EGRB should be smaller.  In addition, with the blazar component reduced by a factor of 2, our calculated normal galaxy contribution will become comparable to that of blazars for energies $\\sim 1$GeV. Therefore, if the relative contributions of blazars and normal galaxies to the minimal model are comparable to our estimates, the shape of the EGRB spectrum should start to exhibit a (detectable in principle) deviation from its single power-law form at $\\sim 1$ GeV, corresponding to the normal galaxy spectrum peak. Were this peak detected, the relative contribution of normal galaxies to the EGRB could be determined observationally.  The observations of GLAST can also be used to improve the minimal model and its predictions. The observations of more blazars will allow a more confident determination of the observational inputs for the SS96 blazar model, as pointed out by \\citet{ss99}. As far as the normal galaxy  component model is concerned, GLAST is expected to detect several Local Group galaxies (the SMC, LMC, M31 and maybe M33; Pavlidou \\& Fields 2001), and therefore it will be possible to check our assumption of the universality of the galactic diffuse gamma-ray emission spectrum.  Finally, with both guaranteed EGRB components well-understood, one can better identify or constrain any other components and any new physics which might generate them."
    },
    "0207/astro-ph0207586_arXiv.txt": {
        "abstract": "We present position-velocity strip maps of the Galactic Center region in the CO $J=7 \\rightarrow 6$ and $J=4 \\rightarrow 3$ transitions observed with the Antarctic Submillimeter Telescope and Remote Observatory (AST/RO) located at Amundsen-Scott South Pole Station. Emission from the two rotational transitions of $^{12}$CO was mapped at $b=0^{\\circ}$ for $3.5^{\\circ}> \\ell > -1.5^{\\circ}$, on a $1'$ grid with a FWHM beamsize of $58''$ at 806 GHz and $105''$ at 461 GHz. Previous observations of CO $J=4\\rightarrow 3$ (Martin et al., in preparation) and [C\\,I] (Ojha et al. 2001) emission from this region show that these lines are distributed in a manner similar to CO $J=1 \\rightarrow 0$ (Stark et al. 1987); the (CO $J=4 \\rightarrow 3$)/(CO $J=1 \\rightarrow 0$) line ratio map is almost featureless across the entire Galactic Center region. In contrast, the CO $J=7 \\rightarrow 6$ emission from the Galactic Center is strongly peaked toward the Sgr~A and Sgr~B molecular complexes. A Large Velocity Gradient (LVG) analysis shows that aside from the two special regions Sgr~A and Sgr~B, the photon-dominated regions within a few hundred parsecs of the Galactic Center are remarkably uniform in mean density and kinetic temperature at $n = 2500$ to $4000 \\, \\mathrm{cm^{-3}}$ and T = 30 to 45 K. The (CO $J=7 \\rightarrow 6$)/(CO $J=4 \\rightarrow 3$) line temperature ratios near Sgr~B are a factor of two higher than those observed in the nuclear region of the starburst galaxy M82 (Mao et al. 2000), while the CO($J=7 \\rightarrow 6$)/CO($J=4 \\rightarrow 3$) line temperature ratios around Sgr~A are similar to M82. The line ratio on large scales from the Galactic Center region is an order of magnitude less than that from M82. ",
        "introduction": "\\label{s:intro}  Observations of photons emitted by the various rotational transitions of the ground vibrational state of carbon monoxide (CO) are the primary means of studying molecular gas in the Galaxy. These spectral lines occur at frequencies of $J \\times 115 \\, {\\mathrm{GHz}}$, for the transition from the $J$ to the $J-1$ rotational state. Numerous galactic surveys (Combes 1991) have studied the lowest-frequency $J = 1 \\rightarrow 0$ line. The brightness of this spectral line is roughly proportional to the total molecular column density within the telescope beam (Liszt 1984), provided that the molecular gas has a relatively low density ($n < 3 \\times 10^3 \\,{\\mathrm{cm}^{-3}}$) and column density ($N < 10^{23} \\, {\\mathrm{cm}^{-2}}$).  Surveys of the $J = 1 \\rightarrow 0$ line therefore give an indication of the extent and distribution of molecular material.  A more complete picture of the thermodynamic state of the molecular gas can be gained through observations of other, higher-$J$ transitions of CO and the various rotational transitions of the isotopically-substituted species $^{13}$CO and C$^{18}$O. If the brightness of several of these lines is known, models of the excitation and radiative transfer can be used to solve for the density, temperature, and cooling rate of the molecular gas. To avoid degeneracy in this solution, it is valuable to have observations of a line transition from a $J$-state that is sufficiently high in energy such that it is only weakly populated. CO in interstellar molecular gas is typically in a thermodynamic state where the low-$J$ transitions are in approximate thermal equilibrium, so that all of the rotational transitions below some $J$-level have roughly the same excitation temperature as the $J = 1 \\rightarrow 0$ transition. The observable quantity that is usually compared with theory is the ratio of line brightnesses. If only low-$J$ lines are observed, it is found that the ratio of these line brightnesses are only weakly dependent on excitation temperature and that maps of all the low-$J$ transitions look similar, regardless of variations in excitation temperature across the map. If, however, we observe a transition from a high-$J$ level that is not strongly excited, it will be found that the line brightness ratios involving that transition do vary significantly across the map, and the value of the excitation temperature can be more readily determined.  Previous observations with the AST/RO telescope (see Figure 1) have shown that the distribution of the  CO $J = 4 \\rightarrow 3$ line emission is remarkably similar to that of the CO $J = 1 \\rightarrow 0$ line throughout the Galactic Center region. Previous observations of CO  $J = 7\\rightarrow 6$ (Harris et al. 1985) showed that line to be unusually bright within $\\pm 200''$ of the Galactic Center.  In this paper, we report observations of the CO $J = 7 \\rightarrow 6$ and $J = 4\\rightarrow 3$ transitions in a strip map several degrees in extent, and find that the $J = 7 \\rightarrow 6$ line has significantly lower brightness temperature and significantly different spatial distribution than the lower-$J$ CO lines. We can therefore produce a new model of the thermodynamic state of the CO gas along this 1-dimensional strip. ",
        "conclusions": "We have presented maps of the Galactic Center region in the CO $J=7 \\rightarrow 6$ and $J=4 \\rightarrow 3$ transitions observed with the Antarctic Submillimeter Telescope and Remote Observatory (AST/RO) located at Amundsen-Scott South Pole Station. Comparing the CO $J=4 \\rightarrow 3$, CO $J=1 \\rightarrow 0$, and $^{13}$CO $J=1 \\rightarrow 0$ maps reveals that the CO $J=7 \\rightarrow 6$ emission is concentrated toward the Sgr~A and Sgr~B complexes. Using an LVG model, we find that the gas kinetic temperatures of the Sgr~B and Sgr~A complexes are 72 K and 47 K and the gas densities of those regions are 10$^{4.15}$ cm$^{-3}$ and 10$^{3.9}$ cm$^{-3}$. Excluding those regions, we estimate the mean kinetic temperature and density throughout the Galactic Center are 30 to 45 K and 10$^{3.7\\pm0.05}$ cm$^{-3}$. The CO $J =7 \\rightarrow 6$ transition accounts for 7.5\\% of the total CO luminosity from the Sgr~A complex. At the Sgr~B complex, with its higher gas temperature, the CO $J =7 \\rightarrow 6$ transition contributes more than 17\\% of the total CO luminosity.  The tragic death of Rodney Marks, who was the year 2000 AST/RO Winterover Scientist, occurred just before this observational program. The remaining winterover crew from the Center for Astrophysical Research in Antarctica---Gene Davidson, Greg Griffin, David Pernic, and John Yamasaki---continued AST/RO operations in tribute to Rodney's memory, allowing these observations to be made. The AST/RO group is grateful for the logistical support of the National Science Foundation (NSF), Antarctic Support Associates, Raytheon Polar Services Company, and the Center for Astrophysical Research in Antarctica during our polar expeditions. SK thanks Min Yan for LVG analysis and Wilfred Walsh for his helpful comments on the manuscript. This work was supported in part by  United States National Science Foundation grant DPP88-18384, and by the Center for Astrophysical Research in Antarctica and the NSF under Cooperative Agreement OPP89-20223."
    },
    "0207/astro-ph0207065_arXiv.txt": {
        "abstract": "{ There exists both theoretical and observational evidence that the magnetic field decay in neutron stars may proceed in a pronounced non--linear way during a certain episode of the neutron star's life. In the presence of a strong magnetic field the Hall--drift dominates the field evolution in the crust and/or the superfluid core of neutron stars. Analysing observations of $P$ and $\\dot{P}$ for sufficiently old isolated pulsars we gain strong hints for a significantly non--linear magnetic field decay. Under certain conditions with respect to the geometry and strength of a large--scale magnetic background field an instability is shown to occur which rapidly raises small--scale magnetic field modes. Their growth rates increase with the background field strength and may reach $\\sim 10^4$ times the ohmic decay rate. Consequences for the rotational and thermal evolution as well as for the cracking of the crust of neutron stars are discussed. ",
        "introduction": "There is still an ongoing scientific debate whether the magnetic fields of pulsars decay at all \\citep[see, e.g.,][ and references therein]{RF01}. Thus, the question whether this questionable decay proceeds in a non--linear way seems to be somewhat academic. However, the neutron star magnetic fields are the strongest known in the universe and already this very fact suggests the idea that, if there is a field evolution at all, it should be a non--linear one. As we will show, there exist both theoretical and observational evidences that the field decay proceeds essentially in a non--linear way, at least in certain regions of the star and surely only during certain episodes of its life.  In order to be able to deal with well defined physical conditions we confine ourselves here to considering the evolution of {\\em sufficiently old isolated neutron stars} (SOINSs). We characterize them by the following properties: \\begin{itemize} \\item the supernova fall--back accretion phase is finished and a stable density stratification $\\rho(r)$ has been established; \\item the rediffusion of the magnetic field has been completed; \\item  the rotation is slow enough so that the irradiation of gravitational waves does not contribute to the spin down; \\item the appearance of glitches is less probable; \\item the crust is almost completely crystallized; \\item the temperature profile has become flat enough, to avert convection in the outermost thin liquid layer and in the core region as well as the occurrence of a thermoelectric instability in the crust; \\item the electron relaxation time is already so large that the magnetization parameter may reach values large enough to make the electron transport processes non--linear; \\item the absence of accretion, as a consequence of isolation, ensures that there are neither external sources of angular momentum nor of heating; \\item therefore, there is a stable compactness ratio $M/R$ and no screening of the magnetic field can happen. \\end{itemize} All these constraints are met by the majority of neutron stars the age of which is $\\gtrsim 10^5$ years. Therefore, if any thermal and/or rotational evolution of SOINSs is observed beyond that predicted when assuming a constant magnetic field, it is with a high probability caused by an evolving one.  Here, we intend to consider as the non--linearity of the field evolution the \\emph{Hall--effect} or \\emph{Hall--drift}, as it occurs in the crust \\citep[see, e.g.,][]{SU97}, i.e., both ambipolar diffusion and any convective motion of the neutron star matter is excluded.\\\\ The effect of the Hall--drift on the magnetic field evolution of isolated neutron stars has been considered by a number of authors \\citep[see, e.g.,][]{HUY90,GR92,M94,NK94,SU95,SU97,US95,US99,VCO00}. They discussed the redistribution of magnetic energy from an initially large--scale (e.g., dipolar) field into small--scale components via the Hall--term. Though the Hall--drift itself is a non--dissipative process, these changes in the field geometry may in principle accelerate the field decay. However, the results presented in \\citet{NK94}, \\citet{M94}, \\citet{SU97}, and \\citet{US99} suggest the conclusion that the field decay is not modified drastically.  Goldreich \\& Reisenegger \\citet{GR92} developed the idea of the {\\em Hall--cascade}, i.e., when starting with a large--scale magnetic field small--scale field components are generated down to a scalelength $l_{\\text{crit}}$, where the ohmic dissipation begins to dominate the Hall--drift.  In some of the above--mentioned investigations numerical instabilities are mentioned if either the field structure becomes too complex \\citep{US95} or the initial field is too strong \\citep{NK94,US99}. Also, when considering the thermomagnetic field generation in the crust of young neutron stars \\citep{WG96}, where small--scale modes are the first ones to be excited, numerical instabilities occurred exclusively caused by the Hall--drift.  Recently we have shown \\citep{RG02} that all the observed instabilities are in their essence very likely not of numerical origin but have physical reasons: A sufficiently strong and inhomogeneous large--scale background field is unstable with respect to small--scale perturbations, which rise very rapidly in comparison with the relevant ohmic decay time of the background field. This rapid transfer of energy from the background into small--scale field components proceeds \\emph{not} via a cascade but jump--like across wide spectral distances. The unstable perturbations show small radial structures close to the neutron star's surface whereas their lateral structures are of medium size. Of course, this Hall--instability can act only during an episode of the field decay, which unavoidably leads to a zero field. This episode, however, may have observable consequences.  In the following we will present the theoretical and observational hints for the non--linear field decay in isolated neutron stars. After a description of the Hall--drift induced field instability we discuss its effects on observable quantities of the above introduced SOINSs. ",
        "conclusions": "Clearly, any assignment of the results gained by help of a very simplified model to real neutron stars has to be done with great care. Even when accepting the plane layer as a reasonable approximation of the crust one has to concede that the very specific profile of $\\Bvec_0$ assumed above may only exemplify the crustal field's structure.  An acceptable approximation of the radial profile of a dipolar crustal field as given, e.g., in \\citet{PGZ00} will in general have to allow for non--zero values at the boundaries which cause a more complicated perfect conductor boundary condition. Moreover, the strong dependence of the conductive properties on the radial co--ordinate should anyway be taken into account.  Recently, \\citet{HR02} solved the non--linear \\Eqref{equ:indeqdimless} for the first time in a spherical shell geometry, again with the transport coefficients constant in space and time, and applying the vacuum  and the perfect conductor  boundary conditions at the surface and at the crust bottom, respectively. This is the best model for the non--linear magnetic field evolution in the crust of neutron stars available so far. They made different assumptions about the large--scale initial field (which corresponds in some sense to our background field) and report, as all the formerly mentioned authors too, about insurmountable numerical problems when the initial magnetization parameter, very roughly to be identified with $B_0$, exceeds $\\approx 200$. Although their investigations can not reveal the existence and the character of the instability presented above (too short integration period in comparison with the growth time of the unstable modes for the maximum initial field considered), they found at least two features which support our results: For the largest feasible initial field their Legendre spectra are no longer convergent and indicate a local minimum at $l\\approx 60$, a hint that small--scale field structures are fed in a non--cascade--like, possibly unstable way. Moreover, their findings confirm the importance of the background field's profile curvature: For an initial (poloidal) field nearly linear in $r$ the excitation of small--scale structures and the acceleration of the decay were insignificant, in contrast to an initial (toroidal) field with a nearly quadratic $r$--profile.  To get an impression of possible consequences for the evolution of neutron stars we now simply assume, that the real $\\Bvec$--profile is sufficiently ``curved\" (i.e., its second spatial derivative is big enough) and associate the parameter $B_0$ with a typical value of the field.  Assuming further electric conductivity and chemical composition to be constant, $\\sigma_0= 5\\times 10^{26}$s$^{-1}$ and the relative atomic weight $A/Z=25$, respectively, we find the normalization field at a density $\\rho = 10^{14}$g cm$^{-3}$ to be $7\\times 10^{10}$G \\citep[see, e.g.,][]{PGZ00}. That is, for typical (inner) crustal magnetic fields ranging from $7\\times 10^{12}$G to $1.4\\times 10^{14}$G we find a $B_0$ between $100$ and $2000$ and the e--folding time of the most rapidly growing unstable mode to be $0.0035$ and $0.0003$ times the ohmic decay time, respectively. Note, when comparing with the values of Fig.~\\ref{fig2} that the dimensionless thickness 2 of the slab causes a factor 4 in the growth times. These values agree roughly with the ratios $\\tau_B/\\tau_{\\text{ohm}}$ derived in Sect.~\\ref{obsmot} from observational data. An initial perturbation will quickly evolve to a level at which the linear analysis is no longer feasible, that is, at which it starts to drain a remarkable amount of energy out of the background field. Note, that the assumption about $B_0$ may well be in agreement with the surface field data given in Table~\\ref{table}.  We want to emphasize again that a sufficient curvature of the background field profile is a necessary condition for the occurrence of an unstable behavior. Therefore neither the derivation of the well--known helicoidal waves (whistlers) nor its modification presented in \\citet{VCO00} could reveal it because a homogeneous background field was used there.  \\begin{figure}[t] \\begin{center} \\hspace*{-5mm}\\epsfig{file=MS26038.ps, width=1.1\\linewidth,clip=} \\end{center}  \\vspace{-1.2cm} \\begin{center} \\hspace*{-5mm}\\epsfig{file=MS26039.ps, width=1.1\\linewidth,clip=} \\end{center}  \\vspace{-.3cm} \\caption{Heat source density perturbation $2\\curl\\Bvec_0\\cdot\\curl\\bvec$ (upper panel) and Lorentz force density perturbation $\\curl\\Bvec_0\\times\\bvec + \\curl\\bvec\\times\\Bvec_0$ (lower panel) of the fastest growing mode for $B_0=2000$ in arbitrary units. Upper panel: dark and light shading denotes cooling and heating by the perturbation, respectively; note the stretched $z$--scale. Lower panel: arrows denote $x$-- and $z$--force components, grey shading the value of the $y$--component (dark -- into, light -- out of the plane).} \\label{fig7} \\end{figure}  With great reserve we may speculate about possible observational consequences. For SOINSs, the small--scale field modes initially generated or existing in the crust have already been decayed and the magnetic field is concentrated almost completely in the large--scale, say, dipolar mode. Simultaneously, in the process of cooling the coefficient $m_e^* c/ e\\tau_e$, i.e., the normalization field $B_N$, becomes smaller and smaller until the Hall--term in \\Eqref{equ:indeqdimless} dominates the linear ohmic term. Of course, as a counteracting tendency, the magnetic field decay will continue, increasingly modified by the Hall--cascade and the gradual onset of the Hall--instability. Therefore, we are confronted with a competition of one process increasing $B_0$ (via $\\tau_e$) and another diminishing it (via the {\\em non--normalized} magnetic field). The final decision whether $B_0$ indeed reaches values $O(1000)$ or even higher has to be left to non--linear coupled magneto--thermal calculations. Once the Hall--term dominates the instability may raise small--scale perturbations down to scale lengths $\\gtrsim l_{\\text{crit}}$ on expense of the dipolar mode, resulting in a spin--down behaviour of SOINSs as discussed in Sect.~\\ref{obsmot}.  Another possible observational consequence follows from the dissipative effects of the rapidly excited small--scale field modes, that is, enhanced Joule heating. As can be seen in the upper panel of Fig.~\\ref{fig7}, the growing perturbations cause heat sources located close to the surface of the crust which, in turn, may produce a patchy neutron star surface. Note, that the plot shows only the part linear in $\\bvec$ of the heat source distribution, since the quadratic term $(\\curl\\bvec)^2$ has to be dropped in the linear analysis. These heat sources are concentrated in a very thin layer comprising only less than 5 \\% of the crustal thickness. Clearly, conclusions about the observability of that hot spots are impossible to be drawn from our simple model, since to infer the temperature distribution from the heat sources demands the knowledge of the absolute perturbation amplitudes. Anyway, we expect the neutron star's surface to be warmer during the episode of the Hall--instability than standard cooling calculations predict.  Third, strong small--scale field components cause small--scale Lorentz forces. The lower panel of Fig.~\\ref{fig7} shows that the maximum of the Lorentz force density perturbation is concentrated in deeper regions of the crust (at about 25 \\% of the thickness). From its pattern one may infer shear, both in $y$-- and $z$-- directions, to be the prevailing type of deformation. \\citet{TD95} discussed a scenario in which the helicoidal waves, being a consequence of the Hall--drift, ignite high energetic bursts by cracking the crust. From our results we can hope that the unstable modes of the Hall--instability act even more efficiently than the at best undamped, but never growing helicoidal waves can.  However, all the questions connected with the actual influences of the presented instability on spin--down, cooling and crust--cracking can be decided only by performing non--linear calculations in a realistic spherical--shell geometry taking into account a realistic density profile and chemical composition of the crust. On the other hand, the question at which field strength the action of the Hall--instability finally ceases, that is, the field decay is again overtaken by ohmic dissipation and, as a consequence, to which extent the distribution of observable pulsars in a $B$--$P$, better, in a $B$--$P$--$\\tau_{\\text{a}}$ diagram is possibly influenced, can again only be answered on the basis of realistic non--linear magneto--thermal simulations. Whether the scenario of an accelerated field decay is in agreement with the observed properties of the pulsar population can only be decided by help of a population synthesis as presented by \\citet{RF01} applying a decay pattern with an initially slow decay followed by an episode of rapid decay and leading finally to a nearly constant field."
    },
    "0207/astro-ph0207199.txt": {
        "abstract": " Space is not a boring static stage on which events unfold over time, but a dynamic entity with curvature, fluctuations and a rich life of its own which is a booming area of study. Spectacular new measurements of the cosmic microwave background, gravitational lensing, type Ia supernovae, large-scale structure, spectra of the Lyman $\\alpha$ forest, stellar dynamics and x-ray binaries are probing the properties of spacetime over 22 orders of magnitude in scale. Current measurements are consistent with an infinite flat everlasting universe containing about 30\\% cold dark matter, 65\\% dark energy and at least two distinct populations of black holes. ",
        "introduction": "Traditionally, space was merely a three-dimensional static stage where the cosmic drama played out over time. Einstein's theory of general relativity \\cite{Einstein,WeinbergBook,MTW} replaced this by four-dimensional spacetime, a dynamic geometric entity with a life of its own, capable of expanding, fluctuating and even curving into black holes. Now the focus of research is increasingly shifting from the cosmic actors to the stage itself. Triggered by progress in detector, space and computer technology, an avalanche of astronomical data is revolutionizing our ability to measure the spacetime we inhabit on scales ranging from the cosmic horizon down to the event horizons of suspected black holes, using photons and astronomical objects as test particles. The goal of this article is to review these measurements and future prospects, focusing on four key issues: \\begin{enumerate} \\item The global topology and curvature of space \\item The expansion history of spacetime and evidence for dark energy \\item The fluctuation history of spacetime and evidence for dark matter \\item Strongly curved spacetime and evidence for black holes \\end{enumerate} In the process, I will combine constraints from the cosmic microwave background \\cite{CMB}, gravitational lensing, supernovae Ia, large-scale structure, the Lyman $\\alpha$ forest\\cite{LyAF}, stellar dynamics and x-ray binaries. Although it is fashionable to use cosmological data to measure a small number of free ``cosmological parameters'', I will argue that improved data allow raising the ambition level beyond this, testing rather than assuming the underlying physics. I will discuss how with a minimum of assumptions, one can measure key properties of spacetime itself in terms of a few cosmological functions: the expansion history of the universe, the spacetime fluctuation spectrum and its growth.  %The references have been limited to a minimum and are focused on mentioning key %papers supplemented by a few complete but more specific reviews and %some recent publications, which include more extensive references. %We apologize in advance to all of those colleagues whose work is not directly %referenced.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure}[pbt] %\\hskip-6cm {\\epsfxsize=9.0cm\\epsffile{zoom.eps}} %\\vskip-4.5cm \\smallskip \\caption{ Summary of the spacetime issues discussed in this article. One can use photons and astronomical objects as test particles to measure spacetime over 22 orders of magnitude in scale, ranging from the cosmic horizon (probing the global topology of and curvature of space --- top) %to distant supernovae (giving evidence of dark energy) down to galaxies (giving evidence for dark matter), galactic nuclei and binary stellar systems (giving evidence for black holes). The figure illustrates how spacetime ripples at the $10^{-5}$ level will be imaged by the cosmic microwave background satellite MAP \\protect\\cite{MAP} and has grown via gravitational instability into cosmic large-scale structure \\protect\\cite{VIRGO}, galaxies and, it seems, black holes \\protect\\cite{Bromley}. } \\label{zoomFig} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\subsection{Goals and tools}  %Before embarking on our survey of spacetime, let us briefly review what it is we want to measure %and the basic tools at our disposal \\cite{WeinbergBook,MTW,WillBook}. Before embarking on our survey of spacetime, let us briefly review what it is we want to measure, the basic tools at our disposal \\cite{WeinbergBook,MTW,WillBook} and the broad-brush picture of how our topics fit together. According to general relativity, spacetime is what mathematicians call a manifold, characterized by a topology and a metric. The topology gives the global structure (\\fig{\\ZoomFig}, top): is space infinite in all directions like in high-school geometry or multiply connected like say a hypersphere or doughnut so that traveling in a straight line could in principle bring you back home --- from the other direction? The metric determines the local shape of spacetime, \\ie,  the distances and time intervals we measure, and is mathematically specified by a $4\\times 4$ matrix at each point in spacetime.  General relativity theory (GR) consists of two parts, each providing a tool for measuring the metric. The first part of GR states that in the absence of non-gravitational forces, test particles (objects not heavy enough to have a noticeable effect on the metric) move along geodesics in spacetime, generalized straight lines, so the observed motions of photons and astronomical objects allow the metric to be reconstructed. I will refer to this as {\\it geometric} measurements of the metric. The second part of GR states that the curvature of spacetime (expressions involving the metric and its first two derivatives) is related to its matter content --- in most cosmological situations simply the density and pressure, but sometimes also bulk motions and stress energy. I will refer to such measurements of the metric as  {\\it indirect}, because they assume the validity of the Einstein field equations of GR.  \\subsection{The broad brush picture}  The current consensus in the cosmological community is that spacetime is extremely smooth, homogeneous and isotropic (translationally and rotationally invariant) on large ($\\sim 10^{23}$m$-10^{26} $m) scales, with small fluctuations that have grown over time to form objects like galaxies and stars on smaller scales. Cosmic Microwave Background (CMB) observations have shown \\cite{Smoot92} that space is almost isotropic on the scale of our cosmic horizon ($\\sim 10^{26} $m), % Actually, 9000 Gpc/h ~ 9000 Gpc/0.7 ~ 3.9e26 with the metric fluctuating by only about one part in $10^5$ from one direction to another, and combining this with the so-called cosmological principle, the assumption that there is nothing special about our vantage point, implies that space is homogeneous as well.  Three-dimensional maps of the galaxy and quasar distribution give more direct evidence for large-scale homogeneity \\cite{Colless01,Zehavi01,Hoyle01}.  The fact that the CMB fluctuations are so small is useful, because it allows the intimidating nonlinear partial differential equations governing spacetime and its matter content to be accurately solved using a perturbation expansion. To zeroth order (ignoring the fluctuations), this fixes the global metric to be of the so-called Friedman-Robertson-Walker (FRW) %Friedman-Lemaitre-Robertson-Walker (FLRW) form, which is completely specified except for a curvature parameter and a free function giving its expansion history. To first order, density perturbations grow due to gravitational instability and gravitational waves propagate through the FRW background spacetime, all governed by {\\it linear} equations. Only on smaller scales ($\\simlt 10^{23}$m) do the fluctuations get large enough that nonlinear dynamics becomes important --- in the realm of galaxies, stars and, perhaps, black holes. %The sections below are organized accordingly, discussing spacetime to 0th, 1st and higher order. This review is organized analogously: Sections 2 and 3 discuss spacetime to 0th order (curvature, topology and expansion history), Section 4 describes spacetime to 1st order (fluctuations) and Section 5 focuses on nonlinear objects, mainly black holes. ",
        "conclusions": ""
    },
    "0207/astro-ph0207409_arXiv.txt": {
        "abstract": "SN~1999aw was discovered during the first campaign of the Nearby Galaxies Supernova Search ({\\it NGSS}) project.  This luminous, slow-declining ($\\Delta$$m_{15}(B)$$= 0.81\\pm 0.03$) Type Ia supernova was noteworthy in at least two respects.  First, it occurred in an extremely low luminosity host galaxy that was not visible in the template images, nor in initial subsequent deep imaging.  Secondly, the photometric and spectral properties of this supernova indicate that it very likely was similar to the subclass of Type Ia supernovae whose prototype is SN~1999aa.  This paper presents the {\\it BVRI} and {\\it J$_{s}$HK$_{s}$} lightcurves of SN~1999aw (through $\\sim100$ days past maximum light), as well as several epochs of optical spectra.  From these data we calculate the bolometric light curve, and give estimates of the luminosity at maximum light and the initial $^{56}$Ni mass. In addition, we present deep {\\it BVI} images obtained recently with the Baade 6.5-meter telescope at Las Campanas Observatory which reveal the remarkably low-luminosity host galaxy. ",
        "introduction": "Type Ia supernovae (SNe~Ia) offer arguably the most precise method to measure cosmological distances.  Over the last ten years, these highly luminous explosions have been used to determine distances accurate to $7\\%$, despite the fact that we know little about their progenitors.  These objects show considerable uniformity in their absolute {\\it B} magnitudes at maximum light with an intrinsic dispersion of less than 0.4 mag (Hamuy et al.  1996b).  This scatter is greatly reduced by the application of empirical relations linking the luminosity at maximum to the width, or decay rate, of the lightcurve (Luminosity-Width Relations, or {\\em LWR}s).  More luminous SNe~Ia decline in brightness after maximum at slower rates than less luminous SNe~Ia.  The $\\Delta$$m_{15}(B)$ relation (Phillips 1993, Hamuy et al.  1996a, Phillips et al.  1999), which is a measure of the decay in the {\\it B} band lightcurve from peak to 15 days after peak, has reduced the scatter around the hubble law to less than 0.2 mag, making it a powerful tool in using SNe~Ia at high redshifts to investigate cosmological parameters.  Including reddening corrections further decreases the dispersion in the Hubble diagram to 0.14 mag (Phillips et al.  1999).  An equally effective method developed by Riess, Press, \\& Kirshner (1996) uses the lightcurve shapes in multiple passbands to simultaneously estimate the SN~Ia luminosity and amount of extinction/reddening.  This MLCS method has demonstrated it can produce Hubble diagrams with dispersions of only 0.12 mag (6\\% in distance).  Attention is now being directed to the possible systematic errors involved in using these objects as high redshift distance indicators. Part of the challenge lies in untangling the dispersion of SNe~Ia lightcurve widths from possible sub-populations of Type Ia supernovae.  With more discoveries of nearby supernovae in various host environments, and the development of new spectroscopic and photometric techniques for isolating these sub-populations, we may soon be able to determine $\\Omega_{M}$ and $\\Omega_{\\Lambda}$ with lower systematic uncertainties as well as hone in on the progenitors of SNe~Ia.  Over the past three years, the Nearby Galaxies Supernova Search Team {\\it(NGSS)} has conducted successful search campaigns for supernovae of all types using the Kitt Peak National Observatory's 36-inch telescope and the Mosaic North camera (Mosaic I).  Each of our 5 - 8 night campaigns have allowed us to search $\\sim$ 250 fields (each nearly 1$\\arcdeg$ square) along the celestial equator and out of the Galactic plane to limiting magnitudes of $R \\sim 21$.  At the project's end, we had searched nearly 750 fields and discovered 42 supernovae.  The goals of this project have been to understand supernova rates (for all SNe types) in both field and galaxy cluster environments, to investigate correlations of SN type with host galaxy environments, and to increase knowledge of observationally rare and peculiar SN types through increased detailed photometric and spectroscopic observations.  Further information concerning the {\\it NGSS} project goals and methods will be described in a forthcoming paper (Strolger et al.  2003b, in preparation, hereafter referred to as Paper 2).  SN~1999aw was one of the supernovae discovered during the first of the {\\it NGSS} campaigns (Feb.  20 - Feb.  24, and Mar.  04 - Mar. 09, 1999) which was conducted in co-operation with the Supernova Cosmology Project ({\\it SCP}; Aldering 2000, Nugent \\& Aldering 2000).  Initial discovery and confirmation images surprisingly showed a bright new object in a location which was devoid of galaxies on the template images taken only a few weeks earlier.  Spectroscopic and photometric evidence show SN~1999aw is a member of an intriguing subclass of Type Ia supernovae, characterized spectroscopically by SN~1999aa (Li et al.  2001).  These unusual supernovae have lightcurve shapes that are slightly different from those of normal Type Ia (Strolger et al.  2003a, also in preparation; see also Section~\\ref{subsec:templates}).  These differences, although not yet completely understood, may be key in understanding not only the physical processes of this subclass, but of all SNe~Ia, and may place limits on models of Type Ia progenitors.  In Section~\\ref{sec:discovery} we discuss the discovery, confirmation, and classification of SN~1999aw.  In Section~\\ref{sec:spectroscopy} we present several epochs of optical spectra and discuss their similarity to 1999aa-like supernovae.  In Section~\\ref{sec:lcurves} we present the optical and infrared photometry, and the calibrations and corrections.  In Section~\\ref{sec:analysis} we show the photometric similarity to 1991T/1999aa SNe, determine the bolometric lightcurve, estimate the luminosity at maximum light, and estimate the initial $^{56}$Ni mass. In Section~\\ref{sec:galaxy} we present the photometry of the host galaxy and discuss the host environment. ",
        "conclusions": "\\label{sec:analysis} \\subsection{{\\it B}-band Template Fits}\\label{subsec:templates}  The lightcurves of SN~1999aw are remarkable for how much slower they evolve compared to the template curves of a typical decline-rate SNe Ia (see Figure~\\ref{fig:lightcurve}).  This is evident not only in the slow initial decline rate of $\\Delta$$m_{15}(B) = 0.81$ (the mean decline rate of SNe Ia is $\\Delta$$m_{15}(B)\\sim 1.1$), but in the delay of the second maximum in {\\it I, J$_{s}$, H,} and {\\it K$_{s}$}.  The {\\it B}-band lightcurve is particularly interesting in that its ``shape'' is subtly different than that for typical SNe Ia.  One parameter LWR relations such as the $\\Delta$$m_{15}(B)$ suggest that a {\\it B}-band lightcurve template can be made to fit the {\\it B}-band lightcurve of a SN by applying a time delaying ``stretch'' factor to the epochs of observation, and the related magnitude offset (Perlmutter et al.  1997).  However, Strolger et al.  (2000 \\& 2003a) show that {\\it B}-band template lightcurves built from well-sampled spectroscopically normal SNe~Ia fit poorly to the {\\it B}-band lightcurves of 1991T/1999aa-like supernovae, even when stretched to fit the observed lightcurve in early epochs (By ``normal'', we mean those those SNe~Ia which displayed strong Si~II$\\lambda6355$ absorption at least 5 days before {\\it B}-band maximum).  The Strolger et al.  (2000 \\& 2003a) analysis was conducted on a few SNe that were 1) spectroscopically identified at least 5 days before maximum light, and 2) frequently observed with well sampled {\\it B}-band lightcurves from just prior to maximum light to around the $+$80 day epoch.  The low-order mean difference between the template curve and the supernovae lightcurves from the beginning of the exponential phase (after the $+$25 day epoch) was determined for each supernova in the sample:   \\begin{equation} \\delta_{ave} = \\sum_{+25}^{+80} \\frac{m(t) - m_{T}(t^{\\prime})}{N} \\end{equation}  Observations made at some epoch, $m(t)$, were compared to the time stretched template curve, $m_{T}(t^{\\prime})$, and then averaged over the number of observations ($N$) from the $+$25 day to $+$80 day epochs.  Results of the analysis show that spectroscopically normal SNe Ia exhibit little to no difference from the template curve during the exponential phase, whereas 1991T/1999aa-like SNe Ia show a substantial overbrightness during this phase.  The analysis, when performed on SN~1999aw, produced a result consistent with that of the 1991T/1999aa-like SNe.  Figure~\\ref{fig:templates} shows four SNe spectroscopically similar to SNe~1991T and/or 1999aa (including SN~1999aw), along with the normal SNe~Ia template, stretched in time to fit the lightcurves within the first 15 days past maximum light.  The template does not fit well to the data past around the $+25$ day epoch, and the data are systematically brighter than the template from that epoch.  \\begin{figure*} \\plotone{strolger.f7.eps} \\caption{{\\it B}-band lightcurves of four 1991T/1999aa-like supernovae, scaled to {\\it B}$_{max}$. Solid line is a Template for spectroscopically normal Type Ia SNe, stretched along time axis to fit observations between peak and $+$15 days past maximum light.\\label{fig:templates} } \\end{figure*}  In the future, with more examples of 1991T/1999aa-like SNe, this analysis may lead to photometric method that, in addition to spectroscopy, indicate the possible 1991T/1999aa-like peculiarity of supernovae.   \\subsection{Color Curves}\\label{subsec:colorcurves}  SNe~Ia show a impressive uniformity in their intrinsic colors in late epochs after maximum light.  As Lira (1995) and Riess et al.  (1996) independently showed, SNe~Ia with $0.85\\lesssim$$\\Delta$$m_{15}(B)$$\\lesssim 1.90$ and little or no reddening from their host galaxies have very uniform $B-V$ color evolution from $+$30 to $+$90 days after maximum light.  Krisciunas et al.  (2000) also observe uniformity in the {\\it V}$-${\\it Near IR} color evolution of spectroscopically normal SNe Ia with mid-range decline rates from $-$9 days to $+$27 days past maximum light.  Figure~\\ref{fig:colors} shows the evolution in multiple colors of SN~1999aw, corrected for Galactic reddening assuming an excess of $E(B-V) = 0.032$ (Schlegel, Finkerbeiner, \\& Davis 1998).  The left-hand side panels of Figure~\\ref{fig:colors} show the optical color evolution, along with some example SNe (zero-reddening corrected) for comparison.  The solid line in the plot of the evolution of $B-V$ in this figure represents the zero-reddening least-squares fit derived by Lira (1995).  SN~1992al is included to show the evolution of a typical SNe Ia, while SN~1992bc is a slow-declining SNe Ia ($\\Delta$$m_{15}(B) = 0.87$) but with normal pre-maximum light spectra.  SNe~1991T and 1999aa, the prototypical 1991T/1999aa-like SNe, have color evolutions nearly parallel to the Lira (1995) line, thus showing that the uniformity holds for {\\em some} spectroscopically extreme SNe~Ia.  \\begin{figure*} \\plotone{strolger.f8.eps} \\caption{Color curves of SN~1999aw with photometric errors. B$-$V curve includes Lira zero-reddening relation (Phillips et al. 1999). V$-${\\it Near IR} curves include zero-reddening relations from Krisciunas et al. (2000) for Type ia SNe with mid-range decline rates.\\label{fig:colors}} \\end{figure*}  This color uniformity has proven useful as an indicator of host galaxy reddening, which has been important to revising the LWR relations (Phillips et al.  1999).  We have used the recipes in Phillips et al.  (1999) to determine the host galaxy reddening for SN~1999aw.  The analysis gave $E(B-V)_{tail}=-0.20\\pm0.05$, $E(B-V)_{max}=0.12\\pm0.10$, and $E(V-I)_{max}=-0.11\\pm0.09$, which when averaged, resulted in a negative color excess, implying zero host galaxy reddening.  Clearly the negative value of $E(B-V)_{tail}$ is heavily influenced by the anomalous $B-V$ color evolution of SN~1999aw; specifically, for most of the period from $+$40 days to $+$60 days SN~1999aw was considerably bluer than the Lira (1995) zero-reddening fit, and perhaps evolving with a steeper slope.  As SN~1999aw was spectroscopically similar to SN~1999aa, we might have expected that its color evolution would be similar to SN~1999aa and/or SN~1991T, both of which appeared generally bluer than SN~1992bc in $V-R$ and $V-I$.  In SN~1999aw, the bluer $V-R$ and $V-I$ evolution were caused by the delay in the appearance of the second lightcurve maximum, which was more clearly seen in redder passbands (see Figure~\\ref{fig:lightcurve}), and therefore is accentuated in the $V-R$ and $V-I$ curves.  SN~1999aw also appeared bluer than SN~1999aa and SN~1991T in $V-R$ and $V-I$, which perhaps was an effect of SN~1999aw having a much slower decline rate than both SN~1999aa and SN~1991T.  The $V-J_{s}$ and $V-H$ curves for SN~1999aw are also considerably bluer than the zero-reddening fits from Krisciunas et al.  (2000) in the period after $+$10 days (right-hand side panels of Figure~\\ref{fig:colors}).  The trend in the period prior to $+$10 days in the $V-J_{s}$ curve also seems to be different than the fit from Krisciunas et al.  It appears that one could apply a ``stretch'' to the Krisciunas lines to force them to fit the SN~1999aw data, that is to say there is an apparent delay in the $V-J_{s}$ evolution. This again is due to the apparent delay of the second maximum in the {\\it I, J$_{s}$, H,} and {\\it K$_{s}$} lightcurves. Note, however, that the $V-J_{s}$, $V-H$, and $V-K_{s}$ colors in this figure may change somewhat once K-corrections for the IR bandpasses become available.  \\subsection{Bolometric lightcurve, Maximum Luminosity, and $^{56}$Ni Mass Estimate.}\\label{subsec:bol}  As nearly all of the bolometric luminosity of a typical Type Ia supernova is emitted in the range of 3000 to 10000{\\AA} (Suntzeff 1996), the integrated flux in the {\\it UBVRIJ$_{s}$HK$_{s}$} bandpasses provides a reliable and meaningful estimate of the bolometric luminosity, which is directly dependent on the amount of nickel produced in the explosion.  The {\\it BVRIJ$_{s}$HK$_{s}$} data were used to calculate ``{\\it uvoir}'' bolometric fluxes using the techniques described in Suntzeff (1996) and Suntzeff and Bouchet (1990).  A table of {\\it UBVRIJ$_{s}$HK$_{s}$} data was made by linearly interpolating the {\\it J$_{s}$HK$_{s}$} data to the optical dates.  For some of the missing optical and {\\it J$_{s}$} and {\\it H} data, we added photometry based on spline fitting of the data to the date of the missing data.  We have added on {\\it U} data because the optical ultraviolet adds significant flux to the early time bolometric light curve.  There is little high-quality {\\it U} data for Type Ia supernovae due to the generally poor ultraviolet sensitivity of the present generation of CCDs.  We have instead relied on $U-B$ photometry from photoelectric measurements.  For dates past $+$9 days from {\\it B} maximum, we have used the $U-B$ data of SN~1972E from Lee et al.  (1972) and Ardeberg \\& de Groot (1973).  For dates before $+$9 days from {\\it B} maximum we have used the $U-B$ data for SN~1980N from Hamuy et al.  (1991), and for SN~1981B compiled by Cadonau \\& Leibundgut (1990).  The $U-B$ for these supernovae were corrected to the reddening of SN~1999aw using the reddening values in Phillips et al.  (1999) and a value of $E(U-B)/E(B-V)=0.72$ from Cardelli et al.  (1989).  The {\\it U} photometry of SN~1999aw was then estimated from spline fits to the $U-B$ data of these supernovae combined with our {\\it B} data of SN~1999aw.   We then converted the broadband magnitudes to equivalent monochromatic fluxes at the effective wavelengths of Vega (Bessell 1979, 1990; Bessell \\& Brett 1988).  A magnitude scale of $(U,B,V,R,I)=0.03$ and $(J,H,K)=0.0$ was used for Vega.  The monochromatic fluxes were then scaled to the magnitude of the supernova, dereddened by $E(B-V)=0.032$ using the reddening law of Cohen et al.  (1981), and corrected to intrinsic fluxes using a distance modulus of 36.28 based on a Hubble flow with a Hubble constant of 63.3 km/s/Mpc (Phillips et al.  1999).  These fluxes were then integrated using a simple trapezoidal integration.  We added on a Rayleigh-Jeans extrapolation to zero frequency to the reddest flux point.  We extrapolated to the ultraviolet by adding a flux point at 3000{\\AA} with zero flux.  We correct the derived bolometric lightcurves for time dilation effects, and plot them in Figure~\\ref{fig:bolcurve}.  The {\\it UBVRI} and the {\\it UBVRIJ$_{s}$HK$_{s}$} integrations track each other well, except that the inflection points around days 20-45 are more pronounced in the {\\it UBVRIJ$_{s}$HK$_{s}$} integrations.  Similar inflection points were noted by Suntzeff (1996) and are indicative of significant flux redistributions which may be related to rapid changes in the wavelength dependence of the opacities (Pinto \\& Eastman 2000 a \\& b).  The peak bolometric luminosity is about $L_{bol} = 1.51 \\times 10^{43}$ erg s$^{-1}$.  \\begin{figure*} \\epsscale{1.0} \\plotone{strolger.f9.eps} \\caption{The bolometric lightcurve of SN~1999aw, constructed from the integrated flux in the {\\it B, V, R, I, J$_{s}$, H,} and {\\it K$_{s}$} passbands, and corrected for time dilation.  Stars show the bolometric curve in UBVRI, while open circles are in UBVRIJ$_{s}$H, and filled circles are in UBVRIJ$_{s}$HK$_{s}$.\\label{fig:bolcurve}} \\end{figure*}   It is fairly straightforward to derive the nickel mass produced in the explosion from the bolometric luminosity at peak.  At maximum light, photons escape the surface at a rate which is equal to the radioactive energy input produced primarily by the $^{56}$Ni decay, and thus it is also related to the $^{56}$Ni synthesized in the explosion (Arnett 1982, Nugent et al. 1995b, Pinto \\& Eastman 2000a).  Contardo et al.  (2000) calculate the luminosity and nickel mass for several SNe~Ia from {\\it UBVRI} bolometric peak fluxes.  Using the same method, and assuming a rise time of 17 days to bolometric peak for consistency with Contardo, et al., we derive an initial nickel mass of $M_{Ni} = 0.76 M_{\\odot}$ for SN~1999aw.  This is brighter and more nickel massive than many of the normal Type Ia SNe discussed in Contardo, et al., and it is comparable in brightness and nickel mass to SN~1991T (see Table~\\ref{table:masses}).  However as they note, a number of the SNe~Ia in their study do not have {\\it U}-band data available at peak, and therefore they have developed a correction curve based on data from the very well-sampled SN~1994D. Although this SN was spectroscopically normal, it did have some rather unusual features, including an unusually blue $U-B$ color at maximum light.  Additionally, Riess et al.  (1999) have shown that the characteristic rise time of SNe~Ia is $19.5\\pm0.2$ days, and that brighter and slower declining SNe~Ia have longer rise times.  For the peak magnitude and decline rate observed in SN~1999aw, we expected a rise time of $\\sim20$ days.  Using this value, we derive an alternative initial nickel mass of $M_{Ni}=1.07 M_{\\odot}$.  There are a number of additional uncertainties in both the calculated luminosity, and the derived nickel mass.  The first is the assumption that more than 80\\% of the true bolometric light is emitted in the optical regime, that less than 10\\% can be expected in the UV below 3200{\\AA}, and that no more than 10\\% (in early epochs) from {\\it JHK} (Suntzeff 1996, Elias et al.  1985, Contardo et al.  2000).  In comparing our derived {\\it UBVRIJ$_{s}$HK$_{s}$} and {\\it UBVRI} bolometric fluxes, we find the IR contribution to be only a few percent ($\\sim 1\\%$ at early epochs, $\\sim 5\\%$ after 35 days past $B_{max}$).  We have not accounted for the space ultraviolet flux.  Also, as we will discuss further in Section~\\ref{sec:galaxy}, we do not have much information about the host galaxy of SN~1999aw. Although we have compensated for extinction and reddening due to our own galaxy, it difficult to do the same for the host galaxy. However, as the $B-V$ color of SN~1999aw is nearly zero at maximum light, as it would be for an unreddened Type Ia SN (Phillips et al. 1999; Garnavich et al.  2001, in Figure 15), we have assumed that the extinction due to the host must be negligibly small.  \\label{sec:summary}  Our photometric and spectroscopic study of SN~1999aw indicate that SN~1999aw was probably a 1999aa-like event.  The light curve decline rate is among the slowest observed at $\\Delta$$m(B)_{15} = 0.81 \\pm 0.03$.  At the redshift of $z = 0.038$, it is also among the brightest Type Ia, with $M_{B}=-19.45\\pm0.11$, $M_{V}=-19.50\\pm0.11$, $M_{R}=-19.38\\pm0.11$, and $M_{I}=-18.97\\pm0.12$.  Although luminous, these magnitudes are slightly less (by a few tenths of a magnitude) than one might expect from the trends of absolute magnitude versus decline rate or $\\Delta$$m_{15}$ (See Figure 11 of Krisciunas et al. 2001).  Perhaps this suggests that there is some host galaxy extinction that needs to be accounted for.  Or it is possible that SN~1999aw is simply slightly less luminous than expected, and therefore may be indicating a possible downward curve to the relations in Fig.  11 for the slowest-declining SNe Ia (similar to the downward curve recently determined by Garnavich et al.  (2001) for fast-declining SN 1991bg-like SNe~Ia).  It is thought that the brightness and decline rate of Type Ia are dependent on the metallicity of the progenitor, and additionally on the opacity in the atmosphere of the event (Hoflich et al.  1998, Pinto and Eastman 2000a \\& b, Mazzali et al.  2001).  The derived luminosity of $1.51 \\times 10^{43}$ erg s$^{-1}$, and $^{56}$Ni mass of $0.76 M_{\\odot}$ for SN~1999aw are both relatively high, but are consistent with what can be expected for a SN Ia with the observed decline rate, as inferred from the trend between $\\Delta$$m_{15}(B)$ and $^{56}$Ni mass evident in Table~\\ref{table:masses}.  The observations made at the Baade 6.5-meter have provided some very interesting information on the host galaxy of SN~1999aw.  The derived absolute magnitudes seem to indicate that it is among the intrinsically faintest and bluest dwarf galaxies observed.  This may indicate that the galaxy consists of either a young population of stars, or is a fairly metal-poor environment.  It will be necessary to eventually obtain deeper images with large telescopes such as the Magellan 6.5-meter or the VLT to obtain sufficient signal-to-noise to put more accurate constraints on not only the colors of this galaxy, but its structure as well.  The discovery of SN~1999aw exemplifies an advantage of magnitude limited field searches over galaxy-targeted searches.  It is presently unclear how many supernova similar to SN~1999aw occur at low redshift, and with the biases associated with current targeted surveys we are less likely to detect them.  As more magnitude limited field searches commence, relations between low luminosity (and probably low metallicity) galaxies and the supernova types that occur in them can be accurately determined.  This will not only place constraints on Type Ia supernova progenitor models, but help to clarify the possibility of supernova ``evolution'' which may explain the difference in SN~1991T-like event rates between the low and high-$z$ surveys (Li et al.  2001)."
    },
    "0207/astro-ph0207123_arXiv.txt": {
        "abstract": "The production and survival of the quasistable isomer $^{180}$Ta during the stellar nucleosynthesis has remained a matter of discussion for years. A careful analysis of the available experimental data and theoretical calculations enabled us to reproduce the observed solar abundance of $^{180}$Ta even in the classical s-process ($kT=28$ keV -- 33 keV). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207137_arXiv.txt": {
        "abstract": "We report on the combined X-ray and radio observations of the type Ic SN 2002ap, using XMM-Newton ToO observation of M74 and the Giant Metrewave Radio Telescope (GMRT). We account for the presence of a nearby source in the pre-supernova Chandra field of view in our measurements of the X-ray flux (0.3 - 10 KeV) 5.2 days after the explosion. The X-ray spectrum is well fitted by a power law spectrum with photon index $\\alpha= 2.6$. Our results suggest that the prompt X-ray emission originates from inverse Compton scattering of photospheric thermal emission by energetic electrons. Radio observations with the GMRT at 610 MHz (8 days after the explosion) and 1420 MHz (70 days after the explosion) are combined with the high frequency VLA observations of SN 2002ap reported by \\cite{Ber02}, and the early radiospheric properties of SN 2002ap are compared with similar data from two other supernovae. Finally, the GMRT radio map reveals four other X-ray sources in the field of view of M74 with radio counterparts. ",
        "introduction": "\\label{sec: Intro} Supernovae are known to be explosions of massive and intermediate mass stars, although the nature of the progenitor star for different supernova types, remains an area of long-standing research.  Currently, it is believed that type II and type Ib or Ic supernovae arise from core collapse of massive stars, while the more homogeneous type Ia supernovae are the results of thermonuclear explosions.  There is a considerable variety among the core collapse supernovae spectroscopically and in their kinetic energy.  The `hypernova' class among these are believed to have explosion energy significantly in excess of $10^{51} \\rm erg$, found in \"normal\" supernovae. The spatial and temporal near-coincidence of type Ic SN 1998bw with the long duration GRB980425 has provided further impetus to observational and theoretical studies of type Ib/Ic SNe and their progenitors (\\cite{Mac01}, \\cite{Mes93}).   During a supernova explosion the interaction of the outer parts of the stellar ejecta with the circumstellar matter gives rise to a high energy density shell. X-ray emission is expected from both shocked circumstellar matter and the shocked supernova matter (see e.g. \\cite{Che82}, \\cite{Che01}). In addition, the interaction region may also accelerate electrons to relativistic energies and amplify pre-existing magnetic fields which gives rise to nonthermal synchrotron emission seen in many supernovae. The radio and the X-ray emission give information on a region of the supernova which may be far removed from the optical photosphere (which has a smaller radius), although the conditions in the optical photosphere may determine the X-ray emission characteristics in some instances.  In general, X-ray and radio observations in early stages of a supernova can be used to determine (1) the total mass lost from the pre-supernova star before explosion and (2) constrain various physical processes leading to X-ray and radio emission.  The type Ic SN 2002ap was discovered on Jan.  29.4, 2002 (Y.  Hirose as reported by \\cite{Nak02} ) in NGC  628 (M74), at a distance of only $7.3$ Mpc. Based on spectral analysis of the early observations, the epoch of explosion was estimated at Jan.  $28.0 \\pm 0.5$ 2002UT (\\cite{Maz02}). For the purpose of the analysis presented in this paper, we will regard Jan.  $28.0$ as the date of explosion.  The broad spectral features (\\cite{Kin02},  \\cite{Mei02}, \\cite{Gal02}), and a subsequent modeling of its spectroscopic and photometric data (\\cite{Maz02}), suggested that this was an energetic event with explosion energy $E \\simeq 4 -10 \\times 10^{51}$ erg. In this paper, we discuss our analysis of the XMM-Newton observation in the X-ray (0.1-15 keV) bands  (Sect. \\ref{sec: XrayObs}), accounting for the presence of a nearby object in the pre-supernova X-ray field, observed earlier by the Chandra X-ray Observatory (Sect.  \\ref{sec: Chandra}). We also observed SN 2002ap in the 0.61 GHz and 1.4 GHz radio bands (see Sect. \\ref{sec: GMRT}), and the implications of the  GMRT upper limits in the context of VLA observations at the same epoch (\\cite{Ber02}) are discussed in Sect. \\ref{sec: combined}.  We also summarize the explosion parameters that we derived from the optical observations and modeling reported by \\cite{Maz02}, as this is used as input for later sections (Sect. \\ref{sec: combined}). We have also combined the GMRT data with the VLA data (\\cite{Ber02}) to derive conditions near the radiosphere. A combined analysis of the early X-ray and radio observations is presented in Sect. \\ref{sec:  massloss}, which attempts to constrain the multiple physical processes (thermal and non-thermal) that are responsible for the early X-ray emission. We also compare SN 2002ap with 2 other SNe (1998bw (Ic) and 1993J (IIb) for which early multi frequency data is available. Finally, in Sect. \\ref{sec:  radiative} we discuss these results in the context of the presupernova star and its evolution. ",
        "conclusions": "\\label{sec: Discussion}  In this paper, we presented an analysis of the XMM data of the SN 2002ap field and have obtained spectral model fits to the prompt X-ray emission. We compare the X-ray image with the GMRT 610 MHz radio image obtained three days apart. While we find no radio counterpart of the SN at such low frequencies, several sources in the field have radio and X-ray counterparts.  We compare the radio data obtained from three different supernovae in their early phases and model these using the synchrotron self absorption model. SN 1998bw(Ic)  with a GRB counterpart had very different radiosphere radius and equipartition angular sizes at approximately the same time in their evolution compared to two other SNe: SN 2002ap(Ic)  and SN 1993J(IIb).  We model the early X-ray emission with inputs from optical photometry and light curve and find that the inverse Compton scattering of optical photons from the supernova photosphere by hot electron plasma can account for the observed early X-ray flux and the spectrum for modest electron temperatures and optical depths. Thermal processes are inefficient initially and would be important only as the supernova ages.  Mass loss rates and stellar wind velocities of the progenitor stars determine the optical depth of shock heated matter due to electron scattering, -- a key parameter in the production of the X-ray flux from the lower energy photons. These in turn depend upon the scenario and progenitor configuration giving rise to type Ic SNe, e.g.:  1. Massive Wolf-Rayet(WR) stars which have lost their hydrogen and helium envelope before the explosion (\\cite{Lan99}) have an empirically determined mass loss rate of $1.5 - 3.2 \\times 10^{-5} M_{\\odot}$ (after taking into account effects due to clumpy medium, (\\cite{Ham98}, \\cite{Wil98}) and terminal wind velocities $\\sim 1000$ km s$^{-1}$ depending upon the type of the WR star.  2. Interacting binaries - in particular, in a case BB mass transfer from a helium star overflowing its Roche-lobe to a companion removes most of the helium rich layers before the type Ic SN. Habets (1985) finds that a 2.5 $M_{\\odot}$ helium star during carbon shell burning stage expands to a red giant dimension of 18 $R_{\\odot}$ and sustains an average mass transfer rate of $10^{-4} \\rm M_{\\odot} yr^{-1}$ lasting about 3000 years. During this time, the terminal wind speed, of the mass losing star would be typically $100$ km s$^{-1}$.   Using the relevant parameters in the above two scenarios, we derive electron optical depth encountered by the intermediate energy photons as listed in the Table \\ref{tab: progenitor}. It is evident that both ranges of electron optical depths for the Compton boosting process remain viable alternatives. For the case of the interacting binary model where the optical depths are somewhat larger, the implied electron temperatures required for the plasma would be lower than in the single star WR model. Such temperatures are well within the range expected for hot circumstellar gas ($T_{cs} \\geq 3.6 \\times 10^9 K$) even for modest velocities of $16,000 - 20,000$ km s$^{-1}$ for the hot electron plasma moving above the optical photosphere."
    },
    "0207/astro-ph0207301_arXiv.txt": {
        "abstract": "We create stacked composite absorption spectra from {\\it Hubble Space Telescope} Faint Object Spectrograph data from four quasi-stellar objects to search for absorption lines in the extreme ultraviolet wavelength region associated with \\lya\\ forest absorbers in the redshift range $1.6 < z < 2.9$. We successfully detect \\ion{O}{5} \\lam 630 in \\lya\\ absorbers throughout the $10^{13}$ to $10^{16.2}$ \\percm\\ column density range. For a sample of absorbers with $10^{13.2} <$ \\nhi\\ $< 10^{14.2}$ \\percm, corresponding to gas densities ranging from around the universal mean to overdensities of a few, we measure an \\ion{O}{5} \\lam 630 equivalent width of $10.9 \\pm 3.7$ m\\AA.  We estimate the detection is real with at least 99\\% confidence. We only detect \\ion{O}{4} \\lam 788, \\ion{O}{4} \\lam 554, \\ion{O}{3} \\lam 833, and \\ion{He}{1} \\lam 584 in absorbers with \\lya\\ equivalent widths $\\gtrsim 0.6$ \\AA, which are likely associated with traditional metal-line systems.  We find no evidence in any subsamples for absorption from \\ion{N}{4} \\lam 765, \\ion{Ne}{5} \\lam 568, \\ion{Ne}{6} \\lam 559, \\ion{Ne}{8} \\lamlam 770, 780, or \\ion{Mg}{10} \\lamlam 610, 625. The measured equivalent widths of \\ion{O}{5} suggest values of \\ovhi\\ in the range $-1.7$ to $-0.6$ for $10^{13.2} <$ \\nhi $\\lesssim 10^{15}$ \\percm. The lack of detectable \\ion{O}{4} absorption except in the strongest absorption systems suggests a hard ionizing background similar to the standard Haardt \\& Madau spectrum.  Using photoionization models, we estimate that the oxygen abundance in the intergalactic medium with respect to the solar value is $[{\\rm O}/{\\rm H}] \\approx -2.2$ to $-1.3$. Comparing to studies of \\ion{C}{4}, we estimate $[{\\rm O}/{\\rm C}]\\approx 0.3$ to 1.2. The overabundance of oxygen relative to carbon agrees with other low-metallicity abundance measurements and suggests enrichment of the intergalactic medium by Type II supernovae. ",
        "introduction": "}  Early studies of the intergalactic medium (IGM) identified two separate groups of absorbers \\citep[e.g.,][]{sybt80}:  (1) ``metal-line systems'' which have observable metal lines and the strongest \\lya\\ absorption, typically with a \\lya\\ rest-frame equivalent width $\\gtrsim$ 1 \\AA, and (2) weaker absorbers having no observable metal lines and comprising what is generally thought of as the ``\\lya\\ forest''.  However, it was suspected that this distinction was due at least partly to observational limitations in the observability of weak metal lines \\citep{tytl87}, a suspicion that has been verified as technological and computational advances have allowed astronomers to detect metal absorption to much lower column densities. Constraining the abundances of metal ions in the IGM has become an active area of research, since this information can in principle be used to determine the chemical abundances in the IGM as well as shape and strength of the typical metagalactic ionizing background radiation from stars and quasi-stellar objects (QSOs).  Of particular interest is the low column density regime (\\nhi $\\lesssim 10^{14.5}$ \\percm), since these systems are presumably far away from local sources of metals and ionizing photons and thus offer the best insights into enrichment mechanisms and the ionizing background.  The production of composite or stacked absorption spectra of many \\lya\\ absorbers is an obvious way to increase observational sensitivity.  \\citet*{nhp83} used the composite technique to claim a detection of \\ion{O}{6} \\lamlam 1032, 1038 absorption.  The reality of this detection was challenged in a similar study by \\citet{wcwb+89}, but \\ion{O}{6} in the IGM was detected unambiguously by \\citet{lusa93} by creating a composite absorption spectrum of \\ion{C}{4} absorbers.  Composite spectra were used to detect \\ion{C}{4} absorption in systems with \\nhi $\\gtrsim 10^{14}$ \\percm\\ at high redshift ($z > 1.7$) by \\citet{lu91}, but \\citet{tyfa94} failed to detect \\ion{C}{4} associated with weaker lines. In a study similar to \\citet{lu91} but at $z < 0.8$, \\citet{baty98} found an order of magnitude stronger \\ion{C}{4} absorption than \\citet{lu91}.  A more precise way to increase sensitivity is to build bigger telescopes and better instruments in order to acquire higher signal-to-noise ratio (S/N), higher resolution data.   With the HIRES instrument on the 10m Keck~I telescope, individual \\ion{C}{4} features can be detected down to a column density of  \\nciv $\\sim 10^{12}$ \\percm. The majority of \\lya\\ absorbers with \\nhi $> 3 \\times 10^{14}$ \\percm\\  exhibit such absorption \\citep{tfbc+95,cskh95,soco96}. Besides allowing for the direct detection of weaker lines, the existence of such high quality data has opened the door to an additional detection method introduced by \\citet{coso98}, that of comparing the individual pixel optical depths of \\lya\\ to the expected position of corresponding metal lines.  \\citet{coso98} and \\citet{essp00} use this technique to detect \\ion{C}{4} absorption down to an \\ion{H}{1} optical depth of $\\tau_{\\rm HI} \\lesssim 1$.  \\citet{srsk00} have used the optical depth technique to detect \\ion{O}{6} to optical depths down to $\\tau_{\\rm HI} \\sim 0.1$.  Studies have concentrated on \\ion{C}{4} not only due to the fact that it is expected to be one of the strongest features associated with the IGM (\\citealt*{rhs97}; \\citealt{hhkw98}), but also largely because of its near ideal location in wavelength space.  With a wavelength of 1549 \\AA, this absorption feature is near enough to \\lya\\ \\lam 1216 that both features can easily be acquired with the same observations, but far enough above \\lya\\ that the feature will be longward of the \\lya\\ forest for a significant span in redshift.  Although \\ion{C}{4} is a good tracer of metals in the IGM, in standard photoionization models the \\ion{O}{6}  \\lamlam 1032, 1038 \\AA\\ doublet is expected to be a stronger feature at \\ion{H}{1} column densities $\\lesssim 10^{15}$ \\percm\\ \\citep{hhkw98}. Unfortunately, although the wavelength of \\ion{C}{4} makes it very appealing for absorption studies, the \\ion{O}{6} \\lamlam 1032, 1038 \\AA\\ doublet could hardly reside in a worse portion of the spectrum.  The region of \\ion{O}{6} absorption is muddled not only by \\lya\\ absorption but by higher order Lyman lines as well, making detections much more difficult.  Fortunately, \\citet{srsk00} have overcome this difficulty to detect \\ion{O}{6} in the diffuse IGM, although they make no quantitative claims as to the implied amount of \\ion{O}{6} in the IGM.  The extreme ultraviolet (EUV) absorption region has been largely ignored, not because this region contains no useful information, but because the best available data in this region, that from the {\\it Hubble Space Telescope} ({\\it HST}) Faint Object Spectrograph (FOS), is inferior in resolution and S/N to that used at longer wavelengths.  However, some of the strongest expected absorption features occur in this region, in particular \\ion{O}{5} \\lam 630 and \\ion{O}{4} \\lam 788 \\citep*{vtb94}.  The \\ion{O}{5} feature is especially worthy of interest because its large oscillator strength ($f = 0.514$, greater than \\ion{H}{1} \\lya) should make it competitive in strength to the \\ion{O}{6} feature even at column densities down to $\\sim 10^{13}$ \\percm. Despite the fact that the FOS has rather low resolution by modern standards ($R \\approx 1300$), searching for \\ion{O}{5} rather than \\ion{O}{6} has the advantage that the features are in a spectral region where the \\lya\\ forest is less dense owing to the strong evolution of the density of lines with redshift. The detection of oxygen lines in the IGM can be used to infer the oxygen abundance, which is expected to be overabundant relative to carbon with respect to the solar ratio if the early universe is enriched by Type II supernovae.  Strong evidence for such abundance patterns are seen at high redshift in Lyman-limit systems \\citep{rvhe+92,krth+99}, damped \\lya\\ absorbers (\\citealt*{plh95}; \\citealt{lsbc+96}), and \\ion{C}{4} absorbers \\citep{dhhk+98,song98}.  Here we extend this work to the lower column density regime by carrying out a statistical search for absorption in the EUV using the presently available data.  In \\S\\ref{sec:data} we discuss the data and the spectral selection criteria that we use for this study.  We begin \\S\\ref{sec:analysis} by discussing the characterization of our \\lya\\ absorber sample from optical data.  This is followed by a presentation of the details of the technique that we use to generate the stacked absorption spectra and the method for measuring the equivalent width of the resulting absorption features.  In \\S\\ref{sec:results} we describe the results of our search for various samples selected by the strength of \\lya\\ absorption, particularly the detections of \\ion{O}{5}.  We consider the implications of our results in \\S\\ref{sec:discuss} by comparing the measurements and limits on the equivalent widths of the absorption features to simulated data to constrain the abundance ratios of the relevant ions to \\ion{H}{1}.  Using photoionization models, we then infer the abundance of oxygen in IGM.  We provide a brief summary of our results in \\S\\ref{sec:summary}. ",
        "conclusions": "}  \\subsection{The \\ion{O}{5} / \\ion{H}{1} Ratio\\label{sec:o5h1}}  The interpretation of composite spectra is not straightforward. In order to reliably interpret the measured equivalent widths, it is necessary to generate simulated data to compare to the results.  As a simplification, we first will assume that the mean ratio \\ovhi\\ is constant within each of our subsamples. This may or may not be a good assumption, depending on the ionization model, although our subsamples span a small enough range of \\nhi\\ that we do not expect \\ovhi\\ to vary by more that a factor of $\\sim 2$. Making this assumption allows us to derive results that are, as much as possible, model independent.  Later in this section we will drop this restriction and simulate data using particular ionization models.  Our Monte Carlo simulations proceed as follows: \\begin{enumerate} \\item{We randomize the redshift of each line to take on any value in the allowable range, as we did for our noise simulations in \\S\\ref{sec:results}.} \\item{We calculate the \\ion{O}{5} column density for each absorber from the \\ion{H}{1} column density and the given relative abundance \\ovhi.} \\item{The equivalent width of \\ion{O}{5} \\lam 630 is calculated from the \\ion{O}{5} column density assuming a particular $b$-parameter.} \\item{For the particular (randomized) redshift of each absorber, we insert a Gaussian absorption feature with a FWHM of 290 \\kms\\ and the appropriate equivalent width at the position of the \\ion{O}{5} \\lam 630 absorption.} \\item{The absorbers are combined into a composite and the equivalent width of \\ion{O}{5} \\lam 630 is measured as we have done for the real data.} \\end{enumerate}  An important assumption in this process is the $b$-parameter used to calculate the equivalent width.  For the weak lines in the low-column-density sample, this is fairly unimportant, since most of the lines will be reasonably near the linear portion of the curve of growth.  It will have some effect, though, and it is worth considering since this will have a more pronounced effect when we model lines at higher column densities. Individual metal lines as observed in \\ion{C}{4} absorption typically have quite small $b$-parameters.  As examples, \\citet{cskh95} find a median $b$ for \\ion{C}{4} lines of 10 \\kms, while \\citet{essp00} find a median $b$ of 13 \\kms.  However, when the \\ion{C}{4} absorption is strong, the absorption corresponding to particular \\lya\\ features often consists not of a single line but several such narrow absorption features. The effective $b$-parameter for such metal-line clusters, when viewed as a single absorption feature, is significantly larger than 10 \\kms. Because individual metal lines are not generally detected at the small \\ion{H}{1} column densities that we probe, we do not know the typical structure of the metal-line absorption. Strong \\ion{C}{4} absorbers often contain perhaps 3--5 components (see, for example, the \\ion{C}{4} profiles in \\citet{elpc+99}).  An effective $b$-parameter of 50 \\kms\\ should therefore provide a reasonable extreme value for the typical $b$-parameter of the metal absorption associated with a sample of \\ion{H}{1} lines. To illustrate the effect of $b$ on our results, we create simulated data for $b$-parameters of 10 and 50 \\kms.  A physically realistic picture is probably somewhere in between, since it seems likely that some absorbers will consist of several components whereas others may be dominated by a single strong component.  The simulations also depend on the scatter in \\ovhi\\ at fixed \\nhi, which depends on the scatter in both the metal abundances and physical conditions of the absorbers. Using \\ion{C}{4} absorption, \\citet{rhs97} and \\citet{hdhw+97} find a scatter of around an order of magnitude, while the analysis of \\citet{dhhk+98} suggests that the true scatter may be smaller as some of the observed scatter could be attributable to fitting uncertainties. The individual systems measured by \\citet{essp00} indicate an rms scatter of around 0.5 dex.  We generate simulated data with a Gaussian distribution in \\ovhi\\ with an rms scatter, which we denote as $\\sigma_{\\rm O}$, of 0.0, 0.5, and 1.0 dex.  The results of our simulations of sample L for different values of $b$ and $\\sigma_{\\rm O}$ are plotted in Figure~\\ref{fig:lowmodelo5}, compared with our observed value.  For each line in Figure~\\ref{fig:lowmodelo5} we have computed 100 simulated data sets for each point in steps of 0.1 in \\ovhi, and the plotted value represents of the mean of these simulations.  The $b$-parameter makes a small difference, but $\\sigma_{\\rm O}$ has a larger effect. If we assume that $\\sigma_{\\rm O}$ is likely in the range 0.5--1.0 and allowing for the uncertainty in $b$, we get agreement with our observations for \\ovhi$\\approx -1.7$ to $-0.8$.  Generating simulated datasets for sample I is not as simple. Because we do not have accurately measured column densities for these lines, we cannot simply calculate the \\ion{O}{5} column densities from the \\ion{H}{1} column densities of the lines in our sample.  Instead, we use simulated lists of random lines drawn from a column density distribution with $\\beta = -1.41$ for \\nhi $< 10^{14.3}$ \\percm\\ and $\\beta = -1.83$ for $10^{14.3} <$ \\nhi $< 10^{16.2}$ \\percm, as shown in Figure~\\ref{fig:cddf}.  We assume a $b$-parameter distribution as derived by \\citet{kity97}:  a Gaussian distribution in $b$ with $\\overline{b} = 23$ \\kms\\ and $\\sigma_b = 14$ \\kms, and a minimum $b$ given by $b_{\\rm min} = 14 + 4 \\times \\log [$\\nhi$/ 12.5]$ \\kms.  The distribution of $b$ is important because we have used a cut in equivalent width to define our sample, so we must know the equivalent width of the lines in the simulated data.  In our full sample, we have 69 \\lya\\ lines with $10^{14.2} <$\\nhi $< 10^{16.2}$ \\percm.  In Figure~\\ref{fig:ew} we plot a histogram of the equivalent width distribution of these lines, as well as the average of 1000 simulated lists of \\lya\\ lines in this column density range.  The notable difference between the two is the excess of large equivalent width lines, EW $\\gtrsim 0.8$ \\AA, in the real data.  These are likely metal-line systems, as discussed earlier, though they could also be blended lines.  The upper-limit in \\lya\\ equivalent width of 0.6 \\AA\\ for sample I is thus chosen to help assure that this sample is relatively free of metal-line systems. At lower equivalent widths, it seems that the simulated data are a good approximation of the real data.  In generating our simulated \\ion{O}{5} data for sample I, we replace our sample absorbers with randomly generated line lists.  Given that we now have lines with a realistic distribution of \\ion{H}{1} column densities, we can calculate the \\ion{O}{5} column densities and generate simulated data as we did for the lower column densities.  Figure~\\ref{fig:midmodelo5} shows the results of the simulations.  At these larger \\ion{O}{5} equivalent widths, $\\sigma_{\\rm O}$ makes less of a difference, but the results are quite sensitive to the assumed $b$-parameter, since for a small $b$-parameter the stronger lines will begin to saturate.  An effective $b$-parameter of 10 \\kms\\ is almost certainly unrealistically low for this sample, since \\ion{C}{4} absorption in this regime often shows structure, but it is useful to consider as an extreme case. Depending on the $b$-parameter, we find \\ovhi\\ should be in the range $-1.6$ to $-0.6$ for sample I.  This agrees with what we derived for sample L, giving us confidence in our analysis and suggesting that \\ovhi\\ does not vary more than an order of magnitude over the range of \\ion{H}{1} column densities that we have considered, from $10^{13.2}$ to $\\sim 10^{15}$ \\percm.  \\subsection{Other Ions\\label{sec:otherions}} To this point we have only analyzed \\ion{O}{5} since it is the only ion we detect in samples L and I.  However, in \\S\\ref{sec:otherlimits} we placed 90\\% confidence limits on the equivalent widths of absorption from other ions. By creating model data as we have done for \\ion{O}{5} \\lam 630, we can convert these 90\\% confidence upper limits on the equivalent widths of metal lines of other ions into constraints on the column density ratio of the corresponding ions to \\ion{H}{1}. For all of the ions we assume an rms scatter in \\xihi\\ of 0.5 dex for the simulated data. Table~\\ref{ta:otherions} lists our derived values and limits on \\xihi\\ for samples L and I.  We do not calculate values for sample H because our lack of knowledge of the \\ion{H}{1} column densities makes the results very uncertain.  \\begin{figurehere} \\centerline{\\psfig{file=f10.eps,angle=-90,width=9cm}} \\caption{Expected \\ion{O}{5} equivalent widths from simulated data for \\lya\\ absorbers with $10^{13.2} <$ \\nhi $< 10^{14.2}$ \\percm, as a function of the mean ratio \\ovhi.  The different curves correspond to different assumptions for the $b$-parameter for \\ion{O}{5} and the rms scatter in \\ovhi, denoted as $\\sigma_{\\rm O}$.  Our measured value is shown as the horizontal solid line, and the $1\\sigma$ limits on the measured value are shown as the horizontal dotted lines. \\label{fig:lowmodelo5}} \\end{figurehere} \\vspace{0.4cm}  \\begin{figurehere} \\centerline{\\psfig{file=f11.eps,angle=-90,width=9cm}} \\caption{{\\it Solid Line}:  Distribution of rest-frame equivalent widths of the 69 \\lya\\ absorption lines in our sample with \\ion{H}{1} column densities $> 10^{14.2}$ \\percm. {\\it Dotted Line}:  Average distribution of rest-frame equivalent widths from 1000 randomly-generated lists of 69 \\lya\\ lines with column densities $> 10^{14.2}$ \\percm, assuming the column density distribution shown in Figure~\\ref{fig:cddf} and a $b$-parameter distribution as derived by \\citet{kity97}.\\label{fig:ew}} \\end{figurehere} \\vspace{0.4cm}  \\subsection{The Oxygen Abundance\\label{sec:abund}}  Converting a value for the \\ovhi\\ into an oxygen abundance $\\langle {\\rm O} / {\\rm H} \\rangle$ requires knowledge of the relative ionization correction for the ions \\ion{O}{5} and \\ion{H}{1}.  This is generally ascertained by means of photoionization models in which the IGM is illuminated by the processed radiation of QSOs and/or stars.  We use the photoionization code CLOUDY \\citep[version 94.00;][]{ferl96}. Because we have not resolved individual \\ion{O}{5} absorption components, in using these models to infer $\\langle {\\rm O} / {\\rm H} \\rangle$ we must necessarily assume that the absorbers are predominantly single-phase; i.e., the \\ion{O}{5} and \\ion{H}{1} absorption occurs in the same gas. The main inputs to the models are the density and temperature of the absorbers as a function of \\ion{H}{1} column density, and the shape and normalization of the ionizing continuum. For the physical conditions in the absorbers we use the results of \\citet{hhkw98}, which are based on hydrodynamical simulations as described by \\citet*{kwh96}.  At $z=3$, \\citet{hhkw98} find that the typical density of the absorbers in their model are well-described by a simple analytic fit, $\\log n_{\\rm H} = -14.8 + 0.7 \\log \\,$\\nhi\\ assuming $\\Omega_b h^2 = 0.0125$.  The typical temperature of the absorbers is described by $\\log T=7.17 + 0.65 \\log n_{\\rm H}$ for $\\log n_{\\rm H} < -3.8$ and $\\log T=3.18 - 0.4 \\log n_{\\rm H}$ for $\\log n_{\\rm H} > -3.8$.  The temperature relation is somewhat dependent on the ionization history of the universe \\citep{hugn97}, particularly for underdense gas, but for the absorbers and redshifts we are studying this is a small effect and our results are fairly insensitive to the temperature.  \\begin{figurehere} \\centerline{\\psfig{file=f12.eps,angle=-90,width=9cm}} \\caption{Same as Figure~\\ref{fig:lowmodelo5} but for absorbers with \\nhi $> 10^{14.2}$ \\percm\\ and a \\lya\\ equivalent width $< 0.6$ \\AA.  In contrast to the weaker lines in Figure~\\ref{fig:lowmodelo5}, the scatter $\\sigma_{\\rm O}$ makes only a small difference but the result is more sensitive to the assumed $b$-parameter, since many of the \\ion{O}{5} features would begin to saturate if $b$ were small. \\label{fig:midmodelo5}} \\end{figurehere} \\vspace{0.4cm}  For our nominal ionizing spectrum we use the standard \\citet[hereafter HM96]{hama96} spectrum.  Specifically, we use their result for $z=2$, although there is very little evolution in this spectrum between $z\\sim 2$ and 3.  \\hm\\ assume that QSOs are the source of the ionizing radiation for the IGM. They assume a QSO spectral index of $\\alpha_{\\rm EUV}=-1.5$ ($f_\\nu \\propto \\nu^\\alpha$) for the continuum shortward of the Lyman limit, consistent with the $\\alpha_{\\rm EUV}=-1.57\\pm 0.17$ for radio-quiet QSOs derived by \\citet{tzkd02}.  Some recent hydrodynamical simulations of the IGM, however, suggest that a softer ionizing continuum than \\hm\\ is necessary to explain the observed ratio of \\ion{He}{2} to \\ion{H}{1} in the IGM.  \\citet{zanm97} require less \\ion{He}{2} ionizing radiation by a factor of $\\sim 4$, while \\citet{tlep+98} require a factor of $\\sim 2$ less.  \\citet{cwkh97} resolve the discrepancy by increasing $\\Omega_b h^2$ by an amount consistent with recent determinations of $\\Omega_b h^2$ from ${\\rm D}/{\\rm H}$ \\citep{otks+01} and the cosmic microwave background \\citep{dabb+02}. Because the ions of interest to our discussion, in particular \\ion{O}{5}, are created by photons with wavelengths shortward of the \\ion{He}{2} break, our results are sensitive to differences in the \\ion{He}{2} photoionization rate.  We create different ionizing spectra for our photoionization models by first creating a simple analytic fit to the \\hm\\ spectrum, which we find produces indistinguishable results for the ions of interest.  We then create softer spectra by simply increasing the strength of the break at the \\ion{He}{2} edge and then extending the softer spectrum to where it meets the \\hm\\ spectrum in the X-ray regime.  The original \\hm\\ spectrum and our model spectra are shown in Figure~\\ref{fig:ionspec}.  For all these models we adopt the normalization of the ionizing continuum of $J_\\nu = 10^{-21.3}$ \\jnu\\ at 1 Ryd as derived by \\hm\\ and in good agreement with \\citet{gcdf+96}.  \\begin{figurehere} \\centerline{\\psfig{file=f13.eps, angle=-90, width=9cm}} \\caption{Various ionizing spectral shapes considered in this work.  The solid line shows the spectrum of \\citet{hama96} at $z=2$.  The dotted line shows an approximation to this spectrum using just a few points.  These two spectra yield virtually indistinguishable results for the ions we consider.  The dashed lines show the same model spectrum, except we have artificially increased the break at the \\ion{He}{2} ionization edge by factors of two and four. \\label{fig:ionspec}} \\end{figurehere} \\vspace{0.4cm}  In Figure~\\ref{fig:modelic} we plot the ionized fractions of \\ion{O}{4}, \\ion{O}{5}, \\ion{O}{6}, and \\ion{C}{4} relative to \\ion{H}{1} for our photoionization models.  For the \\hm\\ spectrum, \\ovhi $-$\\oh\\ remains relatively constant for $10^{13} <$\\nhi $< 10^{14.8}$ \\percm\\ at around 3.6.  Given our values of \\ovhi\\ for sample L, this implies \\oh\\ around $-5.3$ to $-4.4$, or $[{\\rm O} / {\\rm H}]$ around $-2.2$ to $-1.3$ with respect to the standard solar oxygen abundance.  In the softer ionizing spectra models, slightly less oxygen is required to explain our measurement for sample L but slightly more oxygen for sample I. At larger \\ion{H}{1} column densities, \\ovhi\\ varies more dramatically with \\nhi\\ and the interpretation of our measured value of \\ovhi\\ over a range of \\nhi\\ is less obvious.  To include the effect of a changing ionization state with \\nhi, we create simulations using these specific ionization models to determine \\nov\\ for the simulated absorption where the free parameter is now \\oh\\ rather than \\ovhi.  We calculate model equivalent widths for \\ion{O}{4} \\lam 788 as well to compare with our limits since this can provide an interesting constraint on the ionization state.  In Figure~\\ref{fig:lowmodelo} we plot the equivalent widths of \\ion{O}{5} and \\ion{O}{4} as a function of \\oh\\ for sample L, assuming $b=50$ \\kms\\ and $\\sigma_{\\rm O} = 0.5$ dex.  We do not plot the results for different values of the $b$-parameter and $\\sigma_{\\rm O}$, but the effects are very similar to those seen in Figure~\\ref{fig:lowmodelo5}; i.e., the required \\ovhi\\ is more by around 0.1 for $b=10$ \\kms, and less by around 0.4 for $\\sigma_{\\rm O} = 1.0$ dex. As expected, the required \\oh\\ agrees with what we estimated above, with slightly less oxygen required for the softer spectra.  For none of these ionizing spectra does the required amount of oxygen imply that there should be a clearly observable feature for \\ion{O}{4}, although for the \\ion{He}{2} break $\\times$ 4 spectrum the expected equivalent width of \\ion{O}{4} \\lam 788 is around our 90\\% confidence limit. This comparison of \\ion{O}{5} to \\ion{O}{4} is virtually independent of our assumptions about the $b$-parameter and $\\sigma_{\\rm O}$ since the \\ion{O}{5} and \\ion{O}{4} equivalent widths essentially scale by the same factor for changes in these parameters.  \\begin{figurehere} \\centerline{\\psfig{file=f14.eps,width=9cm}} \\caption{Ionized fractions of various ions relative to \\ion{H}{1}.  The solid lines are for \\ion{O}{4}--\\ion{O}{6} and are labeled in the diagram.  The dotted line is for \\ion{C}{4}. The models correspond to the ionizing spectra shown in Figure~\\ref{fig:ionspec}.  A normalization of $J_{\\nu} = 10^{-21.3}$ \\jnu\\ at 1 Ryd is used for all three models. \\label{fig:modelic}} \\end{figurehere} \\vspace{0.4cm}  The \\ion{O}{4} limit is more interesting for sample I. The predicted equivalent widths are plotted in Figure~\\ref{fig:midmodelo_50} assuming $\\sigma_{\\rm O} = 0.5$ dex and $b=50$ \\kms, which as we argued before is probably more realistic than $b=10$ \\kms\\ for \\nhi $> 10^{14.2}$ \\percm. Again, the effects of changing the $b$-parameter and $\\sigma_{\\rm O}$ are similar to those in Figure~\\ref{fig:midmodelo5}. For these absorbers, the simulated data suggest that the observed lack of \\ion{O}{4} favors a hard ionizing spectrum. For the \\hm\\ spectrum the predicted equivalent width is only slightly higher than our 90\\% confidence limit. However, for the \\ion{He}{2} break $\\times$ 4 spectrum, the simulations suggest an equivalent width of \\ion{O}{4} of more than three times our limit.  This constraint is sensitive to the normalization of the ionizing continuum. An increase in the ionizing flux will decrease \\oivov\\ resulting in less predicted \\ion{O}{4}. We rerun the simulations using $J_\\nu = 10^{-21}$ \\jnu\\ at 1 Ryd  as found by \\citet*{cec97}. The resulting \\oivov\\ is smaller by around 30--40\\%, consistent with the \\hm\\ spectrum but still lying above our upper limit for the softer spectra. Density gradients in the absorbers can also influence the results.  By using a constant density, we are assuming that the absorption takes place primarily in the densest portion of the absorber. This should be a good approximation for a given ion provided that the volume density of the ion increases with total density.  However, sample I contains absorbers with \\ion{H}{1} densities near the point where \\ion{O}{5} peaks.  If a significant amount of the absorption were to occur in the density wings of the absorber, the observed \\oivov\\ would be less than what we have calculated for a constant-density model.  \\begin{figurehere} \\centerline{\\psfig{file=f15.eps,angle=-90,width=9cm}} \\caption{Expected equivalent widths of \\ion{O}{5} \\lam 630 and \\ion{O}{4} \\lam 788 as a function of \\oh\\ for \\lya\\ absorbers with $10^{13.2} <$ \\nhi $< 10^{14.2}$ \\percm.  The different lines correspond to different assumed ionization fractions, as shown in Figure~\\ref{fig:modelic}, based on different ionizing spectra.  For all the simulated data we assume an rms scatter in \\oh\\ of 0.5 dex and a $b$-parameter of 50 \\kms.  Our 90\\% confidence upper limit on the equivalent width of \\ion{O}{4} \\lam 788 is shown as the dotted line in the top panel.  Our measured value for \\ion{O}{5} is shown in the bottom panel as it was in Figure~\\ref{fig:lowmodelo5}. \\label{fig:lowmodelo}} \\end{figurehere} \\vspace{0.4cm}  \\begin{figurehere} \\centerline{\\psfig{file=f16.eps,angle=-90,width=9cm}} \\caption{Same as Figure~\\ref{fig:lowmodelo} for absorbers with \\nhi $> 10^{14.2}$ \\percm\\ and a \\lya\\ equivalent width $< 0.6$ \\AA. \\label{fig:midmodelo_50}} \\end{figurehere} \\vspace{0.4cm}  An additional consideration is our ignorance of the column densities of the lines in sample I. Because there are only 49 lines in the sample, it is possible that the true distribution of \\ion{H}{1} column densities is significantly different than our simulated lists.  In particular, there could be a relative lack in the true sample of high-column-density systems where \\oivov\\ is highest.  Our simulations suggest that the uncertainty in the predicted \\ion{O}{4} equivalent width due to statistical fluctuations in the column-density distribution is comparable to the random noise. The additional uncertainly for \\ion{O}{5} is much smaller since in our photoionization models \\nov\\ varies much less with \\nhi\\ over the range of \\ion{H}{1} column densities spanned by sample I. Models with predicted \\ion{O}{4} equivalent widths only moderately above our stated limit are therefore consistent with our results.  However, it is quite difficult to reconcile our limit on \\ion{O}{4} with the \\ion{He}{2} break $\\times$ 4 spectrum.  We conclude that our data are probably inconsistent with such a soft ionizing spectrum and that the true ionizing background is not significantly softer than that of \\hm.  As we noted above, early observational constraints on $N($\\ion{He}{2}$)/$\\nhi, often called $\\eta$, led us to consider ionizing continua softer than \\hm.  The strongest constraint on $\\eta$ generally considered in the literature is from the \\ion{He}{2} opacity measurement of \\citet*{dkz96}.  Recently, \\citet{ksoz+01} analyzed a high-resolution spectrum of the \\ion{He}{2} absorbing region with the Far Ultraviolet Spectroscopic Explorer enabling for the first time the explicit measurement of $\\eta$ for individual absorbers.  Using both measured values of $\\eta$ and lower limits for \\ion{He}{2} absorbers with no corresponding \\ion{H}{1} absorption, they find a mean $\\eta$ of $\\sim 80$.  However, the mean $\\eta$ for absorbers with detectable \\ion{H}{1} (\\nhi $\\gtrsim 10^{12.3}$ \\percm) is $\\sim 30$.  Because our sample is selected by \\ion{H}{1} absorption, it is appropriate to compare to this value. Comparing to the models of \\citet*{fgs98}, $\\eta = 30$ suggests a QSO source spectrum with $\\alpha_{\\rm EUV}$ around $-1.6$, consistent with \\citet{tzkd02} and similar to \\hm. Thus the fact that our data support a hard ionizing continuum is consistent with the \\ion{He}{2} results if one takes into account that an \\ion{H}{1}-selected sample preferentially selects absorbers with lower values of $\\eta$ which are likely photoionized by hard, QSO-like radiation.  Including uncertainty in the scatter, our models suggest \\oh\\ of around $-5.2$ to $-4.6$ for sample I, although the required \\oh\\ could be much higher if the effective $b$-parameter is much smaller than 50 \\kms.  This agrees with the results for sample L, for which the results are more robust due to the fact that the \\ion{H}{1} column densities of the constituent absorbers are directly measured. We therefore conclude that the typical \\oh\\ for the IGM likely lies in the range $-5.3$ to $-4.4$, as derived for sample L.  The abundance with respect to the standard solar value is thus $[{\\rm O}/{\\rm H}]$ around $-2.2$ to $-1.3$.  Having derived this abundance, it is interesting to compare this to results from studies of \\ion{C}{4}.  Using the optical depth ratio technique, \\citet{essp00} find that \\civhi $=-2.6$ for $\\tau_{\\rm Ly\\alpha} \\gtrsim 1$ provides a good match to their data, in agreement with what \\citet{soco96} found for individually detectable \\ion{C}{4} absorbers.  Using our \\hm\\ model for the ionization correction for \\ion{C}{4}, \\civhi $=-2.6$ corresponds to $[{\\rm C}/{\\rm H}] \\approx -2.5$.  This is in good agreement with what other authors have derived assuming a \\hm\\ spectrum \\citep{hdhw+97,dhhk+98}.  The inferred $[{\\rm O}/{\\rm C}]$ is in the range 0.3--1.2, assuming that the chemical composition does not change dramatically from $z\\sim 2$ to 3.5 since our results are for slightly lower redshift than those for \\ion{C}{4}.  Similar results are obtained for the \\ion{He}{2} break $\\times$ 2 spectrum. This model implies less carbon, $[{\\rm C}/{\\rm H}] \\approx -2.7$, to match the \\ion{C}{4} results, but also less oxygen to match our results, resulting in the same $[{\\rm O}/{\\rm C}]$. Thus, since the relative ionization corrections in our models for \\ion{C}{4} and \\ion{O}{5} vary little for the \\ion{H}{1} column densities of interest, the uncertainty in $[{\\rm O}/{\\rm C}]$ is dominated by our uncertainty in \\ovhi.  An overabundance of oxygen relative to carbon is expected if the IGM is enriched in metals by Type II supernovae.  \\citet{eagl+93} find overabundances of oxygen relative to carbon by factors of 3--5 in halo stars, consistent with the low end of the range we infer.  There is also evidence that this overabundance may increase with decreasing metallicity \\citep*{igr98}. A relative overabundance of oxygen is expected from Type II supernova yields \\citep{wowe95}.  However, large values of $[{\\rm O}/{\\rm C}]$ may require nucleosynthesis via pair instability supernovae \\citep{hewo02} from very massive stars ($m \\gtrsim 100 \\ M_\\odot$) in the early universe \\citep{adsl+01,scha02}. A bimodal initial mass function for Populations III stars, including a peak at $\\sim 100 \\ M_\\odot$, is predicted by the hydrodynamical simulations of \\citet{naum01}.  An additional predicted consequence of nucleosynthesis by very massive stars is a large abundance of silicon, $[{\\rm Si} / {\\rm C}] \\sim 1.0$.  In contrast, \\citet{soco96} infer a typical value of $[{\\rm Si} / {\\rm C}] \\sim 0.4$ for metal-line systems from the observed \\ion{Si}{4} / \\ion{C}{4} ratios. There is evidence for absorbers with $[{\\rm Si} / {\\rm C}] \\gtrsim 1.0$ at high redshift ($z>3$), but the assumed implausibility of such abundance ratios are used to argue instead for a substantial softening of the ionizing background \\citep{scdf+97,gish97,song98}.  Further exploration of the ionizing background at $z>3$, such as from measurements of the \\ion{He}{2} / \\ion{H}{1} and \\ion{O}{4} / \\ion{O}{5} ratios to be made possible by the Cosmic Origins Spectrograph, will help to clarify this issue.  Evidence for an overabundance of oxygen at high redshift has been found by other authors.  \\citet{rvhe+92} find $[{\\rm O}/{\\rm C}] \\gtrsim 0.6$ for Lyman-limit systems.  Using the optical depth technique, \\citet{dhhk+98} find a good fit for their data using $[{\\rm O}/{\\rm C}] = 0.5$ for \\ion{C}{4} absorbers at $z\\gtrsim 3$. However, based on the lack of narrow \\ion{O}{6} absorption lines in their data, they claim that the overall metallicity must drop dramatically, by a factor of at least $\\sim 3$, for absorbers with \\nhi $\\lesssim 10^{14.5}$ \\percm. In contrast, our results suggest that there is a large abundance of metals at \\nhi\\ $\\lesssim 10^{14.2}$ \\percm, consistent with what is found at higher \\ion{H}{1} column densities.  The \\ion{O}{6} detection of \\citet{srsk00} supports this result by suggesting that the enrichment of the IGM does indeed extend to densities below the universal mean.  It is instructive to compare our results to \\citet{srsk00}.  Assuming \\ovhi$=-1.1$ and using the \\hm\\ model for the ionization correction, for \\lya\\ optical depths ranging from 0.1 to 1.0, corresponding to $10^{13.6} \\lesssim$ \\nhi $\\lesssim 10^{14.6}$ \\percm\\ for $b=30$ \\kms, we would predict \\ovihi\\ from around $-0.5$ for $\\tau_{\\rm Ly \\alpha} = 0.1$ to $-1.3$ for $\\tau_{\\rm Ly \\alpha} = 1.0$.  The corresponding ratio of the optical depth of \\ion{O}{6} \\lam 1032 absorption to \\lya\\ would range from $-1.1$ to $-1.9$.  This agrees well with the apparent optical depths of \\ion{O}{6} measured by \\citet{srsk00} for the same redshift range, though as they point out, the true optical depth of \\ion{O}{6} could be considerably different from the apparent optical depth due to the effects of noise and Lyman-line contamination.  The existence of metals at low column densities addresses the issue of the source of the metals in the IGM.  The enrichment of the IGM is thought to occur in one of two basic ways:  (1) by a high-redshift era of small star-forming regions, either Population III stars \\citep{osgn96} or protogalaxies \\citep*{mfr01}, or (2) by in-situ star formation, either in the \\lya\\ absorbers themselves or nearby galaxies.  The primary distinction between the two scenarios is the volume extent of the enrichment.  While early supernovae could be expected to pre-enrich the universe in a fairly uniform way, it is difficult to enrich the most diffuse regions of the IGM by in-situ star formation. For example, \\citet{gnos97} predict a strong dropoff in metallicity at column densities below around $10^{13.5}$--$10^{14.5}$ \\percm\\ at $z=3$.  Although our sample L includes absorbers down to $10^{13.2}$ \\percm, we cannot be sure that there are metals at such low column densities, though the fact that the S/N of the detection does not decrease until we include absorbers with \\nhi $<10^{13.2}$ \\percm\\ certainly provides evidence that there are metals at low column density.  Our results apply to slightly lower redshifts than this prediction, so there is more time for metals to diffuse into less dense regions, and the absorbers that we observe correspond to slightly larger overdensities for the same \\nhi.  Given these considerations, we cannot conclusively claim that our results are inconsistent with in-situ enrichment scenarios, which are currently an active area of study (\\citealt*{fps00}; \\citealt{ahsk+01}). However, it is fair to say that our detection of \\ion{O}{5} in sample L, in conjuction with previous detections of \\ion{C}{4} and \\ion{O}{6} in similar density regimes, provide evidence for the presence of metals at quite low column densities, which would favor a uniform enrichment as provided by an early era of star formation.  } We have used {\\it HST} FOS spectra of four QSOs to search for absorption features associated with \\lya\\ forest absorbers in the EUV in the redshift range $1.6 < z < 2.9$. The results of this search can be summed up as follows: \\begin{enumerate} \\item{We detect \\ion{O}{5} \\lam 630 over a large range of \\nhi.  Most interestingly, we detect \\ion{O}{5} in a sample of absorbers with $10^{13.2} <$ \\nhi\\ $< 10^{14.2}$ \\percm\\ with greater than 99\\% confidence.} \\item{We detect \\ion{O}{4} \\lam 788, \\ion{O}{4} \\lam 554, \\ion{O}{3} \\lam 833, and \\ion{He}{1} \\lam 584 only for the strongest absorbers in our sample, those with \\lya\\ equivalent widths $\\gtrsim 0.6$ \\AA.} \\item{We find no evidence for \\ion{N}{4} \\lam 765, \\ion{Ne}{5} \\lam 568, \\ion{Ne}{6} \\lam 559, \\ion{Ne}{8} \\lamlam 770, 780, or \\ion{Mg}{10} \\lamlam 610, 625 absorption in any of our samples.} \\end{enumerate}  For absorbers with \\lya\\ equivalent widths $\\lesssim 0.6$ \\AA, the \\ion{O}{5} detections imply \\ovhi $\\approx -1.7$ to $-0.6$ based on our simulated data, where this range allows for uncertainties in the assumed $b$-parameter and scatter in \\ovhi.  This result implies that \\ion{O}{5} is typically a factor of $\\sim 10-100$ more abundant than \\ion{C}{4}.  The \\ion{O}{5} \\lam 630 line thus has EW$/\\lambda$ a factor of $\\sim 30-300$ larger than \\ion{C}{4} \\lam 1548, or ranging in strength from around that of Ly$\\delta$ to \\lyb, making it an excellent tracer of metal content.  Using photoionization models to calculate the ionization correction, we find that the oxygen abundance in the IGM is $[{\\rm O}/{\\rm H}] \\approx -2.2$ to $-1.3$, implying $[{\\rm O}/{\\rm C}]\\approx 0.3$ to $1.2$, consistent with what is found in halo stars and Lyman-limit systems. The overabundance of oxygen suggests Type II supernova enrichment, but an unusual stellar initial mass function resulting in a significant contribution from pair instability supernovae of very massive ($m \\gtrsim 100 \\ M_\\odot$) Population III stars may be necessary.  The fact that we find no evidence for \\ion{O}{4} \\lam 788 absorption except in the strongest systems provides an interesting constraint on the ionizing background spectrum, specifically that it is unlikely to be more than a factor of $\\sim 2$ softer than the \\hm\\ spectrum at 4 Ryd.  A hard ionizing spectrum is consistent with measurements of \\ion{He}{2} absorption for absorbers with detectable \\ion{H}{1}.  We conclude that studying the absorption of the IGM in the rest-frame EUV is not only a useful but essential tool for gaining a more complete understanding of the metal content and ionization of the IGM.  From longer wavelength data alone one cannot place useful constraints on absorption from multiple ionization stages of the same element in the diffuse IGM.  Ultimately, a more complete understanding requires such analysis, and the clearest path to this end is through the simultaneous study of multiple ionization stages of oxygen (\\ion{O}{3}--\\ion{O}{6}).  We look forward to the installment of the Cosmic Origins Spectrograph on the {\\it HST}, currently planned for 2004, which will enormously increase our ability to observe weak oxygen features on an individual basis."
    },
    "0207/astro-ph0207071_arXiv.txt": {
        "abstract": "The cosmological implications from a new estimate of the local X-ray galaxy cluster abundance are summarized. The results are then compared to independent observations. It is suggested that `low' values for the mean cosmic matter density and the amplitude of mass density fluctuations currently do not appear unreasonable observationally. ",
        "introduction": "A new X-ray selected and X-ray flux-limited galaxy cluster sample has been constructed (\\gcs, the HIghest X-ray FLUx Galaxy Cluster Sample, Reiprich \\& B\\\"ohringer 2002). Based on the ROSAT All-Sky Survey the 63 brightest clusters with galactic latitude $\\vert b_{\\rm II} \\vert \\geq 20$\\,deg and flux $f_{\\rm X}(0.1-2.4\\,{\\rm keV})\\ge 2\\times10^{-11}\\,{\\rm ergs\\,s^{-1}\\,cm^{-2}}$ have been compiled. Gravitational masses have been determined utilizing intracluster gas density profiles, derived mainly from ROSAT PSPC pointed observations, and gas temperatures, as published mainly from ASCA observations, assuming hydrostatic equilibrium. This sample and an extended sample of 106 galaxy clusters has been used to establish the X-ray luminosity--gravitational mass relation.  From the complete sample and the individually determined masses the galaxy cluster mass function has been determined and used to constrain the mean cosmic matter density and the amplitude of mass density fluctuations. Comparison to Press--Schechter type model mass functions in the framework of cold dark matter cosmological models and a Harrison-Zeldovich initial density fluctuation spectrum yields the constraints $\\Omega_{\\rm m}=0.12^{+0.06}_{-0.04}$ and $\\sigma_8=0.96^{+0.15}_{-0.12}$ (90\\% c.l.). The degeneracy between $\\Omega_{\\rm m}$ and $\\sigma_8$ previously encountered for local cluster samples therefore has been broken mainly due to the large covered mass range (Fig.~1; see section 5.3.2 in Reiprich \\& B\\\"ohringer 2002 for more details). Various possible systematic uncertainties have been quantified. Adding all identified systematic uncertainties to the statistical uncertainty in a worst case fashion results in an upper limit $\\Omega_{\\rm m}<0.31$. For comparison to previous results a relation $\\sigma_8=0.43\\,\\Omega_{\\rm m}^{-0.38}$ has been derived.  Two further constraints on $\\Omega_{\\rm m}$ obtained from the \\gcs\\ clusters agree well with the above results. The mean intracluster gas fraction combined with independent estimates of the baryon density yields the upper limit $\\Omega_{\\rm m}\\la 0.34$. Calculation of the median mass-to-light ratio for 18 clusters in common to the sample of Girardi et al.\\ (2000) combined with estimates of the total luminosity density in the Universe yields $\\Omega_{\\rm m}\\approx 0.15$ (Reiprich 2001).  The mass function has been integrated to show that the contribution of mass bound within virialized cluster regions to the total matter density is small; i.e., $\\Omega_{\\rm cluster}=0.012^{+0.003}_{-0.004}$ for cluster masses larger than $6.4^{+0.7}_{-0.6}\\times 10^{13}\\,h_{50}^{-1}\\,M_{\\odot}$. If light traces mass this also implies that most galaxies sit outside clusters. \\begin{figure} \\plottwo{bfmf.ps}{banana.ps} \\caption{{\\bf Left:} \\gcs\\ differential mass function compared to the best fit model mass function with $\\Omega_{\\rm m}=0.12$ and $\\sigma_8=0.96$ (solid line). Also shown are the best fit model mass functions for fixed $\\Omega_{\\rm m}=0.5$ ($\\Rightarrow\\sigma_8=0.60$; dashed line) and $\\Omega_{\\rm m}=1.0$ ($\\Rightarrow\\sigma_8=0.46$; dotted line). {\\bf Right:} Statistical confidence contours. The cross indicates the minimum and ellipses the 68\\%, 90\\%, 95\\%, and 99\\% c.l.\\ for two interesting parameters.} \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207592_arXiv.txt": {
        "abstract": "Two sets of observational carbon stellar yields for low-and-intermediate mass stars are computed based on planetary nebula abundances  derived from C II $\\lambda4267$ and C III $\\lambda\\lambda1906+1909$ lines, respectively. These observational yields are assumed in chemical evolution models for the solar vicinity and the Galactic disk. C/O values observed in stars of the solar vicinity and Galactic \\HII regions are compared with those predicted by chemical evolution models for the Galaxy. I conclude that the C yields derived from permitted lines are in better agreement with the observational constraints than those derived from forbidden lines. ",
        "introduction": "\\label{sec:intro}  The $N({\\rm C}^{++})/N({\\rm H}^+)$ values derived from the optical recombination line (permitted line, PL) C II $\\lambda4267$ are higher, by as much as a factor of 10, than those determined from the collisionally excited lines (forbidden lines, FL) C III $\\lambda \\lambda1906+1909$ (eg. Rola \\& Stasinska 1994, Peimbert, Luridiana \\& Torres-Peimbert 1995a, Peimbert, Torres-Peimbert \\&  Luridiana 1995b, Liu et al. 2001, Luo, Liu \\& Barlow 2001). Several explanations for this discrepancy have been presented in the literature, (see the reviews by Liu 2002, Peimbert 2002, Torres-Peimbert \\& Peimbert 2002) but the problem remains open. Since PNe are important for the C enrichment of the interstellar medium, a successful chemical evolution model for the solar vicinity and the Galactic disk (Carigi 2000) is used to discriminate between the PN-C abundances derived from permitted lines and the PN-C abundances obtained from forbidden lines. Hereafter, all abundances are given  by number.  This work is based on a preliminary study presented by Carigi (2002). ",
        "conclusions": "\\label{sec:conclusion}  From chemical evolution models of the Galaxy, I conclude that:  a) Models with the permitted line yields (C$_{\\rm PN}^{\\rm PL}$) match all the observational constraints, in particular  they reproduce the C/O absolute values observed in dwarf stars of the solar vicinity and in \\HII regions of the Galactic disk.  b) Models with the forbidden line yields (C$_{\\rm PN}^{\\rm FL}$) fail to  reproduce the C/O ratios in dwarf stars of different ages  in the solar vicinity, the Sun, and the inner \\HII region M17.  c) Models with C$_{\\rm PN}^{\\rm PL}$ yields agree with models based on theoretical yields, in particular showing better agreement with models based on the Padova yields  than models based on the Amsterdam yields.  d) The C/O values predicted with ${\\rm C_{PN}^{\\rm PL}}$ yields are about 0.08 dex higher  than those obtained with the Amsterdam yields, and about 0.03 dex lower than those computed with the Padova yields.  e) The C/O values predicted with ${\\rm C_{PN}^{\\rm FL}}$ yields are about 0.10 dex lower than those obtained with the Amsterdam yields, and about 0.20 dex lower than those obtained with the Padova yields.  f) The C$_{\\rm PN}^{\\rm PL}$ yields should increase with $Z$ to obtain a better agreement between models and observations in the C/O versus $t$ and  C/O versus O/H diagrams.  g) The C/O increase with $Z$ is governed by the metallicity dependent yields of both massive stars and LIMS. Massive stars determine the behaviour of C/O with $Z$ in the early and late evolution, while LIMS do so in the middle evolution.  h) The C/O gradient steepens with time, but when the gas acquires supersolar abundances, the gradient flattens with time."
    },
    "0207/astro-ph0207247_arXiv.txt": {
        "abstract": "We present the detailed analysis of {\\sl Hubble Space Telescope} observations of the spatial distributions of different stellar species in two young compact star clusters in the Large Magellanic Cloud, NGC 1805 and NGC 1818.  Based on a comparison of the characteristic relaxation times in their cores and at their half-mass radii with the observed degree of mass segregation, it is most likely that significant primordial mass segregation was present in both clusters, particularly in NGC 1805. Both clusters were likely formed with very similar initial mass functions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207521_arXiv.txt": {
        "abstract": "The processes that disperse the products of massive stars from their birth sites play a fundamental role in determining the observed abundances.  I discuss parameterizations for element dispersal and their roles in chemical evolution, with an emphasis on understanding present-day dispersion and homogeneity in metallicity.  Turbulence dominates mixing processes, with characteristic timescales of order $10^8$ yr, implying significant dilution of metals into the ISM.  This permits a rough estimate of the metallicity distribution function of enrichment events.  Many systems, including the Milky Way and nearby galaxies, show metallicity dispersions that as yet appear consistent with pure inhomogeneous evolution. There are also systems like I~Zw~18 that show strong homogenization, perhaps tied to small galaxy size, high star formation rate, and/or superwinds. ",
        "introduction": "The dispersal of the nucleosynthetic products of massive stars plays a fundamental role in galactic chemical evolution.  Our ability to decipher the chemical signatures of star formation history depend upon understanding the dispersal processes.  The growing database and improving accuracy in elemental abundance determinations increasingly is revealing the magnitude of abundance dispersions, and setting limits on uniformity.  It is apparent that the magnitude of these dispersions results directly from the balance between interstellar transport processes and the timescale for global star formation. This field of study is presently still in its infancy, but I will discuss here some of our rudimentary knowledge.  Dispersal of newly-synthesized elements takes place on three levels: {\\bf 1.} Local {\\bf mixing} with the ambient interstellar medium (ISM); {\\bf 2.} Global {\\bf homogenization} of a system or subsystem; {\\bf 3.  Outflow / inflow} of metals from the considered system.  The local mixing processes describe the immediate dispersal of elements associated with individual star-forming regions into areas of roughly uniform metallicity.  Homogenization of the system then describes the mixing of these individual patches toward globally uniform abundances.  Finally, outflows and inflows consider the transport of metals beyond the system altogether.  Other presentations in this volume (e.g., Strickland; Matteucci) consider the last; here, I will consider only the first two processes, which apply only within a given system. ",
        "conclusions": "In summary, the dispersal of newly synthesized elements takes place on at least three scales:  localized mixing, global homogenization, and outflow/inflow from the system.  It has emerged that the transport mechanisms are dominated by turbulence.  The dispersal length scale for localized mixing determines the dilution of new metals in the ISM, and therefore sets the MDF of the parent enrichment units $f(Z)$ that drive chemical evolution.  We crudely estimate that, depending on the balance of ISM temperature phases, the hot, metal-bearing gas will have cooling times of order $10^8$ yr and disperse to roughly a decade in distance beyond the original radius of the superbubble created by the SNe. This is consistent with analytic and numerical hydrodynamical simulations of mixing by de Avillez \\& Mac Low (2002).  An important constraint on the minimum values for $f(Z)$ should be given by the observed low-metallicity threshold; we find that the current observed limit of [Fe/H]$\\sim -4$ for halo stars (Beers 1999) is compatible with our crude, order-of-magnitude estimate.  This rough understanding of element dispersal and corresponding estimate for the $f(Z)$ of individual enrichment units can then be incorporated into Simple Inhomogeneous Models for chemical evolution. These yield estimates for the present-day dispersion in metallicity, which may be compared with observations.  We estimate dispersions of 0.1 -- 0.3 dex in the metallicity range $0.1 - 1.0 Z_\\odot$, which is thus far consistent with data for the solar neighborhood, Local Group galaxies, and several nearby starburst galaxies, thereby suggesting that global homogenization is thus far unnecessary to explain observations.  However, the scatter in abundance is expected to increase at lower metallicities.  While this is seen in Milky Way stellar abundances, the extremely metal-poor galaxy I~Zw~18 shows metallicities that are more uniform than expected from the SIM.  This implies that this galaxy has a faster homogenization time relative to its star formation duty cycle, perhaps caused by its small size, high star formation rate, and/or outflow of hot, metal-bearing gas.  It is evident that our understanding of these processes are still rudimentary, and more data and modeling are necessary."
    },
    "0207/astro-ph0207184_arXiv.txt": {
        "abstract": "The  detection of spectral variability of the $\\gamma$-ray blazar Mrk~421 at TeV energies is reported.  Observations with the Whipple Observatory 10~m $\\gamma$-ray telescope taken in 2000/2001 revealed exceptionally strong and long-lasting flaring activity.  Flaring levels of 0.4 to 13 times that of the Crab Nebula flux provided  sufficient statistics for a detailed study of the energy spectrum  between 380~GeV and 8.2~TeV as a function of flux level. These spectra are well described by a power law with an exponential cutoff: $\\rm \\: \\:  \\: {{dN}\\over{dE}} \\propto \\: E^{-\\alpha}\\times e^{-E/E_{0}} \\: \\: m^{-2} \\: s^{-1} \\: TeV^{-1} $.  There is no evidence for variation in the cutoff energy with flux, and all spectra are consistent with an average value for the cutoff energy of 4.3 TeV. The spectral index varies between $\\rm 1.89\\pm 0.04_{stat} \\pm 0.05_{syst} $ in a high flux state and $\\rm  2.72 \\pm0.11_{stat} \\pm 0.05_{syst} $ in a low state.  The correlation between spectral index and flux is tight when averaging over the total 2000/2001 data set. Spectral measurements of Mrk~421 from previous years (1995/96 and 1999) by the Whipple collaboration are consistent with this flux-spectral index correlation, which suggest that this  may be a constant or a long-term property of the source. If a similar flux-spectral index correlation were found for other $\\gamma$-ray blazars, this universal property could help disentangle the intrinsic emission mechanism from external absorption effects. ",
        "introduction": "The discovery of more than 70 active galactic nuclei (AGNs) by the EGRET $\\gamma$-ray detector (Hartman et al. 1999) operating at $\\rm E > 30$~MeV gave a fresh perspective on the AGN phenomenon, particularly relevant to understanding  the intrinsic properties of their jets. EGRET-detected AGNs are typically radio-loud and  show a second peak in their $\\rm \\nu F_{\\nu}$ distribution at GeV energies. Blazars detected at TeV energies have a primary peak at X-ray energies and a second component at TeV energies. Both types are $\\rm \\gamma$-ray blazars and the commonly-accepted model is that they  have their  jet oriented towards the observer revealing emission regions that  are strongly Doppler-boosted. Relativistic boosting gives rise to large flux variations (Catanese et al. 1997) and short time scale phenomena (Gaidos et al. 1996).   Two AGNs (Mrk~421 and Mrk~501) show emission extending to energies greater than 10 TeV (Aharonian et al. 1999; Krennrich et al. 2001).    Since the discovery of TeV  $\\gamma$-rays from the blazars Mrk~421 (Punch et al. 1992) and Mrk~501 (Quinn et al. 1996), these objects  played a significant role in  discussions involving  the emission processes in AGN jets and attenuation effects of TeV $\\gamma$-rays propagating over extragalactic distances. Both blazars exhibit episodes of strong flaring activity, providing good statistics for detailed measurements of their average energy spectra from 260~GeV up 17~TeV using ground-based $\\gamma$-ray telescopes. Mrk~421 and Mrk~501 are at approximately the same distance (z=0.031  and z=0.034, respectively). Since  the level of attenuation of $\\gamma$-rays by the diffuse extragalactic background light (EBL) via pair creation (Nikishov 1962; Gould \\& Schr\\`eder 1967; Stecker, De Jager \\& Salamon 1992)  depends on the distance of the source to the observer, it could cause a common spectral feature in the energy spectra of  Mrk~421 and Mrk~501.    Measurements by the Whipple collaboration (Samuelson et al. 1998; Krennrich et al. 2001) imply that the energy spectra of both Mrk~501 and Mrk~421 require a curved fit parametrization, e.g., a power law with an exponential cutoff  with cutoff energies of $\\rm 4.6 \\pm 0.8_{stat}$~TeV   and $\\rm 4.3 \\pm 0.3_{stat} (-1.4 +1.7)_{syst}$~TeV (``stat'' means statistical error, ``syst'' means systematic error), respectively. Data from the HEGRA collaboration suggest that the cutoff energy of Mrk~501 is $\\rm 6.2 \\pm 0.4_{stat} (-1.5 +2.9)_{syst} $~TeV (Aharonian et al. 1999; Aharonian et al. 2001) and that Mrk~421 has a cutoff energy of $\\rm 4.2 \\:  (+0.6 -0.4)_{stat}$~TeV (Kohnle et al. 2001). The results from  the two  groups may be consistent given systematic uncertainties.  The interpretation of the origin of the cutoff requires a better understanding of the emission process in $\\rm \\gamma$-ray blazars as a class of objects. Making progress in unraveling the emission process in blazars from  external attenuation effects requires two types of key observations.  The first is contemporaneous multiwavelength observations at X-ray and TeV energies (see Buckley et al. 1996; Krawczynski et al. 2002). These constrain the emission process by simultaneously probing the synchrotron emission and the mechanism responsible for the $\\rm \\gamma$-ray  component, e.g., the inverse Compton scattering process (Maraschi, Ghisellini \\& Celotti 1992; Marscher \\& Travis 1996). There are also competing models that assume that the high energy emission arises from protons (Mannheim 1998; M\\\"ucke \\& Protheroe 2001; Aharonian 2000).  The second is the study of $\\gamma$-ray spectral variability as a function of flux and time.  Spectral variability is directly tied to the emission process at  the source, whereas external absorption by the EBL is a universal feature that is independent of flux level. Thus, by studying the  spectral variability, we are able to disentangle these two effects.   In  synchrotron-Compton models, spectral variability could be explained by cooling of electrons in the jet causing a shift of the break in the electron spectrum, or by variations of the maximum energy of accelerated electrons.  Spectral variability may also be a  key to the understanding of the energy dissipation processes in the vicinity of a supermassive black hole powering the jet.       Some evidence for spectral variability has been reported by Djannati-Ata\\\"{\\i} et al. (1999) and Krawczynski et al. (2001) for Mrk~501; however, the effect was not highly statistically significant ($\\rm \\approx$~4 standard deviations)  and precluded detailed studies. In fact, the spectral index of Mrk~501 turned out to be surprisingly stable during observations of a strong outburst in 1997 (Aharonian et al. 1999) though the flux varied by a factor of 30.  In this paper we present the discovery of spectral variability of Mrk~421 at TeV energies based on data taken with the Whipple Observatory 10~m $\\rm \\gamma$-ray telescope. We show that Mrk~421 exhibits a remarkable flux-spectral index correlation that appears to be stable  averaged over time-scales  of months to several years. ",
        "conclusions": "Strong and extended flaring of Mrk 421 allows us to study the spectral variability of this source as a function of flux. Averaged over the time period         2000 November 28  to  2001 April 13 the data show a clear flux-spectral index correlation, with the spectral  index varying between $\\rm 1.89\\pm 0.04$  in a high state and $\\rm  2.72 \\pm0.11$  in a low state.  Whether or not this correlation is maintained in different epochs outside of the 2000/2001 observing period can be addressed   using previously published results and  archival data from the Whipple collaboration.   Figure 3 also shows results  from the average spectrum in 1995/1996 in a high state (Krennrich et al. 1999a) and  the average spectrum  from 1999 May 6 through June 8. The  data points  from 1995/1996   and from 1999  fall  into place with the flux-spectral index  correlation as observed for the  2000/2001 data alone. This indicates that  the   correlation between spectral index and flux  holds true when averaging over time scales of months to five years.     Spectral hardening during flares has also been observed for Mrk~421 in X-rays by Fossati et al. (2000)  using BeppoSAX data during X-ray flares in 1997 and 1998.  In  X-rays, the  effect of spectral hardening  has been interpreted as a shift of the  synchrotron peak towards higher frequencies.    The flux-spectral index correlation seen in the TeV spectra  could also be interpreted as  a shift of the high energy peak towards higher frequencies. The shape of the spectrum~I in Figure 1 suggests that the peak in  $\\rm \\nu F_{\\nu}$ is at a few hundred~GeV, significantly above previous levels (Maraschi et al. 1999). The spectral hardening is most evident at energies below 2~TeV; it is not uniform.  Strong spectral hardening at lower energies might be expected in the inverse-Compton (IC) scenario in which the IC peak shifts towards higher energies as the flux increases.  Conversely, at the higher energies, spectral softening occurs due to either a terminating particle distribution in energy,  the  falling cross-section due to the Klein-Nishina effect, or  external attenuation effects from  the EBL or nearby radiation fields.    If the changes are due to a shifting IC peak energy, the flux value would be closely tied to the spectral index, as seen here. Further studies of spectral variability on short time-scales (hours-days) will be presented in a follow-up paper."
    },
    "0207/astro-ph0207651_arXiv.txt": {
        "abstract": "We have obtained spectroscopy with the Far Ultraviolet Spectroscopic Explorer (FUSE) of the supersoft X-ray binary RX~J0513.9$-$6951 over a complete binary orbital cycle.  The spectra show a hot continuum with extremely broad O VI emission and weak Lyman absorptions.  He II emission is weak and narrow, while N III and C III emissions are undetected, although lines from these ions are prominent at optical wavelengths.  The broad O VI emission and Lyman absorption show radial velocity curves that are approximately antiphased and have semiamplitudes of $\\sim117\\pm40$ and $54\\pm10$ km s$^{-1}$, respectively.  Narrow emissions from He II and O VI show small velocity variations with phasing different from the broad O VI, but consistent with the optical line peaks.  We also measure considerable changes in the FUV continuum and O VI emission line flux.  We discuss the possible causes of the measured variations and a tentative binary interpretation. ",
        "introduction": "In this paper we present FUSE observations of the supersoft X-ray binary RX~J0513.9$-$6951 (hereafter called X0513$-$69).  This $ROSAT$ source was identified with a $\\sim$16th mag peculiar emission-line star in the Large Magellanic Cloud (Pakull et al. 1993, Cowley et al. 1993).  With an absolute magnitude of M$_V\\sim-2.0$, X0513$-$69 is the brightest of the supersoft X-ray binaries (e.g. Cowley et al. 1998).  It is known to show high and low optical states differing by $\\sim$1 mag (e.g. Alcock et al. 1996, Cowley et al. 2002).  The optical spectrum is characterized by strong, broad emission lines of He II and weaker lines of O VI, N V, C IV, and C III.  Weak emissions flanking the strongest lines at $\\pm\\sim$4000 km s$^{-1}$ have been interpreted as arising in bi-polar jets (e.g. Crampton et al. 1996).  Previous spectroscopic work has shown that the emission line peaks have very small velocity amplitude (e.g. Crampton et al. 1996, Southwell et al. 1996).  If these are interpreted as motion of the compact star, the source must be seen at a low orbital inclination and the unseen secondary star must have a low mass.  The small amplitude of the optical light curve (Alcock et al. 1996) also suggests the system is viewed from a low inclination angle.  However, using a newly refined orbital period and ephemeris, Cowley et al.\\ (2002) have shown that the phasing of the velocities derived from optical spectra differs between the high and low optical states, so that the previously inferred mass function and its interpretation are in doubt.  In order to access other highly ionized lines (particularly the O VI resonance doublet) and investigate the far-ultraviolet continuum, observations were obtained with the Far Ultraviolet Spectroscopic Explorer (FUSE).  The FUSE data presented in this paper add new information about X0513$-$69. ",
        "conclusions": ""
    },
    "0207/astro-ph0207616_arXiv.txt": {
        "abstract": "Gravitational lensing is a powerful tool to detect compact matter on very different mass scales. Of particular importance is the fact that lensing is sensitive to both luminous and dark matter alike. Depending on the mass scale, all lensing effects are used in the search for matter: offset in position, image distortion, magnification, and multiple images. Gravitational lens detections cover three main mass ranges: roughly stellar mass, galaxy mass and galaxy cluster mass scales, i.e. well known classes of objects. Various searches based on different techniques explored the frequency of compact objects over more than 15 orders of magnitude, so far mostly providing null results  in mass ranges different from the ones just mentioned. In particular, no population of ``compact dark objects'' could be detected so far. Combined, the lensing results offer some interesting limits on the cosmological frequency of compact objects in the mass interval $10^{-3} \\le M/M_\\odot \\le 10^{15}$, unfortunately still with some gaps in between. In the near future, further studies along these lines promise to fill the gaps and to push the limits further down, or they might even detect new object classes. ",
        "introduction": "The basic setup for  a gravitational lens scenario is displayed in Figure \\ref{fig-setup}. The three ingredients in such a lensing situation  are the source S, the lens L, and the observer O. Light rays emitted from the source are deflected by the lens. For a point-like lens, there will always be (at least) two images S$_1$ and S$_2$  of the source.  With external shear -- due to the tidal field of objects outside but near the light bundles -- there can be more images. The observer sees the images in directions corresponding to the tangents of the incoming light paths.   In Figure \\ref{fig-setup} the corresponding angles and angular diameter distances $D_L$, $D_S$, $D_{LS}$ are indicated. In the thin-lens approximation, the hyperbolic paths are approximated by their asymptotes. In the circular-symmetric case, the deflection angle is given as  \\begin{equation} \\label{eq-angle} \\tilde \\alpha (\\xi)  =  { { 4 G M(\\xi)} \\over {c^2} } { 1  \\over \\xi }. \\end{equation} where $M(\\xi)$ is the mass inside a radius $\\xi$. In this depiction the origin is chosen at the observer. From the diagram it can be seen that the following relation holds: \\begin{equation} \\theta D_S = \\beta D_S + \\tilde \\alpha  D_{LS} \\end{equation} (for $\\theta$,  $\\beta$, $\\tilde \\alpha \\ll 1$;  this condition is fulfilled in practically all astrophysically relevant situations). With the definition of the reduced deflection angle as $\\alpha (\\theta)  =   ( D_{LS} / D_{S} ) \\tilde \\alpha (\\theta)$, this can be expressed as: \\begin{equation} \\label{eq-lenseq} \\beta =   \\theta - \\alpha (\\theta). \\end{equation} This relation between the positions of images and source can easily be derived for a non-symmetric mass distribution as well. In that case all angles are vector-valued. The two-dimensional {\\bf lens equation} then reads: \\begin{equation} \\label{eq-lenseq-vec} \\vec\\beta =   \\vec\\theta - \\vec\\alpha (\\vec\\theta). \\end{equation}       For a point lens of mass $M$, the deflection angle is given by equation (\\ref{eq-angle}). Plugging this into equation (\\ref{eq-lenseq}) and using the relation $\\xi =   D_L \\theta $ (cf. Figure \\ref{fig-setup}) one obtains: \\begin{equation} \\beta(\\theta)  = \\theta - { D_{LS} \\over D_L D_S} { 4 G M \\over c^2 \\theta}. \\end{equation} For the special case in which the source lies exactly behind the lens ($\\beta = 0$), due to the symmetry a ring-like image occurs whose angular radius is called {\\bf Einstein radius} $\\theta_E$:  \\begin{equation} \\theta_E  = \\sqrt{ { 4 G M \\over c^2} {D_{LS} \\over D_L D_S}. } \\end{equation} The Einstein radius defines the angular scale for a lens situation. For a massive galaxy with a mass of $M = 10^{12} M_{\\odot}$ at a redshift of $z_L = 0.5$ and a source at redshift $z_S = 2.0$ (we used here a Hubble constant of $H = 50 $km sec$^{-1}$ Mpc$^{-1}$ and an Einstein-deSitter universe), the Einstein radius is \\begin{equation} \\label{eq-angle-gal} \\theta_{E, \\ \\rm galaxy}  \\approx 1.8  \\ \\sqrt{ M \\over 10^{12} M_\\odot } \\ {\\rm arcsec} \\end{equation} (note that for cosmological distances in general $D_{LS} \\ne D_S - D_L$!).  For a galactic microlensing scenario in which stars in the disk of the Milky Way act as lenses for bulge stars close to the center of the Milky Way, the scale defined by the Einstein radius is \\begin{equation} \\theta_{E, \\ \\rm galactic\\ star}  \\approx 0.5  \\ \\sqrt{ M \\over M_{\\odot} }   \\ {\\rm milliarcsec}. \\end{equation} Time scales for galactic microlensing events -- i.e. the duration for crossing the Einstein radius -- typically range from weeks to months. For cosmological/quasar microlensing, this time scale extends to years; however, caustic crossing events can be as short as a few weeks.  More detailed introductions to gravitational lensing including some historic aspects can be found in \\cite{livrev}, or in the textbook \\cite{SEF} by Schneider et al. (1992) and in the more mathematically oriented monograph \\cite{PLW} by Petters et al. (2001).   \\begin{figure}[tb] \\plotone{jkw_fig_lensing_geometry.ps}{100mm} \\caption{\\label{fig-setup} a) Setup of a gravitational lens situation: The lens $L$ located between source $S$ and observer $O$ produces two images $S_1$ and $S_2$ of the background source. Relations between the various angles and distances involved in the lensing setup can be derived for the case $\\tilde \\alpha \\ll 1$, as formulated in the lens equation (\\ref{eq-lenseq}). } \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207420_arXiv.txt": {
        "abstract": "We use a simple analytic model to compute the angular correlation function of clusters identified in upcoming thermal SZ effect surveys. We then compute the expected fraction of close pairs of clusters on the sky that are also close along the line of sight. We show how the expected number of cluster pairs as a function of redshift is sensitive to the assumed biasing relation between the cluster and the mass distribution. We find that, in a $\\Lambda$CDM model, the fraction of physically associated pairs is $70\\%$ for angular separations smaller than $20$ arcmin and clusters with specific flux difference larger than $200$ mJy at $143$ GHz. The agreement of our analytic results with the Hubble volume $N$-body simulations is satisfactory.  These results quantify the feasibility of using SZ surveys to compile catalogues of superclusters at any redshifts. ",
        "introduction": "Superclusters are regions with average mass overdensities larger than a few, on scales larger than a few megaparsecs;  their cluster members show indications of  intense dynamical activity and provide evidence that structures form hierarchically on supercluster scales (e.g. \\citealt*{bardelli01}; \\citealt{rines01}; \\citealt{plionis02}).  Analyses of catalogues of superclusters in the local Universe (e.g. \\citealt{bahcall84}; \\citealt{west89}; \\citealt{zucca93}; \\citealt{einasto01}), beyond the local supercluster \\citep{oort83}, have focused on their shapes and their large scale distribution \\citep{kerscher98}. Shape statistics have been used for discriminating among the cosmological models (\\citealt*{basi01}; \\citealt*{kolo02}), whereas the total mass of superclusters can be used to estimate the density parameter $\\Omega_0$ from the mass-to-light ratio on large scales \\citep{small98}. At high redshift, superclusters, when they exist \\citep{postman02}, are difficult to identify, although preferred directions of radio emission from distant quasars and radio galaxies \\citep{west91} or unusual X-ray morphologies \\citep{rosati02} can provide evidence of their existence at $z\\ga 1$.  The peculiar velocities of clusters are enhanced in superclusters by non-linear effects (\\citealt{zucca93}; \\citealt*{bahcall94}). \\citet{colb00} have indeed shown that the peculiar velocity of dark matter halos with massive neighbours in $N$-body simulations is $\\sim 40-50\\%$ larger than predicted by linear theory. In fact, linear theory predicts that the evolution of the peculiar velocities is independent of the local density, whereas \\citet{sd01} have shown that, in $N$-body simulations, evolved peculiar velocities of halos with the same mass are larger (smaller) in high (low) density regions. Moreover, the high velocity tail of the cluster peculiar velocity distribution can be a sensitive discriminator of cosmological models (\\citealt{bahcall94}; \\citealt{peel02}). In $N$-body simulations of representative volumes of the Universe, high density regions of a few megaparsec size contain two or more massive clusters; these superclusters can be easily spot in thermal \\citet{SZ80} effect surveys as nearby CMB temperature decrements in the Rayleigh-Jeans limit and are the favorite regions for searches of enhanced kinematic SZ effect \\citep*{diaferio00}. Confusion caused by primary Cosmic Microwave Background (CMB) anisotropies, thermal SZ and instrumental noise, complicates the estimation of cluster peculiar velocities from measurements of the kinematic effect \\citep*{aghanim01}.  The kinematic SZ effect is difficult to detect directly; it is roughly an order of magnitude smaller than the thermal effect. One may proceed by looking for kinematic fluctuations in clusters selected according  to  thermal SZ measurements.  However the most massive clusters that are responsible for the largest thermal fluctuations do not necessarily have  the largest peculiar velocities.   The thermal SZ effect, being independent of redshift, is the optimal tool for detecting clusters at high redshifts \\citep{carls02}. Several SZ surveys will produce catalogues of thousands clusters at any redshift in the next few years, e.g. SPT \\citep{carls02}, AMI \\citep{kneissl01}, AMiBA \\citep{lo02}, to mention a few.  Clusters are not randomly distributed on the sky: their angular correlation function mirrors the clustering evolution of mass. The angular correlation function of SZ clusters was first computed by \\citet{cole88} who were interested in the non-Gaussianity of CMB temperature fluctuations caused by galaxy clusters. \\citet{kom99} show that the correlation in the cluster distribution enhances by $\\sim 20-30\\%$ the Poisson contribution to the CMB angular power spectrum at degree angular scales. \\citet{mosca02} compute the correlation function on the light-cone of clusters detectable by the Planck surveyor satellite and show that the correlation length of these SZ selected clusters is more sensitive to the physical properties of the ICM than to the cosmological parameters.  In this paper we use a simple analytic model to compute the angular correlation function and the fraction of pairs that are close along the line of sight as well as on the sky.  Our result shows that for viable cosmological models, this fraction is substantial. Therefore if we define superclusters as regions containing two or more close clusters, this result motivates the use of thermal SZ cluster surveys to identify superclusters at any redshift. We also compute the expected number count of cluster {\\it pairs} within a given separation as a function of redshift. Unlike the number count of individual clusters, the pair number count is sensitive to the assumed form of the bias function between the cluster and the mass distribution.  In section \\ref{sec:count} we outline the basic equations; in section \\ref{sec:stat} we compute the expected number of clusters per solid angle in flux limited thermal SZ surveys, their angular correlation function and the number of physically associated pairs; we then apply our model to a full sky survey like the Planck surveyor. In section \\ref{sec:bias} we show how the fraction of physically associated pairs depend on the assumed biasing relation.  We conclude in section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  One of the main purposes of SZ cluster surveys is to constrain the cosmological model with the cluster redshift distribution. In fact, the redshift distribution of SZ selected clusters depends on the density parameter $\\Omega_0$, the mass power spectrum, and the physical properties of the ICM.  Therefore one hopes that catalogues of SZ clusters can constrain the cosmological models despite the degeneracies between some parameters such as the normalization of the power spectrum $\\sigma_8$ and $\\Omega_0$ (\\citealt{barbosa96}; \\citealt{kneissl01}; \\citealt*{holder01}; \\citealt*{benson02}).  However, for the very same reason why SZ cluster surveys are so attractive, their catalogues will not have any accurate information on the cluster distances, although the morphology of the cluster can provide a rough estimate of its redshift \\citep{diego02}. The redshift distribution of clusters can be determined with follow up optical observations, CO line measurement in galaxy cluster members, or high resolution X-ray spectroscopy (e.g. \\citealt{bleeker02}). The number of objects in SZ cluster surveys depends on the sky coverage, instrumental sensitivity, and observational bands.  In an all sky surveys, this number can be of order several thousands.  Of course, only the redshifts of a subsample of these clusters is sufficient to provide a robust estimate of the redshift distribution. How should one choose this cluster subsample? We suggest to measure the redshift of clusters in close pairs. Because the probability that clusters close on the sky are also close along the line of sight is substantial, this strategy will provide, at the same time, the redshift distribution of clusters and a catalogue of superclusters.  To show that the compilation of a catalogue of superclusters from SZ cluster surveys is indeed feasible, we have computed the angular correlation function of SZ selected clusters for two cosmological models. We have then computed how many cluster pairs, close on the sky, are also close along the line of sight.  For example, for clusters with specific flux difference larger than $200$ mJy at $143$ GHz in a $\\Lambda$CDM model, $70\\%$ of the pairs with angular separations smaller than $20$ arcmin are physical associations. Because an SZ supercluster catalogue will contain superclusters at any redshifts thanks to the redshift independence of the SZ effect, the catalogue will be of great relevance for follow up studies of the evolution of galaxy and cosmic structures in high density regions. Moreover, from the redshift distribution of physically associated cluster pairs, we can extract information on the biasing relation between the cluster and the mass distribution.  To compute the number of detectable clusters, we have used the point source approximation, where the survey can detect the total cluster radio flux, but it cannot resolve the cluster structure. This approximation is thus independent of the cluster profile. It is realistic for poor resolution surveys (for example the Planck surveyor will perform a survey with a beam FWHM of 8 arcmin at 143 GHz), but of course unrealistic for high resolution bolometric or interferometric ground-based surveys. However, resolving the clusters has the net effect of increasing the mass threshold and thus decreasing the number of detected sources \\citep{bart00}. In other words, resolving the cluster structure is similar to considering our results for shallower surveys: the cluster minimum mass is larger and the fraction of physically associated pairs increases."
    },
    "0207/astro-ph0207566_arXiv.txt": {
        "abstract": "{ Recent experiments have shown that it is possible to study a fundamental astrophysical process such as dynamo action in controlled  laboratory conditions using simple MHD flows. In this paper we explore the possibility that Taylor-Couette flow, already proposed as a model of the magneto-rotational instability of accretion discs, can sustain generation of magnetic field. Firstly, by solving the kinematic dynamo problem, we identify the region of parameter space where the magnetic field's growth rate is higher. Secondly, by solving simultaneously the coupled nonlinear equations which govern velocity field and magnetic field, we find a self-consistent nonlinearly saturated dynamo. ",
        "introduction": " ",
        "conclusions": "By solving the kinematic dynamo problem, we have determined that a Taylor-vortex flow pattern can sustain a growing magnetic field.  Like in the models of \\cite{dudley89} we also find that the dynamo is sensitive to the flow pattern.  Further, for flows that are capable dynamo action we see that the growth rate is not a monotonic increasing function of the Reynolds number.  This is not seen in \\cite{dudley89}, most likely due to the prescribed form for the driving flow patterns.  In the Taylor-vortex flow the best growth rates have been obtained with co-rotation. The relative magnitude of the shear and roll in the flow plays an important part in the success of the dynamo mechanism.  Solving the full MHD equations we have demonstrated the existence of a fully self-consistent nonlinearly saturated dynamo. Hopefully these results will stimulate experimental work on the problem. Future theoretical work will investigate dynamo action in hydrodynamically stable flows and address the nature of the magnetic field structure when the dynamo is driven harder -- our dynamo is laminar. Most of the present work is concerned with wider gaps."
    },
    "0207/astro-ph0207085_arXiv.txt": {
        "abstract": "We have identified a population of `\\bfs ' (or `blanks') among the {\\em ROSAT} bright unidentified X-ray sources with faint optical counterparts. The extreme X-ray over optical flux ratio of blanks is not compatible with the main classes of X-ray emitters  except for extreme BL Lacertae objects.\\\\ From the analysis of {\\em ROSAT} archival data we found no indication of variability and evidence for only three sources, out of 16, needing absorption in excess of the Galactic value. We also found evidence for an extended nature for only one of the 5 blanks with a serendipitous HRI detection; this source (1WGA~J1226.9$+$3332) was confirmed as a z=0.89 cluster of galaxies. Palomar images reveal  the presence of a red ($O-E \\geq$ 2) counterpart in the X-ray error circle for 6  blanks. The identification process brought to the discovery of another high z cluster of galaxies, one (possibly extreme) BL Lac, two ultraluminous X-ray sources in nearby galaxies and two apparently normal type~1 AGNs. These AGNs, together with 4 more AGN-like objects  seem to form a well defined group: they present unabsorbed X-ray spectra but red Palomar counterparts. We discuss the possible explanations for the discrepancy between the  X-ray and optical data, among which: a suppressed big blue bump emission, an extreme  dust to gas ($\\sim 40-60$ the Galactic ratio), a high redshift ($z \\geq 3.5$) QSO nature,  an atypical dust grain size distribution and a dusty warm absorber. These AGN-like blanks seem to be the bright (and easier to study) analogs of the sources which are found in deep {\\em Chandra} observations. Three more blanks have a still unknown nature. ",
        "introduction": "The X-ray sky is not as well known as sometimes thought. There might exist classes of quite common sources, comprising from one to a few percent of the  high Galactic latitude population, which are currently thought to be ``rare'' because of the difficulty of finding them. We are actively searching for such `minority' populations (Kim \\& Elvis 1999). We have found an interesting high fraction of extreme X-ray loud sources ($\\sim 7-8$\\%) among the {\\em ROSAT} high Galactic latitude sources at all fluxes: (this paper, Bade et al. 1995; Schmidt et al. 1998). Here we designate as  `blank field sources' or `blanks' all the bright {\\em ROSAT} PSPC (Position Sensitive Proportional Counter) X-ray sources ($f_X > 10^{-13}$ erg cm$^{-2}$ s$^{-1}$) with no optical counterpart on the Palomar Observatory Sky Survey (POSS) (to O=21.5\\footnote{O-band effective wavelength = 4100 \\AA ; passband = 1100 \\AA}) within their $39^{\\prime \\prime}$ (99\\%) radius error circle. For comparison, at these X-ray fluxes a normal type 1 AGN should appear on the POSS some 3.5 magnitudes brighter ($O \\sim 18$). The nature of these sources has never been systematically investigated before, and their nature is still mysterious. We decided to select these sources and to study them to understand their nature  because of the important cosmological and astrophysical consequences that could derive from their identification.\\\\ The outline of this paper is the following: we introduce  {\\em ROSAT} blank field sources, review the open possibilities for their nature and the important consequences that could derive from the identification of these sources  in \\S ~2; we present the selection criteria used for the definition of the sample in \\S ~3, while \\S ~4, 6 and 7 focus on the analysis of {\\em ROSAT}, {\\em ASCA} and radio archival data. \\S ~7 contains detailed information on  the sources and  the discussion, while in \\S ~8 we compare the results of this work with the findings of other surveys. A summary is presented in \\S ~9. We will use H$_0$=75~km~s$^{-1}$~Mpc$^{-1}$ and $q_0 = 0.5$; errors   represent 1$\\sigma$ confidence levels, unless explicitly stated otherwise. ",
        "conclusions": "\\label{blanks_summary}   We have identified a population of \\bfs \\/  among the {\\em ROSAT} bright X-ray sources with faint optical counterparts, i.e. $O>21.5$ on the Palomar Sky survey. Their extreme X-ray over optical flux ratio (\\fxv $> 10$) is not compatible with the main classes of X-ray emitters, except for BL Lacs for the less extreme cases. The identification process brought to the discovery of two high z, high luminosity clusters  of galaxies (\\S ~\\ref{blanks_clusters}), one BL Lacertae object (\\S~\\ref{1WGAJ134012743}), and two type~1 AGNs (\\S ~\\ref{blanks_agns}). Four  blanks are  AGN-like sources (\\S ~\\ref{blanks_agns}) which seem to form a well defined group: they present  type~1 X-ray spectra and red Palomar counterparts. These sources are similar to the galaxies with bright X-ray emission ($L_X \\sim 10^{42} - 10^{43}$ erg s$^{-1}$) but weak or absent AGN features in the optical band found since {\\em Einstein} observations (Elvis et al. 1981; Tananbaum et al. 1997) and which are  now being found in large numbers (e.g. Della Ceca et al. 2001; Hornschemeier et al. 2001, Fiore et al. 2001). We considered possible explanations for the discrepancy between  X-ray and optical data: a suppressed BBB emission; % an extreme % dust to gas ratio; a dust grain size distribution different from the Galactic one; a dusty warm absorber and  an high redshift ($z \\geq 3.5$) QSO nature. Two sources (\\S ~\\ref{blanks_ulx}) are candidate ultraluminous X-ray binaries within nearby galaxies (see also Cagnoni et al. 2002). Three sources (\\S ~\\ref{blanks_individual_unknown}) have a still unknown nature; for each of them we listed and justified the possibilities excluded.\\\\  To make progress in understanding the nature of blanks we need to identify them with optical or near infrared counterparts. Since both in the case of obscured AGNs and of high redshift clusters of galaxies red counterparts are expected, we obtained optical (R band) and infrared (K band) imaging for the 16 fields in our sample at a 4m-class telescope and we will present the results in a future paper.\\\\ The obvious next step is to construct a larger and statistically complete sample of \\bfs . For this purpose, the {\\em XMM-Newton} and {\\em Chandra} catalogs soon available will form a good basis. The smaller positional uncertainty will cut down on false optical identifications and so will considerably {\\em increase} the percentage of blanks found ($\\sim 14$\\% at $f_X \\geq 10^{-13}$ \\cgs ,  Maccacaro \\& Della Ceca private communication). The smaller PSF of these new detectors will also allow the direct separation of extended sources, enabling high redshift clusters of galaxies to be found even more readily."
    },
    "0207/hep-ph0207323_arXiv.txt": {
        "abstract": "s{I will give here an overview of the present observational and theoretical situation regarding the question of the matter-antimatter asymmetry of the universe and the related question of the existence of antimatter on a cosmological scale. I will also give a simple discussion of the role of $CP$ violation in this subject.} ",
        "introduction": "One of the most fundamental questions in cosmology is that of the role of antimatter in the universe. This question, which is intimately connected to the question of the nature of $CP$ violation at high energies, is the important subject of this conference. It is a question for theorists, but ultimately as in all scientific endevours, a question which must be answered empirically if possible.  The discovery of the Dirac equation \\cite{di28} placed antimatter on an equal footing with matter in physics and opened up speculation as to whether there is an overall balance between the amount of matter and the amount of antimatter in the universe. The hot big bang model of the universe added a new aspect to this question. It became apparent that in a hot early epoch of the big bang there would exist a fully mixed dense state of matter and antimatter in the form of leptonic and baryonic pairs in thermal equilibrium with radiation. As the universe expanded and cooled this situation would result in an almost complete annihilation of both matter and antimatter.  The amount of matter and antimatter expected to be left over in an expanding universe can be calculated from the proton-antiproton annihilation cross section. Antinucleons ``freeze out'' of thermal equilibrium when the annihilation rate becomes smaller than the expansion rate of the universe. This would have occurred when the temperature of the universe dropped below $\\sim 20$ MeV. The predicted freeze out density of both matter and antimatter is only about $4\\times 10^{-11}$ of the closure density of the universe ({\\it i.e.}, $\\Omega_{baryon} = 4\\times 10^{-11}$) \\cite{ch66}.  On the other hand, big-bang nucleosynthesis calculations \\cite{ol02} and studies of the anisotropy of the 2.7 K cosmic background radiation \\cite{si02} have indicated that baryonic matter makes up about 4\\% of the closure density of the universe, ({\\it i.e.}, $\\Omega_{baryon} \\simeq 4 \\times 10^{-2}$) as shown in Figure \\ref{bbn}. Thus, there is a nine order of magnitude difference between the simple big-bang prediction and the reality of the amount of baryonic matter which is found in the universe and which makes up the visible matter in galaxies as well as the matter in you and me. Clearly, there is something missing. It was elegantly shown by Sakharov \\cite{sa66} that what is missing is the breaking of symmetries. In order to make an omelet you have to break some eggs; in order to make a universe you have to break some symmetries. It is in this context that the question of the nature of the violation of $CP$ symmetry arises.  \\begin{figure} \\centerline{\\psfig{figure=bbn.eps,height=3.5in}} \\caption{Predicted abundances of light nuclides from big-bang nucleosynthesis. \\protect\\cite{bu99} \\label{bbn}} \\end{figure} ",
        "conclusions": "The fact that simple hot big bang ``freeze out'' calculations predict a baryon density in the universe which is nine orders of magnitude too low indicates that $B$, $C$ and $CP$ symmetries must be broken in the early universe at times corresponding to a temperature greater than 20 MeV, the simple freeze-out temperature for nucleons and antinucleons. The violation of these symmetries, especially $CP$, and their consequences for cosmology are the subject of this meeting. If $CP$ violation is predetermined, than only matter will remain in the present universe. If, on the other hand, $CP$ violation ($CPV$) is the result of spontaneous symmetry breaking, domains of positive and negative $CPV$ may result. If this domain structure is stretched to astronomical size by a subsequent period of moderate inflation, then fossil baryons may survive as galaxies in some regions of the universe and fossil antibaryons survive as antigalaxies in other regions. We have referred to this possibility as ``locally asymmetric domain cosmology (LADC).'' A longer period of inflation would result in the entire visible universe being in one domain region.  As of this writing, there is no evidence for large scale extragalactic antimatter and, by inference, for LADC. Cosmologically significant sub-galaxy size antimatter regions are ruled out by their potential effect on big-bang nucleosynthesis. \\cite{ku00} Significant antimatter in our own galaxy is ruled out by low energy cosmic ray measurements. Although presently unclear, \\gray background measurements indicate that in a LADC cosmology, the size of the separate regions of matter and antimatter must be at least of galaxy supercluster extent.  However, in the present search for cosmological antimatter, absence of evidence is not necessarily evidence of absence. A dedicated MeV background satellite detector experiment designed to be as clean from radiation induced intrinsic contamination as possible would help to clarify the situation. A possible determination of departures from isotropy at 20 MeV by the GLAST (Gamma Ray Large Area Telescope) satellite, to be launched in 2006, may provide another test. However, this test is compromised by the real possibility that the 20 MeV background may be dominated by unresolved blazars \\cite{st96}. Another interesting test would be to look for departures from isotropy in the cosmic background radiation caused by the interactions of high energy electrons from the decay of $\\pi^{\\pm}$'s produced by annihilation at the boundaries of matter and antimatter regions. \\cite{ki97}"
    },
    "0207/astro-ph0207200_arXiv.txt": {
        "abstract": "We  determine the   evolution  of the   faint, high-redshift,  optical luminosity function of Active Galactic Nuclei (AGN) implied by several observationally-motivated models of the  ionizing background from $3 < z  < 5$.  Our  results depend crucially  on  whether we  use the total ionizing rate measured by the proximity effect  technique or the lower determination favored by the flux decrement distribution of Ly$\\alpha$ forest lines.  Assuming a faint-end luminosity  function slope of 1.58 and  the SDSS estimates of the  bright-end slope and normalization, we find that the luminosity function must break at $M_B^* = -24.2, -22.3, -20.8$ at $z=3, 4, 5$ if we adopt the lower ionization rate and assume no  stellar contribution to the background.   The breaks must occur at $M_B^*= -20.6, -18.7, -18.7$ for the proximity effect estimate.  Since stars  may also contribute  to the background,  these values are lower limits  on the  break luminosity, and   they brighten  by  as much  as $\\sim2$ mag if the  escape fraction  of  ionizing photons from  high-z galaxies is consistent  with recent estimates: $f_{\\rm esc}=0.16$.  By comparing our expectations to faint AGN searches in the HDF and high-z galaxy  fields, we   find  that   typically-quoted proximity    effect estimates of the background imply an over-abundance of AGN compared to the faint counts (even with $f_{\\rm esc}=1$).  Even adopting the lower bound  on proximity effect measurements,   the stellar escape fraction must be  high:   $f_{\\rm  esc}  \\gsim 0.2$.   Conversely,   the  lower flux-decrement-derived   background  requires   a  smaller  number  of ionizing  sources, and   faint  AGN counts  are   consistent with this estimate only if there is a limited stellar contribution, $f_{\\rm esc} \\lsim  0.05$.   Our derived  luminosity   functions together with  the locally-estimated black  hole  density suggest that the  efficiency of converting mass to light in optically-unobscured AGN is somewhat lower than expected,  $\\epsilon \\lsim 0.05$  (all  models).  Comparison with similar estimates  based on X-ray  counts suggests that more than half of all  AGN  are obscured in   the UV/optical.  We   also derive lower limits on typical AGN lifetimes  and obtain $\\gsim 10^7$yr for favored cases. ",
        "introduction": "Among the long-standing goals in extragalactic astronomy is to explain and characterize the   population  of Active Galactic    Nuclei (AGN). Their  large luminosities,  compact sizes, and  association with radio jets  have lead to  the assumption that AGN   are powered by accretion onto supermassive  black  holes (Salpeter 1964; Zel'dovich  \\& Novikov 1964;  Lynden-Bell    1969).  Although   this framework     provides a theoretical starting point,    there are many questions that    remain largely unresolved.  These  include    explaining the origin   of  the central  black  holes (e.g.   Eisenstein \\&  Loeb  1995; Madau \\& Rees 2001; Koushiappas, Bullock, \\&  Dekel 2002), understanding the fueling process,  lifetime, and efficiency of   the central engine (see,  e.g. Rees 1984;  Koratkar \\& Blaes  1999), and, ultimately, determining how quasar activity fits within   our  cosmological theory for   structure formation.  For many years, the role of AGN in structure formation was believed to be  that of a tracer  population,  important in  their own right,  but cosmologically interesting mainly  for  their contribution to  the  UV ionizing background   (and in their ability  to  track the collapse of structure).  Recent indications  have changed  this view dramatically. It now seems likely  that AGN play an important  role in the formation of galaxies.  In a reversal of sorts, this paper  focuses on using the observed ionizing  emissivity at high-redshift  in  order to constrain the evolution of  the AGN luminosity  function.  Derived in this  way, our luminosity functions relate  directly to the  long-standing desire to   pinpoint the dominant  ionizing   sources  in the  Universe,  and additionally help constrain models that  attempt to explain AGN within a cosmological context.  The AGN  luminosity  function  has  long  served as  a  benchmark  for understanding  the  formation and   evolution of quasi-stellar objects (QSOs)\\footnote{We use   the terms  AGN  and QSO  interchangeably.  In common parlance,  a QSO is  a high luminosity AGN ($M_B\\lesssim-23$).} (Efstathiou  \\& Rees  1988;  Carlberg  1990;  Haehnelt  \\&  Rees 1993; Cavaliere, Perri \\& Vittorini 1997; Haiman \\&  Loeb 1998; Richstone et al.~1998, Haehnelt, Natarajan,  \\&  Rees 1998; Cattaneo, Haehnelt,  \\& Rees 1999;  Haiman,  Madau, \\& Loeb  1999, Kauffman  \\& Haehnelt 2000; Haiman \\&  Hui 2001; Haehnelt   \\& Kauffman 2001;  Steed, Weinberg, \\& Miralda-Escud{\\'{e}} 2002).  Recent indications  that AGN activity  is linked in a fundamental way with the  formation of galaxies make these studies  all the  more   relevant  (Heckman  et  al.~1984; Sanders  et al.~1988; Kormendy  \\& Richstone 1995;  Sanders \\& Mirabel 1996; Boyle \\&  Terlevich 1998;  Dickinson  et  al.~1998;  Magorrian et  al.~1998; Richstone  et al.~1998; Laor  1998; Wandel 1999;   van der Marel 1999; Franceschini  et al.~1999;  Mathur 2000;  Canalizo  \\& Stockton  2001; Levenson, Weaver, \\& Heckman 2001; Ferrarese 2002).  Specifically, the relation between black   hole   mass and  bulge   velocity  dispersion (Gebhardt et al.~2000a,  ~2000b; Ferrarese \\&  Merritt 2000; Ferrarese et al.~2001)   is so tight   that it  seems impossible  to  understand without some significant cross-talk  (in the form of feedback) between the AGN phase and the formation of the galaxy and its stellar bulge.   Modelers attempting to understand these  relations have been forced to test and refine their assumptions by comparing to the low-redshift AGN luminosity function, or   to bright  quasar  counts at  high-redshift, because the  faint  population of  AGN is relatively  unconstrained at high-$z$.  This lack  of knowledge about low-luminosity objects allows considerable, unwanted freedom for model builders. For example, Haiman \\& Loeb (1998) have explored  the idea that faint  AGN are linked in a simple way with low-mass cold  dark matter (CDM) halos.  This predicts a  large number of   low-luminosity systems at  high-$z$  because less massive halos are   relatively   abundant at  early   times.   Another possibility is   that AGN activity  in  small halos is   suppressed by feedback processes.\\footnote{ For   example, the binding energy  of  a galaxy of mass $f_b M$ in a halo of  mass $M$ and circular velocity $V \\propto M^{1/3}$ will  scale  as $E_{\\rm   gal} \\simeq 0.5  f_b  M V^2 \\propto M^{5/3}$.  Compare this to the energy released by a black hole of mass $M_{\\bullet} \\propto M$ shining for a time  $t_{\\rm agn}$ at a fixed fraction $\\lambda$  of its Eddington luminosity $L_{Edd} \\propto M_{\\bullet}$:  $E_{\\rm agn} =  \\lambda  L_{E} t_{\\rm agn}  \\propto M$. This gives $E_{\\rm agn}/E_{\\rm gal} \\propto M^{-2/3}$, suggesting that AGN feedback  should  be  more important for   low-mass  halos.  If we insert  appropriate  numbers, we find  that  the energy  released by a bright AGN  over its  lifetime  should be  comparable  to  the binding energy of a galaxy-sized halo: $E_{\\rm agn} \\simeq 1.5 \\times 10^{16}$ $M_{\\odot}$      km$^2$    s$^{-2}$ (M$_{\\bullet}/10^{8}$ $M_{\\odot}$) ($t_{\\rm agn}/10^{7}$ yr) ($\\lambda/0.1$), while $E_{\\rm gal} \\simeq 2 \\times 10^{15}$ $M_{\\odot}$  km$^2$ s$^{-2}$ (M/10$^{12}$ $M_{\\odot}$) (V/200  kms$^{-1}$)$^2$ (f/0.1).  Although  just how this energy might manifest itself as a suppression  mechanism is unclear, the energetics suggest that a significant amount of feedback is plausible.}  A recent example of  this idea is explored  in Kauffman \\& Haehnelt (2000), who utilize a qualitatively plausible feedback  scheme to model black hole properties  and to match the observed  evolution in AGN number density from $z \\sim 0$ to $z \\sim 2$.  Because their model relies on feedback that scales with host population  and redshift, a high-$z$  constraint on the number of dim objects would serve as a useful test.  The  AGN   luminosity   function  (LF)   is   typically    written  as $\\phi(L,z)dL$, and is defined   as  the number  of objects  per   unit comoving volume  at  redshift  $z$,  with luminosity  between  $L$ and $L+dL$.  In the optical, the majority of studies use B magnitudes.  So unless otherwise stated, $L$ will denote  B band luminosity throughout this paper.  For low-redshift  AGN,  $\\phi$ is  well represented  by a broken power law  \\beq       \\phi(L,z)       =          \\frac        {\\phi_*/L_*       } {(L/L_*)^{\\gamma_f}+(L/L_*)^{\\gamma_b}}, \\label{eq:diff_phi} \\eeq  \\noindent which has a break at luminosity $L_*$, a characteristic number density $\\phi_*$,  and  asymptotic faint and bright  slopes  of $\\gamma_f$ and $\\gamma_b$  respectively.    As  will  be discussed  in    \\S2, out to $z\\sim2.5$ the AGN LF seems to evolve only in luminosity, in the sense that $L_*$  gets  brighter   with  increasing $z$,  while    the other parameters  stay fixed ($\\gamma_f \\simeq  1.6$, $\\gamma_b \\simeq 3.4$, and $\\phi_* \\simeq 10^3$Gpc$^{-3}$).  This kind  of evolution is known as pure luminosity evolution (PLE), and  the current best-estimate for $L_*(z)$  under this assumption has   it rising dramatically from  its local value  of $\\sim   10^{11}L_{\\odot}$  at $z=0$  to nearly   $\\sim 10^{13}L_{\\odot}$ at  $z = 2.5$ (Boyle,  Shanks, \\& Peterson 1988; Koo \\& Kron 1988;   Hewett et al.~1993;  Pei  1995a;  Boyle  et al.~2000). Beyond this redshift, only the brightest quasars have been observed in significant numbers, and there is as of yet no evidence for a break in the luminosity function.   In terms of the  double power-law (1),  the high-$z$  observations   only measure the  bright-end  slope $\\gamma_b \\simeq 2.6-2.9$ (Schmidt, Schneider  \\& Gunn 1995, hereafter  SSG; Fan et al.~2001a) and fix an overall integrated normalization that roughly imposes  a constraint on the  quantity $\\phi_*  L_*^{\\gamma_b-1}$ as a function  of z (see  \\S 2).  Most  interestingly, the space density of the observable bright quasars falls steadily from $z \\sim 3$ out to $z \\sim 6$ (Warren, Hewett, \\& Osmer 1994 [WHO94]; Kennefick, Djorgovski, \\& de Carvalho 1995 [KDC95]; SSG; Fan  et al.~2001a), but this decline cannot be  faithfully represented in terms  of PLE,  as the data shows $\\gamma_b$ is flatter at early times.  This leaves us at high redshift without   a   natural extension  of   the   bright-end  LF to  fainter magnitudes.  Presumably, future QSO surveys will detect a break in the LF  at $z \\gsim 3$  and measure a faint-end slope.   Until then, it is useful to examine other constraints.  A  popular technique  for    constraining a population of   unresolved sources is to  set an upper limit  based on their contribution  to the diffuse background   light.    For instance,   AGN are   strong  X-ray emitters, so measurements of the cosmic X-ray background might be used to provide upper bounds on the density of AGN.   The problem with this is that   the  X-ray  background is   measured  locally,  and   we are interested in constraining a high-$z$ population that contributes only a small fraction to the $z=0$ signal  (see e.g.  Hasinger 2002).  What is  preferable is a  measurement of some background at  high $z$ by an indirect method.   It turns out  that the  UV ionizing  background  is ideal for  this purpose.  As discussed  in \\S 3, the hydrogen ionizing emissivity  can be measured at   high-redshift by studying Lyman alpha absorbers    along  the  line of sight    to  distant  quasars.   This measurement is especially  useful because the  derived background at a specific  redshift   $z$  is  roughly local:    there   is very little contribution  from higher redshift sources  because the mean free path to photo-electric    absorption  is  short compared   to  cosmological distances (see Madau, Haardt, \\& Rees 1999 [MHR] and our Appendix).   \\begin{figure}[t] \\centerline{\\epsfxsize=8.5cm \\epsfbox{phiL_Lz4.ps}} \\caption{ Schematic representation of the limits on $\\phi_*$ and $L_*$ for $z=4$.  The  integrated   LF ($\\Phi(<M_B)$)  basically  constrains $\\phi_*L_*^{\\gamma_b-1}$.   We assume   that  all  reasonable  choices (shaded region) for the  bright-end slope are  bracketed by the values from  SSG  (2.87) and  SDSS  (2.58).   Since the  observations at high redshift   show no sign   of a break   for $M_B\\lesssim-26$, we do not consider  this area of   parameter space (vertical  line).  The filled square is    the point in parameter space    consistent with  both the integrated LF and the  parameters from 2dF (i.e.   $\\gamma_b=3.41$ and $\\phi_*=1070\\rm{Gpc}^{-3}$).   We also  bring   in  upper limits  from ionizing background measurements: proximity effect (``PRX.'') and flux decrement (``F.D.'').  }\\label{fig:phiL_Lz4} \\end{figure}    In what follows we use this idea to construct AGN luminosity functions that reproduce the ionizing  background measurements, allowing in some cases for significant contributions from other  sources (e.g.  stars). We can    demonstrate  our   general   program using     some   simple approximations.   Let  us assume (as  we   do throughout) that  we can approximate each AGN  as emitting light  with  a self-similar spectrum that varies only in normalization from object to object (\\S2.2).  This implies a fixed ratio between an AGN's specific luminosity at the $912 \\mbox{\\AA}$  Lyman edge to  its  luminosity in the B-band: $\\eta_{\\nu} \\equiv L_{\\nu}/L  \\simeq 10^{18}$ erg/s/Hz/L$_{\\odot}$.   The comoving emissivity at the Lyman  edge is then: $\\varepsilon_{\\nu} = \\eta_{\\nu} \\ell$, where $\\ell$ is the luminosity density of  AGN (in the B-band). For a   reasonable LF, $\\ell$  will  be dominated  by objects near the break, so we expect $\\ell \\sim \\phi_* L_*$, or $\\varepsilon_{\\nu} \\sim \\eta_{\\nu}\\phi_*  L_*$ .  A slightly more  accurate estimate of $\\ell$ comes from integrating  the LF (1) from   $L=0$ to $\\infty$, and  with this   we  obtain $\\varepsilon_{\\nu}   \\simeq  \\eta_{\\nu}   \\phi_* L_* [(2-\\gamma_f)^{-1} + (\\gamma_b-2)^{-1}]$.   Thus, a measurement of the ionizing emissivity at  some redshift mostly  constrains the parameter combination $\\phi_* L_*$, with a weak dependence on  the LF slopes (as long as neither slope  approaches  $2$).  Direct observations, on  the other hand, measure the   bright-end   slope $\\gamma_b$ and  fix   the normalization of the bright-end tail.  Together, then, the limits from $\\varepsilon_\\nu$ and direct observations can effectively restrict the acceptable values of $\\phi_*$ and $L_*$.  A schematic illustration of  this is shown in Figure  1, where we plot the available parameter space   at $z=4$.  Direct observations  of the density of bright  QSOs leads to a constraint  of $\\phi_* L_*  \\propto L_*^{2- \\gamma_b}$ (see \\S  \\ref{sect:AGNLF}), which we draw (diagonal lines) for  two values of $\\gamma_b$ from  separate surveys (2.87 from SSG; 2.58 from  Fan et al.~2001a).  We presume  that the best estimate of the bright-end  slope   lies between these two   measurements,  and therefore, that  the  values  of  $\\phi_*$   and  $L_*$ are   situated somewhere in the shaded region.   The fact that  the break has not yet been  observed down to   the  limiting magnitude $L_{\\rm lim}   \\simeq 10^{12.6} L_{\\odot}$ of SDSS allows us to set an  upper limit on $L_*$ (vertical line).  We can further narrow down  the parameter space with two different (conflicting) measurements of the ionizing emissivity at the same redshift (the higher line comes from the proximity effect and the  lower  line   comes  from  the flux   decrement distribution,  as discussed in \\S  \\ref{sect:UVBM}).   We  draw  these as  upper  limits (horizontal bands), since contributions to the background from non-AGN sources will require fewer AGN,  further restricting the allowed range of $\\phi_*L_*$.  The width of the bands reflects the slight dependence of the emissivity on $\\gamma_b$.  As  for the faint-end slope, we have fixed it to  the  low redshift  value  of 1.58, but we   explore other values of  $\\gamma_f$ in \\S  \\ref{sect:vary}.  What is evident  in the figure  is  the  fact   that the emissivity    measurements provide an extremely useful limit on the  parameter space left available from the direct observations of bright QSOs.  In  the next   section  we review what   is  known about  the  optical luminosity   function  from  direct   observations, and   discuss some possible caveats  associated with dust  and gravitational lensing.  We also discuss our  assumed template AGN spectrum.  In \\S\\ref{sect:UVBM} we  review  the different  ionizing   background measurements, and  in \\S\\ref{sect:ionsource},  we  describe   our  calculation of  the   AGN contribution to the background as well  as the contribution from stars and  IGM reemission.  We present our  results in \\S \\ref{sect:results} for   three separate  models of   AGN emissivity, and  we compare  the derived    LFs   to past   and   future     faint  surveys.   In    \\S \\ref{sect:Implics} we  discuss  our results in  the  context of  other constraints on the relative AGN and  stellar emissivities at high-$z$. We also examine what  our results imply for  the AGN  efficiencies and lifetimes.  Our conclusions are summarized in \\S 7  Unless otherwise stated, we assume that the cosmology  is one with the universe   made  flat  by a   cosmological    constant $\\Omega_m =  1- \\Omega_{\\Lambda} = 0.4$,  and that  the  Hubble parameter  at $z=0$ is $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sect:SandC}  In this paper, we have explored the implied evolution of the faint AGN luminosity  function using three   different  models for  the ionizing background from $3 \\lsim z \\lsim 6$.  Our  LFs are derived by matching the implied faint-end  emissivity.  Unfortunately,  the value of   the ionizing background rate is not universally agreed upon.  Measurements obtained using the proximity  effect generally give larger values than those based on the flux decrement distribution.  Although the current data on faint AGN  come from small-field searches with  only modest statistical  accuracy, we have   used them to obtain rough evaluations  of our Models.  If AGN  were producing a background at  the level  measured by   the   proximity effect technique,    then significantly more would  have been seen  in  the HDF and LBG  fields. For typical proximity effect values  of the background, even an escape fraction of unity would require  more than three  times the number  of AGN  observed in  the Keck  fields.   We were   able to obtain  modest agreement only  by taking the   lowest bounds on  the proximity effect measurements and by including a high galactic escape fraction: $f_{\\rm esc} \\gsim 20 \\%$.  Conversely, if the flux decrement rate is adopted, there is little room for a significant  stellar component, and $f_{\\rm esc} \\lsim 5\\%$ is required to match the faint AGN counts.  Future AGN searches and   developments in our  ability  to measure the background intensity  will be   useful   for  further constraining the    stellar contribution to the ionizing background at high-$z$.  We used our derived luminosity  functions and local determinations  of the black hole  relic density to  determine the typical AGN efficiency of converting mass  to light.  Although  many previous estimates based on X-ray counts have obtained $\\epsilon\\gsim  0.1$ (e.g.  Elvis et al. 2002),   we find significantly  lower  values  $\\epsilon\\simeq 0.028 - 0.056$, suggesting that more than half of all  AGN are obscured in the UV/optical.  Alternatively,  lower than expected values  of $\\epsilon$ may  be derived if  much of a  black hole's mass is   set by a massive ``seed'' before  the onset of  its active phase.   A joint analysis of UV-background  measurements, optical  counts,  and X-ray  observations would likely resolve this issue.  A similar analysis allows us to limit the fraction of black holes that could have been active at any redshift.   For our models that are most consistent with faint  AGN counts, we get $f_{\\rm  on} \\gsim 0.5\\%$ at $z\\sim2.5$.  We can interpret  this in terms of a  lower limit  on the average AGN  lifetime: $t_{agn} \\gsim  10^7\\rm{yrs}$.  For comparison, an AGN luminosity   function    that matches the   proximity    effect background (but over-produces the  faint counts) requires $f_{\\rm  on} \\gsim 5\\%$ and $t_{agn} \\gsim  10^8\\rm{yrs}$.  Note that these numbers are firm  lower limits  because we assume  that there  are no obscured AGN.   If a fraction  of AGN were obscured in  the optical, then these limits would be increased by the inverse of the same fraction.   Finally, we conclude with some remarks about CDM-based theories of AGN formation and how they  might  be constrained.  Usually   cosmological models of this kind rely  on simple mappings  between dark halo masses and AGN luminosities, but a fundamental unknown concerns the amount of feedback that takes place (e.g. Haehnelt \\&  Rees 1993, Haiman \\& Loeb 1998).   We point  out  that one  potentially worrisome difficulty for models with very  little feedback arises  from  direct observations of the AGN LF at bright luminosities.  Specifically, the bright-end slope measured by  the  SDSS at high  redshift   is even flatter  than  that measured by the 2dF at low redshift.  This is  the opposite of what is expected  if the AGN LF maps  simply to the   halo mass accretion rate function (see,  e.g., Haiman  \\&  Loeb 1998).  Therefore  the SDSS+2dF results alone seem to suggest that  some feedback is needed for bright to intermediate luminosity AGN.  Another question is whether any  (additional) feedback might be needed in low-mass/ low-luminosity objects in  order to explain the  observed ionizing background rate at high-$z$.  As  discussed above, our Models A and B span what we  believe to be  the range allowed by the ionizing rate measurements.  Model B corresponds to the case where AGN dominate the background, and, interestingly, its  evolution is very reminiscent to that of  the halo mass function in  CDM: faint AGN are numerous  at early times,  but  their numbers  fall-off  slowly at late times  (see Figure 8).   It   would seem that   if  AGN do dominate  the  ionizing background,  then CDM-models would  require  no additional feedback on low mass  scales (relative to  what might already  be required  at the high-mass end in order to match the SDSS results.)   If instead stars  contribute  with an escape fraction  consistent with the Steidel et al.  (2002)  report (Model A),  then many fewer AGN are permitted, and CDM-based  models would need  to be adjusted further in order  to additionally suppress  the  formation of low-luminosity AGN. Both Haiman, Madau  \\& Loeb (1999)   and Kauffmann \\&  Haehnelt (2000) presented models with this  kind of luminosity-dependent feedback, but it   is  unclear if, as   presented,  these models  would   be able to reproduce both the ionizing rate (with an allowance for stars) and the faint AGN counts in the  HDF and LBG  fields.  It would be interesting to  see such a comparison.  The  model of Kauffmann \\& Haehnelt (2000) is  ideally   suited for  this   test because  it  predicts  the joint population    of      galaxies  and  AGN      in   a   self-consistent manner.\\footnote{Indications that obscuration  plays   a key role   in determining which AGN are  seen in the optical  add an extra  layer of difficulty in this already difficult problem.}  It is becoming increasingly clear that galaxy formation, AGN activity, and  the resulting  ionizing  background are  intimately connected and take part in an   important  feedback loop (Gebhardt  et   al.  2000a; Ferrarese \\&  Merritt   2000; Kauffmann  \\&  Haehnelt 2000;   Bullock, Kravtsov, \\&  Weinberg 2000; Benson  et  al.  2002; Somerville  2002). For this reason the need is high for self-consistent models that treat all of these  processes   together, as  is  the need   for  additional observational   constraints on  models    of  this kind.   The  limits presented here  provide one small step in  this direction, but if some agreement  can be reached  on measurements  of the ionizing background emissivity, the resulting constraints would prove remarkably important for piecing together the story of cosmological structure formation."
    },
    "0207/astro-ph0207036_arXiv.txt": {
        "abstract": "We present an implicit finite difference representation for general relativistic radiation hydrodynamics in spherical symmetry. Our code, {\\sc agile-boltztran}, solves the Boltzmann transport equation for the angular and spectral neutrino distribution functions in self-consistent simulations of stellar core collapse and postbounce evolution. It implements a dynamically adaptive grid in comoving coordinates. A comoving frame in the momentum phase space facilitates the evaluation and tabulation of neutrino-matter interaction cross sections, but produces a multitude of observer corrections in the transport equation. Most macroscopically interesting physical quantities are defined by expectation values of the distribution function. We optimize the finite differencing of the microscopic transport equation for a consistent evolution of important expectation values. We test our code in simulations launched from progenitor stars with 13 solar masses and 40 solar masses. Half a second after core collapse and bounce, the protoneutron star in the latter case reaches its maximum mass and collapses further to form a black hole. When the hydrostatic gravitational contraction sets in, we find a transient increase in electron flavor neutrino luminosities due to a change in the accretion rate. The \\( \\mu  \\)- and \\( \\tau  \\)-neutrino luminosities and rms energies, however, continue to rise because previously shock-heated material with a non-degenerate electron gas starts to replace the cool degenerate material at their production site. We demonstrate this by supplementing the concept of neutrinospheres with a more detailed statistical description of the origin of escaping neutrinos. Adhering to our tradition, we compare the evolution of the 13 solar mass progenitor star to corresponding simulations with the multi-group flux-limited diffusion approximation, based on a recently developed flux limiter. We find similar results in the postbounce phase and validate this {\\sc mgfld} approach for the spherically symmetric case with standard input physics. ",
        "introduction": "\\label{section_neutrino_transport}  A supernova explosion is a dramatic event which includes such a rich diversity of physics \\citep{Bethe_90,Burrows_Young_00,Mezzacappa_Bruenn_00,Janka_Kifonidis_Rampp_01} that self-consistent numerical simulations on current computer hardware do not allow to include all relevant pieces at once. After stellar core collapse, a nascent neutron star is formed at the center of the event, requiring a description in general relativity. Neutrinos generated in the compactifying region are tightly coupled to the matter at high densities, while they leak or stream out at lower densities. The collapsing stellar core bounces as the equation of state stiffens when nuclear densities are exceeded, and a shock wave is formed that ploughs outwards through the still infalling outer layers. Multi-frequency radiation hydrodynamics must be used to quantify the energy that the neutrinos deposit in the material behind the shock. This energy deposition has been considered to be essential for the success or failure of the supernova explosion \\citep{Colgate_White_66,Wilson_85,Bethe_Wilson_85}. Observational and theoretical evidence suggests that this neutrino heating drives convection behind the shock. Instabilities in the protoneutron star, significant rotation, and strong magnetic fields further complicate the picture. Observations of neutron star kicks, mixing of nuclear species, inhomogeneous ejecta, and polarization of spectra support the presence of asymmetries in supernova explosions \\citep{Tueller_et_al_91,Strom_et_al_95,Galama_et_al_98,Leonard_et_al_00}. Motivated by such observations, the neutrino driven explosion mechanism has been explored in multidimensions \\citep{Herant_Benz_Colgate_92,Miller_Wilson_Mayle_93,Herant_et_al_94,Burrows_Hayes_Fryxell_95,Janka_Mueller_96,Mezzacappa_et_al_98b,Fryer_Heger_00,Fryer_Warren_02}, and alternative jet-based explosion scenarios have received new momentum \\citep{Hoeflich_Wheeler_Wang_99,Khokhlov_et_al_99,MacFadyen_Woosley_99,Wheeler_et_al_00}.  Because of their excessive computational demand, multi-dimensional simulations have had to rely on physically significant approximations. Simulations are performed that impose the neutrino radiation field externally or uniformly, that prescribe the neutrino spectrum ab initio, that artificially seed instabilities, or that insert jets in order to explore important phenomena of supernova explosions in observational data. A complementary approach is to model supernovae in spherical symmetry. Spherically symmetric simulations of stellar core collapse reach back to the late sixties \\citep{Colgate_White_66,May_White_67,Arnett_67,Schwartz_67,Wilson_71} when computers first became available. Although the assumption of pure spherical symmetry is also a physically significant simplification, it can provide an essential contribution to the quality of numerical supernova models and their interpretation. A spherically symmetric model can implement full general relativity. Additionally, the reduced number of computational zones allows an accurate treatment of neutrino radiation transport and the inclusion of sophisticated microscopic input physics. During the search for a robust supernova mechanism, standards in the nuclear input physics \\citep{Lattimer_Swesty_91} and weak interaction physics \\citep{Bruenn_85} have been established. In recognition of the importance of neutrino transport in the supernova, its numerical treatment has been improved from simple leakage schemes \\citep{VanRiper_Lattimer_81,Baron_Cooperstein_Kahana_85b} to multi-group flux-limited diffusion ({\\sc mgfld}) approximations \\citep{Arnett_77,Bowers_Wilson_82,Bruenn_85,Myra_et_al_87} to investigations of the full Boltzmann transport equation \\citep{Mezzacappa_Bruenn_93a,Messer_et_al_98,Yamada_Janka_Suzuki_99,Burrows_et_al_00,Rampp_Janka_02}. Only recently have self-contained simulations of stellar core collapse and postbounce evolution with Boltzmann neutrino transport been performed, in Newtonian gravity and the {\\cal O}\\( (v/c) \\) limit \\citep{Rampp_Janka_00,Mezzacappa_et_al_01,Thompson_Burrows_Pinto_03} and in general relativistic space-time \\citep{Liebendoerfer_et_al_01}. They allow, on the one hand, the exploration of nuclear physics under extreme conditions, and, on the other hand, provide a solid point of reference for the construction of accurate multidimensional simulations \\citep{Buras_et_al_03b}.  Basic numerical techniques for radiative transfer have been designed and continuously improved over the last three decades (see e.g. \\citep{Mihalas_Mihalas_84} and \\citep{Lewis_Miller_84} for a review). In most applications, photons or neutrons are the transported particles. In our application to supernova dynamics, the transport of neutrinos requires a new combination of capabilities derived from both fields, photon and neutron transport.  As in \\emph{neutron transport} the radiation particles are fermions. In contrast to photons, the neutrinos assume a Fermi-Dirac distribution in thermal equilibrium. Weak interactions with nuclear matter can be suppressed due to Pauli blocking of final states. This makes the dependence of the collision integral on the distribution function nonlinear. Several quantities have to be evolved with special care: An important constraint for the evolution of neutrino and electron abundances is lepton number conservation. While applications to photon transport do not require an exact count of photons as particles, the outcome of a supernova simulation sensitively depends on the deleptonization. The transfer of lepton number occurs via the reactions \\( e^{-}+p^{+}\\rightleftharpoons n+\\nu _{e} \\) and \\( e^{+}+n\\rightleftharpoons p+\\bar{\\nu }_{e} \\) where the protons, \\( p \\), and neutrons, \\( n \\), are free or bound in nuclei. The electron neutrino, \\( \\nu _{e} \\), and antineutrino, \\( \\bar{\\nu }_{e} \\), may escape to infinity or be absorbed at distant locations. These reactions determine the electron fraction, and hence, the partial pressure of the electron gas in the fluid. This pressure contributes importantly to the dynamics if the electron gas is degenerate and the density lower than nuclear density. Beside of the transport of lepton number, the transport of energy is a crucial phenomenon in supernova dynamics. Neutrinos escaping from the accreting material and the protoneutron star may be absorbed behind the stalled shock. This energy transfer is believed to add to the thermal pressure behind the shock such that ultimately a supernova explosion is launched with the ejection of the outer layers \\citep{Colgate_White_66,Bethe_Wilson_85}. Last, but not least, the evolution of the global energy should be monitored. Just before collapse, the progenitor star is marginally bound. During collapse, the binding energy of the nascent protoneutron star increases dramatically. It is balanced by the internal energy of the compressed matter and the trapped neutrinos. Only on a longer time scale is the energy of the radiation field transported away from the star while the emptied states in the neutrino phase space are immediately replenished by neutrino emission in the cooling protoneutron star. Since the gravitational energy, the internal energy, and the energy in the radiation field rise to the order of \\( 10^{53} \\) erg, a detailed analysis of global energy conservation is advised in order to make the comparatively small explosion energy of \\( 10^{51} \\)erg predictable. Whether we choose a kinetic equation for the propagation of particle number or an equation prescribing the evolution of radiation intensity, it is a numerical challenge to conserve both lepton number and energy with the same finite difference representation of the transport equation.  The following transport requirements more closely resemble the astrophysical applications of \\emph{photon transport} than neutron transport. Owing to the protoneutron star or black hole formed at the center of the event, space-time is curved. The particles follow geodesics in general relativistic space-time, i.e. their angular distribution is affected by gravitational bending. Additionally, the particle energies are subject to a gravitational frequency shift. In contrast to many neutron transport applications, the fluid is highly dynamic during collapse and even more so after shock formation. Thus, observer corrections for Doppler shift and angular aberration have to be implemented as in photon transport applications. However, there is some freedom in the choice of where to apply these corrections. In an inertial frame they complicate the description of interactions in the collision integral. In a comoving frame they enter the transport terms in the Boltzmann equation.  We proceed with the latter choice in an \\( S_{N} \\)-method. Our code, {\\sc agile-boltztran}, emerged from the following components: The neutrino transport part, {\\sc boltztran}, has been developed for the simulation of stellar core collapse in an implicitly finite differenced {\\cal O}\\( (v/c) \\) approximation \\citep{Mezzacappa_Bruenn_93a,Mezzacappa_Bruenn_93b,Mezzacappa_Bruenn_93c}. It has been compared to {\\sc mgfld} in selected stationary state phases, and standard test problems for radiative transfer in supernovae have been performed \\citep{Messer_et_al_98,Messer_00}. {\\sc agile} is an implicit general relativistic hydrodynamics code that evolves the Einstein equations based on conservative finite differencing on an adaptive grid \\citep{Liebendoerfer_Rosswog_Thielemann_02}. In this paper, we describe how these codes are merged and extended to enable accurate simulations of the very dynamic postbounce phase. In particular, we detail the finite differencing of the observer corrections for the dynamic conservation of particle number and total energy in the transport scheme, and describe the extension of {\\sc boltztran} to general relativistic flows \\citep{Liebendoerfer_00}. In section \\ref{section_physical_model}, we start with the characterization of the equation of state and a list of the included neutrino-matter interactions, followed by a collection of the basic equations of general relativistic radiation hydrodynamics in spherical symmetry. Two exemplary runs from collapse through bounce and postbounce evolution are also described in this section. In order to test our code in a broad range of conditions, we launched simulations from a small \\( 13 \\) M\\( _{\\odot } \\) progenitor star \\citep{Nomoto_Hashimoto_88} and a very massive \\( 40 \\) M\\( _{\\odot } \\) progenitor star \\citep{Woosley_Weaver_95}. We investigate the different regions of neutrino emission in the star. Section \\ref{section_computer_model} is entirely devoted to the documentation of the finite differencing in our computer model. In section \\ref{section_code_verification}, we analyze and verify the performance of our code in various example situations encountered in the evolution of the two simulation runs. We close the section with the comparison of our results with Bruenn's multi-group flux-limited diffusion code which implements a recently developed new flux limiter \\citep{Bruenn_02}. ",
        "conclusions": "{\\sc agile-boltztran} directly solves the general relativistic Boltzmann transport equation for the specific neutrino distribution function in spherically symmetric space-time. In combination with general relativistic hydrodynamics, we present a consistent finite differencing for fermion radiation hydrodynamics in spherical symmetry. We test our code in the context of stellar core collapse and postbounce evolution of a smaller \\( 13 \\) M\\( _{\\odot } \\), and a more massive \\( 40 \\) M\\( _{\\odot } \\) progenitor star. They embrace the progenitor mass range expected to produce explosions after core collapse. But no explosions are obtained in our simulations. We analyze the regions of neutrino emission by a statistical description allowing the detailed presentation of the specific production reactions and energies of escaping neutrinos.  The Boltzmann transport equation plays a fundamental role in kinetic gas theory. One may derive from it the diffusion equation, or---if one drives the limit further to vanishing transport---the hydrodynamics equations. Thermodynamical quantities are given by the expectation values of various operators on the fundamental particle distribution functions. It is quite unique to numerically solve a single equation for the fundamental particle distribution functions in an object of astronomical size! However, this basic approach becomes intricate in the detailed finite difference representation. In theory, the solution of the Boltzmann equation guarantees the correct evolution of expectation values in the energy conservation equation or the diffusion limit. Also the equation of state for radiation, \\( p\\propto \\rho /3, \\) will emerge correctly. This is not automatically the case in a numerically solved Boltzmann equation. Especially in the astrophysical application, a distant observer will rather observe macroscopic phenomena than the microscopic local state of the matter and the radiation field. Therefore we compose the finite difference representation of the microscopic physics such that important macroscopic quantities are accurately met. This optimizes the accuracy of the simulations at moderate resolution settings. However, it would be more straightforward to comply with this ambition if there would not be an additional complication. Physically, the Boltzmann transport equation is a trivial equation, describing free particle propagation on geodesics between independent collisions (even more so in spherical symmetry, where only two independent spatial degrees of freedom have to be considered). With respect to the collisions, the supernova has been investigated as a site of interesting nuclear physics and neutrino-matter interactions. The cross sections depend on the local neutrino energies and scattering angles. They are most conveniently evaluated in the fluid rest frame and most easily described in combination with comoving coordinates. The latter are also favored by the initial core collapse where the relevant computational domain collapses with the matter from several thousands to some hundreds of kilometer radius, before it hopefully expands again with the explosion. The comoving coordinates convert the simple statement of the Boltzmann equation, that the derivative of the invariant particle distribution function along the phase flow vanishes between collisions, into a complicated sum of partial derivatives along the comoving coordinates. The finite differencing of all comoving frame correction terms have to be adjusted in a mutually dependent way to correctly transform to conservation equations for the observer in a laboratory frame. We construct number conservation to be guaranteed. With respect to energy conservation, we match the largest {\\cal O}\\( \\left( v/c\\right)  \\) terms to machine precision, while higher order gravitational terms are matched approximately, guided by specific limiting cases. The adaptive grid applied to the neutrino distribution function in the comoving frame introduces additional deviations from perfect energy conservation in the laboratory frame. In the worst case, i.e. in the evolution of the very massive \\( 40 \\) M\\( _{\\odot } \\) progenitor star, all these deviations together reach the order of a supernova explosion energy (\\( \\sim 10^{51} \\) erg) at the time the protoneutron star collapses to a black hole. However, when the conditions are most favorable for an explosion (we should rather say least unfavorable), the numerical energy gain has not yet exceeded \\( 10^{50} \\) erg and we are confident that the remaining energy drift does not influence our physical conclusions. We have also tested the implementation of the gravitational redshift and bending, the angular advection in the free streaming regime, and the propagation of the neutrino burst. Quite generally, we find that the choice of first order upwind differencing is the limiting factor for the accuracy in our numerical solution. In space, it causes artificial pulse spreading. In the advection terms of the adaptive grid, it causes artificial diffusion. In momentum space, it causes angular diffusion in the evolution of the particle distribution function. Although these limitations can be reduced by the choice of higher resolution, more elaborate advection schemes could generally be beneficial for the accuracy of the simulations. With an investigation of the resolution dependence of the results in each phase space dimension, however, we demonstrate that our physical results are sufficiently converged in the supernova application. The undesired effects of low order upwind differencing predominantly affect the regions far from the neutrinospheres, where the radiation field is decoupled from the matter such that an adverse influence on the dynamics of the model can be excluded. We derived the moments entering the energy conservation equation. Moreover, we algebraically derived the finite difference representation of the diffusive limit, the nonrelativistic limit, and the radial dependence of the radiation quantities in a stationary state free streaming radiation field. These algebraic gymnastics led to a better understanding of the internal mechanisms in our finite difference representation and, by comparison with the code results, enhanced the confidence that the implementation exactly corresponds to the programmer's intention. In order to perform simulations over a second or more, the radiation transport and hydrodynamics have been implemented with implicit finite differencing. This allows reasonably sized time steps during the neutrino heating phase. Finally we compared the evolution of a \\( 13 \\) M\\( _{\\odot } \\) star with an independently implemented general relativistic supernova code which applies the multi-group flux-limited diffusion approximation with a recently developed flux limiter \\citep{Bruenn_02}. We find significant agreement that is especially impressive in the neutrino heating phase. Other phases show regions with moderate deviations, but those are far from affecting any physical conclusions.  With the exponentially growing power of computer hardware, computational astrophysics may reach a comparable status in astronomical, nuclear and particle physics research as terrestrial experiments have received in the previous century. We may sort scientific progress into three phases: First, the astronomical observation or a new theoretical concept lead to a primary idea about the basic physics involved in an event. At this stage, order of magnitude estimates are made and compared to the qualitative observational data. The second step involves plausibility studies of the suggested scenario, which may include detailed numerical simulations and comparisons to observations. Often, however, approximations have to be used that are dictated by technical limitations in the newly developing field. Only in the simplest cases it can be shown that all assumptions are based on undeniably justifiable physical considerations. This exciting phase of discovery and controversy should be followed by more rigorous numerical simulations in a third step, when technical limitations fade and remaining approximations become quantifiable. This phase has to be accompanied by code documentation that allows independent researchers to analyze and eventually reproduce the numerical results, similar to the way the value of experimental data is enhanced by a detailed description of the reproduceable experiment. The recent emergence of functional Boltzmann solvers leads the radiation hydrodynamics in spherically symmetric supernova models into this third phase. The third phase, by principle, should be accompanied by solid agreement in the numerical solution found with independent implementations of the same physical ingredients. In order to perform such comparisons, we appreciate the situation that a basic set of standard nuclear and weak interaction physics has been defined and used for over \\( 15 \\) years (e.g. \\citep{Tubbs_Schramm_75,Schinder_Shapiro_82,Bruenn_85}). In a comparison with the second phase---which used the former standard scheme for neutrino transport in supernovae---multi-group flux-limited diffusion, we hope to increase the confidence in our results. On the other hand, we support the validity of the {\\sc mgfld} approximation in spherical symmetry if the flux limiter is chosen carefully. As satisfying as these close results are, they should not be mistaken as an example of convergence to agreement in the sense of the third phase. The third phase is only reached if similar agreement is found between independent groups and other documented codes that solve physically complete neutrino transport equations with general relativistic effects, e.g. the variable Eddington factor method implemented by \\citet{Rampp_Janka_02}. We hope that spherically symmetric simulations with neutrino radiation hydrodynamics have reached phase three and that accurate and detailed neutrino information continues to be useful for the exploration and improvement of the local microscopic input physics."
    },
    "0207/astro-ph0207346_arXiv.txt": {
        "abstract": "s{ The thermal conductivity of the color-flavor locked phase of dense quark matter is calculated. The dominant contribution to the conductivity comes from photons and Nambu-Goldstone bosons associated with the breaking of baryon number, both of which are trapped in the quark core. Because of their very large mean free path the conductivity is also very large. The cooling of the quark core arises mostly from the heat flux across the surface of direct contact with the nuclear matter. As the thermal conductivity of the neighboring layer is also high, the whole interior of the star should be nearly isothermal. Our results imply that the cooling time of compact stars with color-flavor locked quark cores is similar to that of ordinary neutron stars.} ",
        "introduction": "At sufficiently high baryon density the nucleons in nuclear matter should melt into quarks so that the system becomes a quark liquid. It should be weakly interacting due to asymptotic freedom \\citebk{ColPer}, however, it cannot be described as a simple Fermi liquid. This is due to the nonvanishing attractive interaction in the color antitriplet quark-quark channel, provided by one-gluon exchange, which renders a highly degenerate Fermi surface unstable with respect to Cooper pairing.  As a result the true ground state of dense quark matter is, in fact, a color superconductor \\citebk{old}.  Recent phenomenological \\citebk{W1S1} and microscopic studies \\citebk{PR1-Son,2nd-wave,3rd-wave,us2} have confirmed that quark matter at a sufficiently high density undergoes a phase transition into a color superconducting state. Phenomenological studies are expected to be appropriate to intermediate baryon densities, while microscopic approaches are strictly applicable at asymptotic densities where perturbation theory can be used. It is remarkable that both approaches concur that the superconducting order parameter (which determines the gap $\\Delta$ in the quark spectrum) lies between $10$ and $100$ MeV for baryon densities existing in the cores of compact stars.  At realistic baryon densities only the three lightest quarks can participate in the pairing dynamics.  The masses of the quarks are much smaller than the baryon chemical potential, thus, to a good approximation, all three flavors participate equally in the color condensation. The ground state is then the so-called color flavor locked (CFL) phase \\citebk{ARW}. The original gauge symmetry $SU(3)_{c}$ and the global chiral symmetry $SU(3)_{L} \\times SU(3)_{R}$ break down to a global diagonal ``locked\" $SU(3)_{c+L+R}$ subgroup. Because of the Higgs mechanism the gluons become massive and decouple from the infrared dynamics. The quarks also decouple because large gaps develop in their energy spectra. The breaking of the chiral symmetry leads to the appearance of an octet of pseudo-Nambu-Goldstone (NG) bosons ($\\pi^{0}$, $\\pi^{\\pm}$, $K^{\\pm}$, $K^{0}$, $\\bar{K}^{0}$, $\\eta$). In addition an extra NG boson $\\phi$ and a pseudo-NG boson $\\eta^{\\prime}$ appear in the low energy spectrum as a result of the breaking of global baryon number symmetry and approximate $U(1)_{A}$ symmetry, respectively.  The low energy action for the NG bosons in the limit of asymptotically large densities was derived in Refs.~\\citebk{CasGat,SonSt}. By making use of an auxiliary ``gauge\" symmetry, it was suggested in Ref.~\\citebk{BS} that the low energy action of Refs.~\\citebk{CasGat,SonSt} should be modified by adding a time-like covariant derivative to the action of the composite field. Under a favorable choice of parameters, the modified action predicted kaon condensation in the CFL phase. Some unusual properties of such a condensate were discussed in Ref.~\\citebk{MirSho}.  While the general structure of the low energy action in the CFL phase can be established by symmetry arguments alone \\citebk{CasGat}, the values of the parameters in such an action can be rigorously derived only at asymptotically large baryon densities \\citebk{SonSt,BS}. Thus, in the most interesting case of intermediate densities existing in the cores of compact stars, the details of the action are not well known. For the purposes of the present paper, however, it suffices to know that there are 9 massive pseudo-NG bosons and one massless NG boson $\\phi$ in the low energy spectrum. If kaons condense \\citebk{BS} an additional NG boson should appear. These NG bosons should be relevant for the kinetic properties of dense quark matter.  It has been found \\citebk{JPS} that neutrino and photon emission rates for the CFL phase are very small so that they would be inefficient in cooling the core of a neutron star. The purpose of the present investigation is to determine quantitatively the thermal conductivity of the CFL phase of dense quark matter in order to see whether it can significantly impact the cooling rate.  We shall argue that the temperature of the CFL core, as well as the neighboring neutron layer which is in contact with the core, falls quickly due to the very high thermal conductivities on both sides of the interface. In fact, to a good approximation, the interior of the star is isothermal. A noticeable gradient of the temperature appears only in a relatively thin surface layer of the star where a finite flux of energy is carried outwards by photon diffusion. A more complete account of this work is to be found in Ref.~\\citebk{SE-prc}. ",
        "conclusions": "\\label{conclusion}  Our analysis shows that the thermal conductivity of CFL color superconducting dense quark matter is very high for typical values of the temperature found in a newborn compact star. This is a direct consequence of the existence of the photon and the massless NG boson $\\phi$ whose mean free paths are very large. Note that the photons are trapped in the core because of reflection by the electron plasma in the neighboring nuclear matter. The NG bosons are also confined to the core since they can only exist in the CFL phase.  It is appropriate to mention that the (pseudo-) NG bosons and photons should also dominate other kinetic properties of dense quark matter in the CFL phase. For example, the shear viscosity should be mostly due to photons and the same massless NG bosons associated with the breaking of baryon number. The electrical conductivity, on the other hand, would be mostly due to the lightest {\\em charged} pseudo-NG boson, i.e., the $K^{+}$. Thus, in the limit of small temperatures, $T\\to 0$, the electrical conductivity will be suppressed by a factor $\\exp(-m_{K^{+}}/T)$.  Since the neutrino emissivity of the CFL core is strongly suppressed, the heat is transferred to the outer nuclear layer only through direct surface contact. While both the core and the outer layer contribute to the heat capacity of the star, it is only the outer layer which is capable of emitting this heat energy efficiently in the form of neutrinos. The combination of these two factors tends to extend the cooling time of a star. However, because of the very small thermal energy of the quark core, the time scale for cooling could be noticeably longer than that for an ordinary neutron star only if the outer nuclear layer was very thin. (Note that, while little is known about the properties of thin boundary layers outside CFL matter, it is possible that photon emission from this layer might quickly drain the relatively small amount of CFL thermal energy.) Thus it appears that the cooling of stars with not too large CFL quark cores will differ little from the cooling of typical neutron stars. A similar conclusion has been reached for stars with regular, non-CFL quark interiors \\citebk{ppls}.  In passing it is interesting to speculate about the possibility that a bare CFL quark star made entirely of dense quark matter could exist. If it were possible, it would look like a transparent dielectric \\citebk{RW-diamond}. Our present study suggests such a star would also have very unusual thermal properties. Indeed, if the star has a finite temperature $T\\lesssim \\tilde{T}$ after it was created, almost all of its thermal energy would be stored in the NG bosons. Notice that all the photons would leave the star very soon after transparency set in because the star is assumed to have no nuclear matter layer. The local interaction as well as the self-interaction of the NG bosons is very weak so that we argued in Sec.~\\ref{mfp} that their mean free path would be limited only by the geometrical size of the star. This suggests that, since photon and neutrino emission inside the CFL phase is strongly suppressed \\citebk{JPS}, the only potential source of energy loss in the bare CFL star would be the interaction of the NG bosons at the stellar boundary and photon emission there. It is likely, therefore, that such stars might be very dim and might even be good candidates for some of the baryonic dark matter in the Universe."
    },
    "0207/astro-ph0207170_arXiv.txt": {
        "abstract": "{We present radiation-driven wind models for Luminous Blue Variables (LBVs) and predict their mass-loss rates. We study the effects of lower masses and modified abundances in comparison to the normal OB supergiants, and we find that the main difference in mass loss is due to the lower masses of LBVs. In addition, we find that the increase in helium abundance changes the mass-loss properties by small amounts (up to about 0.2 dex in log $\\dot{M}$), while CNO processing is relatively unimportant for the mass-loss rate. A comparison between our mass loss predictions and the observations is performed for four relatively well-studied LBVs. The comparison shows that (i) the winds of LBVs are driven by radiation pressure on spectral lines, (ii) the variable mass loss behaviour of LBVs during their S\\,Doradus-type variation cycles is explained by changes in the line driving efficiency, notably due to the recombination/ionisation of Fe\\,{\\sc iv}/{\\sc iii} and Fe\\,{\\sc iii}/{\\sc ii}, and finally, (iii) the winds of LBVs can be used to derive their masses, as exemplified by the case of AG\\,Car, for which we derive a present-day mass of 35 \\msun. ",
        "introduction": "\\label{s_intro}  Luminous Blue Variables (LBVs) are massive stars in a brief ($\\sim$\\,10$^{5}$ yr) but violent post-main-sequence phase of evolution, during which they lose the remaining parts of their hydrogen-rich outer layers, before entering the Wolf-Rayet stage. They show a complex pattern of visual brightness variations (Humphreys \\& Davidson 1994; Lamers 1995; van Genderen 2001), but most characteristic are their visual light variations of $\\Delta V$\\,$\\simeq$\\,1 -- 2 magnitudes on timescales of years. These changes are referred to as ``typical'' or ``S\\,Doradus-type'' outbursts. During these variability cycles, the stars move horizontally in the Hertzsprung-Russell Diagram (HRD) as they expand in radius at approximately constant luminosity. At minimum visual brightness the star is relatively hot ($\\teff$\\,$\\sim$\\,20\\,000 -- 30\\,000 K) and small, whereas the star is large and relatively cool ($\\teff$\\,$\\sim$\\,8\\,000 K) at visual maximum. Although it has been established that the typical variations are related to changes in stellar radius rather than the formation of a pseudo-photosphere (Leitherer et al. 1989; de Koter et al. 1996), the physical reason for these radius changes -- which can be as large as factors of ten -- is as yet unidentified (see Nota \\& Lamers 1997 for overviews).  The strong stellar winds of LBV stars ($\\mdot$\\,$\\sim$\\,$10^{-4}$ -- 10$^{-6}$ \\msunyr) show a wide variety of mass loss behaviour, which is poorly understood. There are cases like R\\,71 where the mass loss increases while the star expands (Leitherer et al. 1989), whereas R\\,110 shows a behaviour that is the exact opposite: while the star increases in size, its mass-loss rate drops (Stahl et al. 1990). These results imply that there is no straightforward correlation between mass loss and radius (or effective temperature) for LBVs (Leitherer 1997).  Recent radiation-driven wind models for normal O and B stars however show that one does not expect such a strict monotonous behaviour of mass loss with effective temperature. Vink et al. (1999, 2000) have found an overall trend that cooler stars -- at constant luminosity -- experience smaller mass loss, but at specific effective temperatures the mass-loss rate jumps upward. The overall trend can be explained by a growing mismatch between the position of the bulk of the driving lines (mostly in the ultraviolet) and the location of the flux maximum, which shifts gradually towards the optical for cooler stars. Superimposed on this there are jumps, where \\mdot\\ increases steeply due to recombinations of dominant line-driving ions, specifically those of iron. At $\\teff$\\,$\\simeq$\\,25\\,000 K, the mass loss is predicted to increase by about a factor of five due to the recombination of Fe\\,{\\sc iv} to Fe\\,{\\sc iii}. Observational evidence for this so-called `bi-stability' jump has been presented by Lamers et al. (1995), who discovered an abrupt discontinuity in the carbon ionisation and the terminal wind velocities of stars at spectral type B1 ($\\teff$\\,$\\simeq$\\,21\\,000 K). They also present evidence for a second jump at spectral type A0 ($\\teff$\\,$\\simeq$\\,10\\,000 K), which is anticipated to be due to the recombination of Fe\\,{\\sc iii} to Fe\\,{\\sc ii} (Leitherer et al. 1989).  The main goal of this paper is to investigate whether radiation pressure on spectral lines is responsible for the mass-loss rates and mass-loss variability of LBV winds, or that alternative mechanisms, such as pulsations or magnetic fields, need to be invoked. If radiation-pressure is the main driver, just as it is in OB-type stars, then one expects the effective temperatures at the epochs at which $\\mdot$ has been deduced, to determine whether an increase or decrease in mass loss occurs. In such a picture, the seemingly contradictory mass loss behaviour of e.g. R\\,71 and R\\,110 may be explained in terms of their differing effective temperature intervals with respect to the bi-stability jumps (de Koter 1997; Lamers 1997).  The mass-loss rates of LBVs are expected to differ from those of normal OB stars -- at the same position in the HRD -- as the result of two effects. First, their stellar masses (and consequently their mass over luminosity ratios) are expected to be lower as LBVs are more evolved. Second, their surface abundances are expected to be different as products of nucleosynthesis may have reached the stellar surface. Observations show that LBVs are indeed enriched in helium and nitrogen, and depleted in carbon and oxygen (e.g. Smith et al. 1994; Najarro et al. 1997). The LBV mass-loss predictions presented in this paper take these effects into account.  As LBVs have the property that their effective temperatures change substantially on timescales of years, they offer the unique opportunity to study mass loss as a function of effective temperature for individual stars. In this respect it is interesting to note that Stahl et al. (2001) have recently observed and analysed the Galactic LBV AG\\,Carinae and found an increase in the mass-loss rate by a factor of five during its excursion from minimum to maximum visual light. Intriguingly, this increment in mass loss is similar in magnitude to the predictions of OB star mass loss around the bi-stability jump (Vink et al. 1999).  The paper is organised as follows. We first describe the Monte Carlo method used to predict mass-loss rates (Sect.\\,\\ref{s_method}). In Sect.\\,\\ref{s_massloss} we present the LBV mass-loss predictions. A comparison between these theoretical mass-loss rates and all available mass-loss observations is subsequently presented in Sect.\\,\\ref{s_obs}. Special emphasis is attributed to AG\\,Car (Sect.\\,\\ref{s_agcar}), which is by far the best studied member of its group. Interestingly, the set of \\mdot-determinations for this star provides such tight constraints that, in addition, the stellar mass of AG\\,Car can be determined (Sect.\\,\\ref{s_mass_agcar}). We present our conclusions in Sect.\\,\\ref{s_concl}. ",
        "conclusions": "\\label{s_concl}  We have presented radiation-driven wind models for a set of stellar masses and luminosities characteristic for LBVs and have predicted their mass-loss rates. We studied the effects of lower masses and modified abundances on the mass-loss rates of LBVs in comparison to normal OB supergiants. We have found that the main difference is due to the lower masses of LBVs. In addition, we have found that an increased helium abundance changes the mass loss properties by small amounts (of up to about 0.2 dex in log $\\dot{M}$), while CNO processing is relatively unimportant for the mass-loss rate.  The comparison of our predictions with the observed mass-loss behaviour of LBVs has shown that:  \\begin{enumerate}  \\item The dominant driving mechanism of LBV winds is radiation pressure on spectral lines. Other mechanisms are expected to play only a minor role.\\\\ \\item The observed variable mass loss behaviour of LBVs of up to $\\sim$\\,0.5 dex during their S\\,Doradus-type variation cycles can be explained in terms of changes in the efficiency of line driving, as a result of the recombination/ionisation of Fe\\,{\\sc iv}/{\\sc iii} and Fe\\,{\\sc iii}/{\\sc ii}.\\\\ \\item The \\mdot(\\teff) behaviour of AG\\,Car can be matched to within 0.1 dex when a shift in temperature of $\\Delta \\teff = -6\\,000$ K is applied to our predictions. This shift is required to correct for an inaccurate calculation of the ionisation balance of iron, as we have used a modified nebular approximation for this element. Note that the shift is consistent with constraints set by observations of supergiants which indicate that the Fe\\,{\\sc iv} to {\\sc iii} and the Fe\\,{\\sc iii} to {\\sc ii} recombinations indeed occur at spectral types B1 and A0 (Lamers et al. 1995).\\\\ \\item We have deduced a value of 35 \\msun\\ for the present-day mass of AG\\,Car. If we adopt a luminosity of 10$^{6}$ \\lsun (which implies an initial mass of $\\sim$\\,65\\,\\msun) and correct for the amount of ejected material found in the circumstellar environment of the star ($\\sim$\\,25\\,\\msun; Voors et al. 2000), we arrive at a present-day mass upper limit of $\\sim$\\,40\\,\\msun, which is consistent with our finding of 35\\,\\msun.\\\\ \\end{enumerate}"
    },
    "0207/astro-ph0207493_arXiv.txt": {
        "abstract": "Ultra-light scalar fields provide an interesting alternative to WIMPS as halo dark matter. In this paper we consider the effect of embedding a supermassive black hole within such a halo, and estimate the absorption probability and the accretion rate of dark matter onto the black hole. We show that the accretion rate would be small over the lifetime of a typical halo, and hence that supermassive central black holes can coexist with scalar field halos. ",
        "introduction": "\\label{one}  The standard assumption concerning galaxy dark matter halos is that they are comprised of some weakly-interacting massive particle (WIMP). However recently there has been increased interest in an alternative possibility, that the dark matter halo may be comprised of some ultra-light scalar field \\cite{sin,lee,schunck,peebles,bento,riotto,varun,luis1,siddh,arbey,miguel}. A large number of such scalar particles, all in their ground state, can be bound by their self-gravity; the configurations possess a core radius related to the Compton wavelength of the particles in question, and for suitable choices of parameters such halos can give a good description of observed rotation curves \\cite{siddh,arbey}, and optimistically may even provide possibilities to alleviate the `cuspy core' \\cite{peebles,riotto,arbey,miguel} and `substructure' \\cite{peebles,varun,luis1} problems of the standard WIMP hypothesis \\cite{primack,salucci}.  Development of this scenario is at a primitive stage compared to the WIMP hypothesis. While it is known that the linear theory evolution of perturbations matches the standard scenario, and that time-independent equilibrium configurations can broadly reproduce desired halo properties, the scenario has not been developed in a full cosmological setting where halo formation is tracked. Nevertheless, what is known so far is sufficiently intriguing that the scenario merits further study.  In this paper, we address one requirement of the scalar-field halo model, which is that such a halo must be able to survive the existence of a supermassive black hole at its centre, as it is widely believed that such black holes reside within many or perhaps even all galaxy halos \\cite{laura}. In the WIMP scenario, the angular momentum of the individual dark matter particles, combined with their low interaction rate, ensures that the capture cross-section for halo particles by the central black hole is sufficiently small. However, the scalar-field halo regime is markedly different; the individual particles do not possess angular momentum and indeed are expected to have a Compton wavelength upwards of one parsec so that the individual particles occupy a considerable volume of space. It is important to verify that the halos are able to survive the presence of a central black hole if the scenario is to remain feasible.  This paper is constructed as follows. In Section~\\ref{two}, we describe the basic steps to model a spherical scalar halo without luminous matter. The main intention is to provide a simple panorama of the modeling and its appealing properties, such as the smooth scalar profiles. In Section~\\ref{three}, we use two complementary views of the interaction of a scalar halo and a black hole: the classical Newtonian picture and the semiclassical approximation. The latter will give us information about the absorption probability and the accretion rate of scalar matter onto the black hole, the main result of this paper. Finally, we discuss the main results and some points deserving further investigation. ",
        "conclusions": "\\label{four}  We have analyzed the impact of a central supermassive black hole on galactic halos comprised of ultra-light scalar particles. From simple physical grounds, we should expect that the accretion rate of a scalar halo onto a black hole is small, since the boson particles cannot `fit' into the horizon due to their large Compton length. Besides, the absorption probability should be proportional to the ratio of the effective `area' of the two objects, and hence proportional to  $(mr_{\\rm s})^2$ as happens for massless scalar fields \\cite{das}.  We found that the absorption probability is decreased by an extra factor $mr_{\\rm s}$, which assures the coexistence of the bosonic halo and the central black hole. For this, we showed how to construct consistent and regular solutions on different scales. In addition, the accretion rate of scalar matter onto the black hole is so small that the matter absorbed by the black hole is much less than a solar mass in the whole lifetime of the Universe. On the other hand, this result would indicate that the current observed accretion in galaxy black holes would not be due to matter provided by a scalar halo.  We have only investigated the equilibrium state of a relaxed scalar halo and a central black hole, and it would be interesting to have a more dynamical view studying the formation (simultaneously or not) of the two objects. This would require of the evolution of the full Einstein equations, which is well beyond the scope of this paper.  A related issue is the interaction of primordial black holes with cosmic scalar fields, prior to the gravitational collapse of density perturbations, as recently outlined in Refs.~\\cite{jacobson,bean}. For a cosmic scalar field endowed with a quadratic potential, the accretion rate would also be given by formula Eq.~(\\ref{3e16}), and then the field would have the oscillatory behavior Eq.~(\\ref{3e6}) near the black hole horizon. That is, the mass of the bosonic field still prevents a strong interaction between black holes and cosmic scalar fields. Other kind of scalar potentials would lead to more interesting pictures \\cite{bean}."
    },
    "0207/astro-ph0207487_arXiv.txt": {
        "abstract": "{Using numerical simulations and the Sheth-Tormen approximation we study the mass function of dark matter halos in voids. We find that the void mass function is significantly lower and its shape is different than that of the field halos. We predict that in the standard LCDM model a void with radius $10\\hMpc$ should have 50 halos with circular velocity $v_c>50$ km/s and 600 halos with $v_c>20$ km/s.}   \\addkeyword{Large--scale structure of Universe} \\addkeyword{Cosmology: theory}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207164_arXiv.txt": {
        "abstract": "When  intermediate  mass stars reach their last stages of evolution they show pronounced  oscillations.  This phenomenon happens when these stars reach the so-called  Asymptotic Giant Branch (AGB), which is a region of the  Hertzsprung-Russell  diagram  located  at about the same  region of effective  temperatures but at larger luminosities than those of regular giant  stars.  The period of these  oscillations  depends on the mass of the star.  There is growing evidence that these  oscillations are highly correlated  with mass loss and  that, as the mass  loss  increases,  the pulsations  become  more  chaotic.  In this  paper  we  study  a  simple oscillator  which  accounts for the observed  properties of this kind of stars.  This  oscillator was first proposed and studied in  \\cite{IFH92} and we extend  their study to the region of more  massive  and  luminous stars --- the region of Super-AGB  stars.  The oscillator  consists of a periodic  nonlinear  perturbation of a linear  Hamiltonian  system.  The formalism  of  dynamical  systems  theory has been used to  explore  the associated  Poincar\\'e map for the range of parameters  typical of those stars.  We have studied and characterized the dynamical behaviour of the oscillator  as the  parameters  of the model are  varied,  leading us to explore a sequence  of local and  global  bifurcations.  Among  these, a tripling  bifurcation  is  remarkable,  which allows us to show that the Poincar\\'e map is a nontwist area  preserving  map.  Meandering  curves, hierarchical-islands  traps and sticky  orbits also show up.  We discuss the  implications  of the  stickiness  phenomenon  in the  evolution and stability of the Super-AGB stars. ",
        "introduction": "The  study  of  pulsating   stars  has  attracted  much  attention  from astronomers.  Pulsational  instabilities  are  found in many  phases  of stellar  evolution,  and also for a wide range of  stellar  masses  (see \\cite{GS96}   for   an   excellent   review).   Moreover,    pulsational instabilities provide a unique opportunity to learn about the physics of stars  and  to  derive  useful   constraints  on  the  stellar  physical mechanisms that would not be accessible otherwise.  Within the theory of stellar pulsations, there are three basic characteristics of the motions that the  associated  model may include or not, namely the  oscillations can be linear or  nonlinear,  adiabatic  or not, and radial or nonradial \\cite{C80}.  Actual pulsations of real stars must certainly involve some degree of  nonlinearity  \\cite{B93,GS95,BEA01}.  In fact, the  irregular behaviour  observed in many variable stars is, definitely, the result of those  nonlinear  effects  \\cite{PEA96}.  Moreover,  the  fact  that the observed  pulsation  amplitudes of variable stars of a given type do not show huge  variations  from star to star  suggests  the  existence  of a (nonlinear)  limit-cycle type of behaviour \\cite{P90}.  However the full set of nonlinear equations is so complicated that there are no realistic stellar  models  for  which  analytic   solutions  exist  and  thus  the investigations  of  nonlinear  pulsations  rely on  numerical  analysis. Accordingly,  most  recent  theoretical  studies of  stellar  pulsations proceed   either   through   pure   numerical    hydrodynamical    codes \\cite{HEA98,H99,WEA00,SEA01,BEA00}  or, conversely,  they adopt a set of simplifying   assumptions   in  order   to  be  able  to  deal   with  a extraordinarily complex problem.  Most of these simple models of stellar pulsation are based on a one-zone type of model which may be  visualized as a single,  relatively  thin,  spherical  mass shell on top of a rigid core.  These models have helped  considerably in clarifying  some of the complicated  physics involved in stellar  pulsations and the role played by different physical mechanisms.  AGB stars are defined as stars that develop  electron  degenerate  cores made of matter  which  has  experienced  complete  hydrogen  and  helium burning,  but not  carbon  burning.  More  luminous  stars  with  highly evolved  cores are called  Super-AGB  stars and their  core is made of a mixture of  oxygen-neon  \\cite{REA96,GBEA97}.  It is a well known result that the observational counterparts of Super-AGB stars --- the so-called Long Period Variables  (LPVs) --- are radial  pulsators.  Moreover, long term  photometry  of these stars has shown that their light  curves (the variation of the luminosity with time) usually have some  irregularities \\cite{B98,BEA99}  which lead to a high degree of  impredictability.  Due to  the  lack  of  appropriate   tools  for  analyzing  these  irregular fluctuations  of the  luminosity or the stellar  radius, the  scientific community did not pay much attention,  until  recently, to this category of  variable  stars.  The  development  of  new  nonlinear   time-series analysis  tools during the last decade has changed  this  situation.  In particular, it has been found that these tools have rich applications in a broad range of astrophysical  situations,  which include time analysis of gamma-ray  bursts  \\cite{NEA94},  of  gravitationally  lensed quasars \\cite{H92}  or  X-rays  within  galaxy  clusters  \\cite{SEA94},  and  of variable white dwarfs  \\cite{GEA91}.  These tools have also been used to analyze theoretical models of stellar pulsations.  In particular, it has been proven that numerical  hydrodynamical  models  display  cascades of period doublings  \\cite{BK87,KB88,A90}  as well as tangent  bifurcations \\cite{BGK87,A87}.  To  be  more   specific,   it  has   been   confirmed \\cite{SEA96}  that the irregular  pulsations  of W~Vir models are indeed chaotic and,  furthermore,  it has been shown that the  physical  system generating  the time series is equivalent to a system of three  ordinary differential  equations.  A similar  approach has been used also for the study of the  pulsations  of other types of stars like, for instance two RV Tau stars,  R~Scuti  \\cite{BEA95}  and AC~Her  \\cite{KEA98},  and has provided  significant results concerning the underlying  dynamics.  This result is not a trivial one since it is not obvious at all that  stellar pulsations  can be fully  described by such a simple  system of ordinary differential  equations.  This  stems  from the fact that the  classical method to physically  describe stellar pulsations is based on the use of a hydrodynamical code where the partial differential  equations of fluid dynamics are replaced by a discrete approximation consisting of $N$ mass shells.  Therefore,   a  set  of   $3N$   coupled   nonlinear   ordinary differential  equations  must be solved,  where $N$ is typically  of the order of 60.  Thus,  simple  models do not only help in  clarifying  the basic behaviour of stellar pulsations but, actually, they may be able to reproduce with a reasonable accuracy the oscillations of real stars.  Several   such   simple   models   have  been   proposed   and   studied \\cite{B66,BMS66,RR70,S72,BR82,TT88,STT89,UX93}   including  the  one  in \\cite{IFH92}.  This model was intended  for study the linear,  adiabatic and radial  pulsations  of AGB stars.  The  purpose of this  paper is to analyze in depth the linear oscillator proposed by these authors, and to extend  their study to more  massive and  luminous  stars.  The paper is organized as follows.  In section II, we summarize the basic assumptions of the model and we mathematically  describe the oscillator,  whereas in section III we  characterize  its  dynamics  and  describe a sequence of bifurcations  which occur as the parameters of the model are varied.  In section  IV  we  compare  our  results  with  those  for  the  perturbed oscillator, in order to get a better physical  insight.  In section V we discuss our results and, finally, in section VI we draw our conclusions. ",
        "conclusions": "We  have  presented  and  analyzed  the  complex  dynamics  of a  forced oscillator which is interesting not only from the mathematical  point of view, but also because it describes with a reasonable degree of accuracy the main  characteristics of some stellar  oscillations.  This model has been  previously  used to study the irregular  pulsations  of Asymptotic Giant  Branch  stars  \\cite{IFH92},  and we have  used it to  study  the pulsations   of  more  massive  and   luminous   stars,  the   so-called Super-Asymptotic Giant Branch stars.  In doing this we have extended the previous studies to a range of the parameters  specific for this stellar evolutionary  phase.  We have found, in agreement with the  observations of Long Period  Variables which are the  observational  counterparts  of Super-Asymptotic Giant Branch stars, that the oscillator shows a chaotic behavior.  It is  important  to realize  as well that this kind of stars shows a more pronounced  chaotic behavior than regular  Asymptotic Giant Branch stars of smaller mass and luminosity.  We have also characterized in depth the full sequence of bifurcations as the physical parameters of the model  are  varied.  We have  found a rich set of local  and  global bifurcations  which were not  described  in  \\cite{IFH92}.  Among these, perhaps the most important one is a tripling bifurcation, but meandering curves,  hierarchical  islands  traps and  sticky  orbits  also  appear. Correspondingly,  the resulting time series also show a rich  behaviour. In particular, we have found that although  there are light curves which show a rather regular  behaviour for certain values of the parameters of the physical system and given initial conditions, there are as well some other light curves which show clear beatings or linear  combinations  of two main  frequencies  up to  terms  of  $2f_0+7f_1$,  being  $f_0$  the fundamental  frequency  and $f_1$ that of the first  overtone,  and even more complex orbits.  For the parameters and initial conditions  leading to more  irregular  behaviour,  we noticed the  existence  of both clear chaotic  pulsations  and sudden  changes from a  limit-cycle  to chaotic pulsations, the latter being  associated with the stickiness  phenomenon characteristic  of  some  Hamiltonian  systems.  For  these  orbits  the velocity of the very outer layers clearly  exceeds the escape  velocity. Hence, for these chaotic  pulsations  mass-loss is very likely to occur, in good agreement with the  observations,  which correlate the degree of irregularity  with the mass-loss rate.  Regarding the stickiness of some orbits  which we have found for a set of  parameters,  perhaps  the most important  result is that the long-term  effects found in real stars are reproduced  by our model,  despite of its  simplicity,  even  though the driven  oscillator  studied  here does not  incorporate  the  effects of secular  changes.  Hence it is quite  likely that this kind of behaviour which  has  been  already  found  in  full  hydrodynamical   simulations \\cite{YT96},   is    intrinsically    associated   with   the   physical characteristics  of the  oscillations  of  real  stars  and  not  to the long-term thermal changes."
    },
    "0207/astro-ph0207508_arXiv.txt": {
        "abstract": "Using data obtained in 1994 June/July with the {\\it Extreme Ultraviolet Explorer\\/} deep survey photometer and in 2001 January with the {\\it Chandra X-ray Observatory\\/} Low Energy Transmission Grating Spectrograph, we investigate the extreme-ultraviolet (EUV) and soft X-ray oscillations of the dwarf nova SS~Cyg in outburst. We find quasi-periodic oscillations (QPOs) at $\\nu_0\\approx 0.012$ Hz and $\\nu_1\\approx 0.13$ Hz in the EUV flux and at $\\nu_0\\approx 0.0090$ Hz, $\\nu_1\\approx 0.11$ Hz, and possibly $\\nu_2\\approx\\nu_0 + \\nu_1\\approx 0.12$ Hz in the soft X-ray flux. These data, combined with the optical data of Woudt \\& Warner for VW~Hyi, extend the Psaltis, Belloni, \\& van der Klis $\\nu_{\\rm high}$--$\\nu_{\\rm low}$ correlation for neutron star and black hole low-mass X-ray binaries (LMXBs) nearly two orders of magnitude in frequency, with $\\nu_{\\rm low} \\approx 0.08\\, \\nu_{\\rm high}$. This correlation identifies the high-frequency quasi-coherent oscillations (so-called ``dwarf nova oscillations'') of cataclysmic variables (CVs) with the kilohertz QPOs of LMXBs, and the low-frequency QPOs of CVs with the horizontal branch oscillations (or the broad noise component identified as such) of LMXBs. Assuming that the same mechanisms produce the QPOs in white dwarf, neutron star, and black hole binaries, we find that the data exclude the relativistic precession model and the magnetospheric and sonic-point beat-frequency models (as well as {\\it any\\/} model requiring the presence or absence of a stellar surface or magnetic field); more promising are models that interpret QPOs as manifestations of disk accretion onto any low-magnetic field compact object. ",
        "introduction": "Rapid periodic oscillations are observed in the optical flux of high accretion rate (``high-$\\Mdot $'') cataclysmic variables (CVs) (nova-like variables and dwarf novae in outburst) \\citep{pat81, war95a, war95b}. These oscillations have high coherence ($Q\\approx 10^4$--$10^6$), short periods ($P\\approx 7$--40 s), low amplitudes ($A\\lax 0.5$\\%), and are sinusoidal to within the limits of measurement. They are referred to as ``dwarf nova oscillations'' (DNOs) to distinguish them from the longer period, low coherence ($Q\\approx 1$--10) quasi-periodic oscillations (QPOs) of high-$\\Mdot $ CVs, and the longer period, high coherence ($Q \\approx 10^{10}$--$10^{12}$) oscillations of DQ~Her stars. DNOs appear on the rising branch of the light curve of dwarf nova outbursts, typically persist through maximum, and disappear on the declining branch of the light curve. The period of the oscillation decreases on the rising branch and increases on the declining branch, but because the period reaches minimum about one day after maximum optical flux, dwarf novae describe loops in plots of oscillation period versus optical flux.  Rapid periodic oscillations have been detected in the flux of numerous high-$\\Mdot $ CVs, but the dwarf nova SS~Cyg in outburst is one of the best studied. Oscillations in the optical flux have been detected with periods ranging from 7.3 s to 11 s \\citep{pat78, hor80, hil81, pat81}, and oscillations in the extreme-ultraviolet (EUV) and soft X-ray flux have been detected with periods ranging from 2.8 s to 11 s in data from {\\it HEAO 1\\/}, {\\it EXOSAT\\/}, the {\\it Extreme Ultraviolet Explorer\\/} ({\\it EUVE\\/}), and {\\it ROSAT\\/} \\citep{cor80, cor84, jon92, mau96, tes97,mau01b, mau02a}. \\citet{mau96, mau02a} showed that the EUV oscillation period is a single-valued function of the EUV flux, explained the loops observed in plots of oscillation period versus optical flux as the result of the delay between the rise of the optical and EUV flux at the beginning of dwarf nova outbursts \\citep{mau01a}, and so provided strong evidence for the proposal, first mooted by \\citet{pat81}, that the oscillation period depends solely on the mass-accretion rate onto the white dwarf. \\citet{mau01b} observed ``frequency doubling'' of the EUV oscillation of SS~Cyg during the rising branch of its 1996 October outburst, demonstrated that the optical and EUV oscillation periods are equal, and that their relative phase delay is consistent with zero.  In this communication, we present an analysis of observations of SS~Cyg in outburst obtained in 1994 June/July with the {\\it EUVE\\/} deep survey (DS) photometer and in 2001 January with the {\\it Chandra X-ray Observatory\\/} Low Energy Transmission Grating Spectrograph (LETGS). In \\S 2 we present the observations and data analysis, finding that the EUV flux was oscillating at frequencies $\\nu_0\\approx 0.012$ Hz and $\\nu_1 \\approx 0.13$ Hz, while the soft X-ray flux was oscillating at frequencies $\\nu_0\\approx 0.0090$ Hz, $\\nu_1\\approx 0.11$ Hz, and possibly $\\nu_2\\approx\\nu_0+\\nu_1\\approx 0.12$ Hz. In \\S 3 we discuss these results in the context of similar results from optical observations of the dwarf nova VW~Hyi in outburst and X-ray observations of neutron star and black hole low-mass X-ray binaries (LMXBs). As first pointed out by \\citet{war02a, war02b}, white dwarf, neutron star, and black hole binaries share a common correlation in the frequencies of their QPOs, extending over nearly five orders of magnitude in frequency. The implications of this result for QPO models are discussed, we discuss how additional observations of CVs can provide unique and quantitative tests of QPO models, and close in \\S 4 with a summary. ",
        "conclusions": "\\subsection{CV DNOs and QPOs}  Summarizing the results from the previous section, we have found that during the 1994 June/July {\\it EUVE\\/} observation of SS~Cyg in outburst the EUV flux was oscillating at frequencies $\\nu_0\\approx 0.012$ Hz and $\\nu_1\\approx 0.13$ Hz, while during the 2001 January {\\it Chandra\\/} LETG observation the soft X-ray flux was oscillating at frequencies $\\nu_0\\approx 0.0090$ Hz, $\\nu_1\\approx 0.11$ Hz, and possibly $\\nu_2 \\approx\\nu_0+\\nu_1\\approx 0.12$ Hz. We note that because the frequency of the $\\nu_0$ oscillation is not constant from 1994 to 2001, it cannot be due to the spin of the white dwarf, as suggested by \\citet{mau97b}. Instead, the near-constancy of the ratio $\\nu_0/\\nu_1 \\approx 0.09$ suggests that the $\\nu_0$ oscillation is related to the dominant $\\nu_1$ oscillation.  To place these results in a broader context, we note that other high-$\\Mdot $ CVs have been observed to display multiple periodicities in their power spectra. First, \\citet{wou02} list a number of instances in the literature when pairs of DNOs have been detected in the optical flux of nova-like variables and dwarf novae in outburst. Second, \\citet{ste01} detected a pair of DNOs at $\\nu_1 =0.0336$ Hz and $\\nu_2 =0.0356$ Hz in the optical continuum and Balmer emission line flux of V2051 Oph during its 1998 July outburst. Although it was not possible to detect the difference frequency $\\nu_2 -\\nu_1 =0.002$ Hz directly in the power spectrum (D.\\ Steeghs 2002, personal communication), the amplitude of the oscillation varied considerably on a timescale equal to the inverse of the difference frequency (8 min). Third, \\citet{wou02} discuss a number of instances when multiple periodicities have been detected in the optical flux of VW~Hyi in outburst. During the decline of the 2000 February outburst, DNOs with periods $P_{\\rm DNO} =27$--37 s and QPOs with periods $P_{\\rm QPO} =400$--580 s were detected simultaneously. The ratio $P_{\\rm DNO}/ P_{\\rm QPO} =0.064$--0.071, which is similar to the ratio observed in the EUV and soft X-ray oscillations of SS~Cyg. Furthermore, in data from the 1972 November outburst of VW~Hyi, three oscillations were detected simultaneously: a pair of DNOs with periods $P_1=28.77$ s and $P_2= 31.16$ s, and a QPO with a period $P_{\\rm QPO}= 349$ s. The ratio $P_2/ P_{\\rm QPO}= 0.089$ and the difference frequency $1/P_1 - 1/P_2 =0.00267$ Hz, which is close, but not equal, to the QPO frequency $1/P_{\\rm QPO} =0.00287$ Hz.  \\subsection{Possible Connection Between CV and LMXB QPOs}  Having established these properties of the QPOs of CVs, it is interesting to investigate their possible connection to the QPOs of LMXBs. \\citet{kli00} provides a comprehensive review of LMXB QPOs, so it is sufficient to note here that among neutron star binaries, the luminous Z sources have pairs of 200--1200 Hz ``kilohertz QPOs,'' 15--60 Hz ``horizontal branch oscillations'' (HBOs), and 5--20 Hz ``normal branch oscillations,'' while the less luminous atoll sources have 500--1250 Hz kilohertz QPOs, as well as 20--60 Hz QPOs and broad noise components with properties that are similar to HBOs. \\citet[hereafter PBK99]{psa99} showed that in five Z sources a tight correlation exists between the HBO frequency $\\nu_{\\rm HBO}$ and the frequency $\\nu _l$ of the lower-frequency member of the pair of kHz QPOs (the ``lower kHz QPO''). Specifically, when $\\nu_l \\lax 550$ Hz, $\\nu_{\\rm HBO}\\approx 0.12\\, \\nu_l^{0.95\\pm 0.16}.$ Furthermore, by identifying with $\\nu_{\\rm HBO}$ and $\\nu_l$ the frequencies of various types of peaked noise components in atoll sources, other neutron star binaries, and black hole binaries, PBK99 and subsequently \\citet[hereafter BPK02]{bel02} extended this correlation over nearly three orders of magnitude in frequency. As noted by \\cite{war02a, war02b}, the optical data for VW~Hyi in outburst lie on an extrapolation of this correlation, extending it an additional two orders of magnitude in frequency. Figure~4 shows that our EUV and soft X-ray data for SS~Cyg in outburst also lie on an extrapolation of this correlation, further establishing the connection between the CV and LMXB QPOs.  This connection identifies the DNOs of CVs with the kHz QPOs of LMXBs, and the QPOs of CVs with the HBOs (or the broad noise component identified as such) of LMXBs. We note that the frequencies of the DNOs of CVs and the kHz QPOs of neutron star binaries are similar in that they are comparable to the Keplerian frequency $\\nu_{\\rm K}(r) = {1\\over{2\\pi }} (GM_\\star /r^3)^{1/2}$ at the inner edge of the accretion disk of, respectively, a white dwarf and neutron star: $\\nu_{\\rm K}\\le 0.14$ Hz for a $M_\\star =1\\, \\Msun $ white dwarf with $r\\ge R_\\star=5.5\\times 10^8$ cm, while $\\nu_{\\rm K}\\lax 1570$ Hz for a $M_\\star =1.4\\, \\Msun $ neutron star with $r\\gax 3\\, R_{\\rm S} =6\\, GM_\\star/c^2 =12.4$ km, as required by general relativity. In addition to their frequencies, the DNOs of CVs and the kHz QPOs of neutron star binaries are similar in that they have relatively high coherence and high amplitudes, their frequency scales with the inferred mass-accretion rate, and they sometimes occur in pairs.  \\subsection{Implications for QPO Models}  Given the apparent connection between the QPOs of white dwarf, neutron star, and black hole binaries, it is appropriate that we investigate the implications for the theories of QPO formation. In the beat-frequency models, QPOs occur at the Keplerian frequency $\\nu_{\\rm K}(r)$ at a special radius in the accretion disk, and the beat of this frequency with the stellar rotation frequency $\\nu_\\star $. In the magnetospheric beat-frequency model \\citep{alp85, lam85}, the HBO frequency is identified with the beat frequency between the Keplerian frequency at the magnetospheric radius $r_{\\rm m}$ and the stellar rotation frequency: $\\nu_{\\rm HBO}= \\nu_{\\rm K}(r_{\\rm m})- \\nu_\\star $. In the sonic-point beat-frequency model \\citep{mil98}, some of the disk plasma makes its way past the magnetospheric radius to the ``sonic point'' radius $r_{\\rm s}$ where the disk is effectively terminated because of either radiation drag or general relativistic corrections to Newtonian gravity. In this model, the upper kHz QPO frequency is identified with the Keplerian frequency at the sonic point radius: $\\nu_u =\\nu_{\\rm K} (r_{\\rm s})$, and the lower kHz QPO frequency is identified with (one or two times) the beat frequency between the Keplerian frequency at the sonic point radius and the stellar rotation frequency: $\\nu_l = n\\, [\\nu_{\\rm K}(r_{\\rm s}) -\\nu_\\star ]$, where $n=1$ or 2. Note that in this model, the stellar rotation frequency $\\nu_\\star= (n\\, \\nu_u-\\nu_l)/n$. In the relativistic precession model \\citep{ste98, ste99}, the QPO signals are due to the fundamental frequencies of disk plasma orbiting a rapidly rotating compact star in slightly eccentric and tilted orbits. In this model, the upper kHz QPO frequency is identified with the Keplerian frequency: $\\nu_u=\\nu_{\\rm K}$, the lower kHz QPO frequency is identified with the periastron precession frequency: $\\nu_l=\\nu_{\\rm pp}$, and the HBO frequency is identified with the nodal precession frequency: $\\nu_{\\rm HBO}= \\nu_{\\rm np}$.  How does the existence of QPOs in white dwarf, neutron star, and black hole binaries constrain the theories of QPO formation? First, as has been pointed out by other authors, the existence of QPOs in both neutron star and black hole binaries excludes both flavors of the beat-frequency models (as well as {\\it any\\/} model requiring the presence or absence of a stellar surface or magnetic field). Second, the DNOs of CVs exclude the sonic-point beat-frequency model. First, there is no reason to expect a sonic point in the inner disk of a CV: $R_\\star >3\\, R_{\\rm S}$ and radiation drag is unimportant because the luminosity is a small fraction of the Eddington rate ($L\\approx GM_\\star\\Mdot/2R_\\star \\lax 3\\times 10^{35}~{\\rm erg~s^{-1}}\\approx 0.002\\, L_{\\rm Edd}$) and because the flow is only mildly relativistic ($v_{\\rm K}\\lax [GM_\\star/R_\\star ] ^{1/2}\\lax 0.02\\, c$). Second, in VW~Hyi at least, the DNO frequency separation is not equal to one or two times the white dwarf spin frequency. Using the {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) Goddard High Resolution Spectrograph, \\citet{sio95} measured the projected rotation velocity of the white dwarf in VW~Hyi in quiescence to be $v\\, \\sin i\\approx 600~\\rm km~s^{-1}$. With a binary inclination $i\\approx 60^\\circ $ and a white dwarf mass $M_\\star = 0.63~\\Msun $ (hence $R_\\star = 8.4\\times 10^8~\\rm cm$), $\\nu_\\star =v/2\\pi R_\\star \\approx 0.013$ Hz, whereas the DNO separation frequency $\\nu_u-\\nu_l\\approx 0.00267$ Hz. Third, the QPOs and DNOs of CVs exclude the relativistic precession model. First consider the nodal precession frequency, which is the sum of the relativistic (Lense-Thirring) precession frequency $\\nu_{\\rm LT}$ due to frame dragging around a rapidly rotating compact star and the classical precession frequency $\\nu_{\\rm cl}$ due to the quadrupole term in the gravitational potential of an oblate star. Because $\\nu_{\\rm cl}$ is negative for prograde orbits, $\\nu_{\\rm np} \\le \\nu_{\\rm LT}=8\\pi ^2 \\nu_{\\rm K}^2\\nu_\\star I_\\star / M_\\star c^2$, where $I_\\star\\approx 0.1\\, M_\\star R_\\star ^2$ is the moment of inertia of the star. For a $M_\\star = 1\\, \\Msun$ white dwarf, $\\nu_{\\rm K}\\le 0.14$ Hz and $I_\\star \\approx 6\\times 10^{49}~\\rm g~cm^2$, and since stability requires $\\nu_\\star < \\nu_{\\rm K},$ $\\nu_{\\rm np}\\lax 10^{-5}$ Hz, whereas the observed CV QPO frequencies $\\nu_{\\rm QPO}\\gax 10^{-3}$ Hz. Next consider the periastron precession frequency, which is the difference between the Keplerian and epicyclic frequencies. The latter term is $\\nu_{\\rm r} \\approx \\nu_{\\rm K}\\, (1-6\\, GM_\\star /c^2 r)^{1/2}$ to sufficient accuracy for a white dwarf accretor, so $\\nu_{\\rm pp}\\approx \\nu_{\\rm K} [1-(1-6\\, GM_\\star /c^2 r)^{1/2}]\\lax 10^{-3}$ Hz, whereas the observed CV DNO frequencies $\\nu_{\\rm DNO} \\gax 10^{-2}$ Hz.  In summary, assuming that the same mechanisms produce the QPOs in white dwarf, neutron star, and black hole binaries, we find that the data exclude the relativistic precession model and the magnetospheric and sonic-point beat-frequency models (as well as {\\it any\\/} model requiring the presence or absence of a stellar surface or magnetic field). More promising are models, such as the transition layer model of Titarchuk and collaborators \\citep{tit98, tit99, tit02}, that interpret QPOs as manifestations of disk accretion onto any low-magnetic field compact object.  \\subsection{Conclusion}  The results from the previous sections suggest that there is a close relationship between the QPOs of CVs and LMXBs, with the DNOs of CVs being the analogues of the kHz QPOs of LMXBs and the QPOs of CVs being the analogues of the HBOs (or the broad noise component identified as such) of LMXBs. The proposed equivalence of CV DNOs and LMXB kHz QPOs strengthens the commonly-held belief that the frequency of the kHz QPOs is equal to the Keplerian frequency at the inner edge of the accretion disk because (1) the frequency ratio $\\nu_{\\rm DNO}/\\nu_{\\rm kHz} \\approx \\nu_{\\rm K} (R_{\\rm WD})/\\nu_{\\rm K}(R_{\\rm NS}) = (M_{\\rm WD} /M_{\\rm NS})^{1/2} (R_{\\rm NS}/R_{\\rm WD})^{3/2}\\approx 10^{-4}$ and (2) in every case the DNO frequencies are consistent with the requirement that $\\nu_{\\rm DNO}\\le \\nu_{\\rm K}(R_{\\rm WD})$ \\citep{pat81, kni98}. The proposed equivalence of CV DNOs and LMXB kHz QPOs on one hand and CV QPOs and LMXB HBOs on the other hand can be strengthened by finding more CVs with pairs of DNOs and more CVs in which DNOs and QPOs are observed simultaneously. Conversely, systems known to show only QPOs could be searched more carefully for DNOs. A interesting candidate is the dwarf nova U~Gem in outburst, whose $\\approx 0.04$~Hz QPO \\citep{cor84, mas88, lon96} implies the existence of $\\approx 0.5$~Hz (2 s) DNOs (an oscillation period shorter than the integration times typically employed to search for optical oscillations). For the Keplerian frequency to be this high, the mass of the white dwarf in U~Gem must be quite high: $M_\\star\\gax 1.3~\\Msun $, assuming the \\citet{nau72} white dwarf mass-radius relation. This requirement is consistent with the white dwarf mass derived by \\citet{fri90} ($M_\\star =1.26\\pm 0.12~\\Msun $), but is inconsistent with the values derived by \\citet{lon99} ($M_\\star =1.14\\pm 0.07~\\Msun $) and \\citet{sma01} ($M_\\star =1.07\\pm 0.08~\\Msun $).  Aside from filling in the lower-left corner of Figure 4, additional observations of CVs are warranted because they provide unique and quantitative tests of QPO models. Data can be obtained from the ground in the optical and from space in the ultraviolet, EUV, and soft X-rays; the system parameters (binary inclination, white dwarf mass, radius, and rotation velocity) can be measured; eclipse mapping in edge-on systems allows the sites of flux modulations to be located and dissected; dwarf nova outbursts provide a dramatic and systematic variation in the mass-accretion rate; diagnostic emission lines are available from the optical through soft X-rays, and general relativistic affects are minimal. Furthermore, technological improvements now allow {\\it spectroscopic\\/} observations of the flux oscillations of CVs. \\citet{ste01} used the Keck Low Resolution Imaging Spectrograph ($\\lambda = 3600$--9200~\\AA ) in continuous readout mode to study the DNOs in the optical continuum and Balmer line flux of the eclipsing dwarf nova V2051 Oph in outburst. \\citet{mar98} used the {\\it HST\\/} Faint Object Spectrograph with the G160L ($\\lambda = 1150$--2510~\\AA ) grating to study the wavelength dependence of the ultraviolet DNOs of the eclipsing dwarf nova OY~Car in outburst. \\citet{mau02b} used the {\\it HST\\/} Space Telescope Imaging Spectrograph in time-tag mode with the E140H ($\\lambda = 1495$--1690~\\AA ) echelle grating to study the ultraviolet DNOs of SS~Cyg in outburst, while \\citet{mau97a} used the {\\it EUVE\\/} SW spectrometer to study the wavelength dependence of the EUV DNOs of SS~Cyg in outburst. Such observations shed new light on the QPOs of CVs and LMXBs."
    },
    "0207/astro-ph0207022_arXiv.txt": {
        "abstract": "{ At 147\\,GHz (2\\,mm wavelength), we detected three prominent AGN (NRAO\\,150, 3C\\,279, 1633+382) with {\\bf V}ery {\\bf L}ong {\\bf B}aseline {\\bf I}nterferometry (VLBI) with an angular resolution of only $\\sim 18$ micro-arcseconds. This is a new world record in radio interferometry and astronomical imaging and opens fascinating future possibilities to directly image and study the innermost regions in Quasars and other Active Galactic Nuclei. }  \\titlerunning{VLBI at 147 GHz} \\authorrunning{Krichbaum et al.} ",
        "introduction": "Even after more than 40 years after Marten Schmidt's discovery of the cosmological redshift of the hydrogen lines in 3C\\,273, and of comprehensive astrophysical research on Active Galactic Nuclei (AGN), the enigma of the origin of their extreme luminosity (ranging from radio to Gamma-ray bands) and the creation mechanism for the highly relativistic plasma jets (often extending over many hundred kpc) is still not solved. Although the majority of the scientific community regards accretion onto supermassive Black Holes as the most plausible explanation for the `quasar phenomenon', many details of the astrophysical processes taking place in the centers of these most luminous objects in the Universe still remain unexplained. In particular is the question of how the relativistic jets are made, accelerated and confined not satisfyingly answered. In order to test existing theories, most of which try to explain energy release and jet production via coupling to the accretion process onto a supermassive Black Hole (e.g. the Blandford-Payne magnetic sling-shot mechanism), the direct imaging of the innermost regions of AGN becomes of great importance. The technique of interferometry is the only astronomical observing method, which leads to such direct images.  In {\\bf V}ery {\\bf L}ong {\\bf B}aseline {\\bf I}nterferometry (VLBI) the angular resolution can be improved, either by increasing the distance between the radio telescopes or by observing at shorter wavelengths. The first approach leads to VLBI with orbiting radio antennas in space (e.g. VSOP, Hirabayashi et al. 2000), which however at present gives only an angular resolution of 0.2--0.3\\,mas (1\\,mas = $10^{-3}$ arcsec) at 5\\,GHz. The second possibility leads to VLBI at millimeter wavelengths (mm-VLBI), which furthermore facilitates the imaging of compact structures, which are self-absorbed (opaque), and therefore not directly observable, at the longer centimeter wavelengths.  Nowadays, mm-VLBI observations are regularly performed at 86\\,GHz ($\\lambda=3.5$\\,mm), where images with angular resolutions of up to $\\sim 50 \\mu$as ($1 \\mu$as = $10^{-6}$ arcsec) are obtained (e.g. Rantakyro et al. 1998, Lobanov et al. 2000).\\\\ VLBI observations at even shorter wavelengths are technically more difficult and have not yet passed the stage of test experiments on relatively short continental baselines. In 1989 and at the so far highest VLBI frequency of 223\\,GHz, the quasar 3C\\,273 was marginally detected (with SNR $\\leq 7$) on the baselines Owens Valley to Kitt Peak (845 km, 0.65G$\\lambda$) (Padin et al. 1990). In 1994 and at 215\\,GHz, first fringes were seen between the IRAM 30\\,m antenna on Pico Veleta (Spain) and a single antenna of the IRAM interferometer on Plateau de Bure (France) (Greve et al. 1995). A second experiment on this 1147 km (0.88G$\\lambda$) long baseline in 1995, resulted in the successful (SNR $\\leq 35$) VLBI detection of 8 out of 9 observed compact flat spectrum sources, including the source in the  Galactic Center Sgr\\,A* (Krichbaum et al. 1997). This experiment lead to a number of conclusions, which are important for the future: (i) a large fraction of the known cm-VLBI sources are compact enough, so that they can be observed with VLBI at short millimeter wavelengths, (ii) the VLBI jets can be traced to sub-parsec scales, however, the curvature of the jets usually increases towards the nucleus, and (iii) the internal structure of the Galactic Center source Sgr\\,A* becomes visible through the foreground scattering IGM, and the size of Sgr\\,A* must be smaller than $\\sim  20$ Schwarzschild radii (Krichbaum et al. 1998).  Between 1995 -- 2000 several attempts with various telescopes were made to achieve VLBI detections on the longer transatlantic baselines. These experiments were performed in the 2\\,mm and 1.3\\,mm bands, but failed due to technical difficulties. The recent promising detection of 3C\\,273 and 3C\\,279 at 147\\,GHz on the 3100 km (1.5G$\\lambda$) baseline between Pico Veleta and Mets\\\"ahovi (Finland) in March/April 2001 (Greve et al. 2002), and the availability of VLBI equipment and a new 2\\,mm receiver at the Heinrich Hertz telescope (HHT) on Mt. Graham (Arizona), stimulated a transatlantic VLBI experiment at 147\\,GHz, which we will describe in the following. ",
        "conclusions": ""
    },
    "0207/astro-ph0207214_arXiv.txt": {
        "abstract": "{  We suggest a new scenario to explain the outburst light curves of black hole soft X-ray transients together with the secondary maximum and the bump seen on their decay phases. Our explanations are based on the disk instability models considering  the effect of X-ray irradiation. The scenario is consistent with the observed X-ray delays by a few days  with respect to the optical rise for both the main outburst and the later maxima. We  test our scenario by  numerically solving the disk diffusion equation. The obtained model curve fits well to the observed X-ray outburst photon flux curve of the black hole soft X-ray transient GS/GRS 1124-68, a typical representative of the four BH SXTs with very similar light curves. ",
        "introduction": "Soft X-ray transients (SXTs), a subclass of low mass X-ray binaries, contain either a neutron star (NS SXTs) or a black hole (BH SXTs). Their sporadic outbursts with observed or estimated recurrence time scales changing from months to more than 50 years show a variety of light curves (Chen et al. 1997). Among these sources the BH SXTs GRO~J0422+32, A0620-00, GS/GRS~1124-68 and GS~2000+25 show very similar X-ray  outburst light curves. Rise to maximum is fast (few days). Decay can be fitted with exponentials, or with power laws. These four sources are usually labeled as FRED (fast-rise-exponential-decay) sources. We refrain from using the term FRED because the initial decays can also be fitted with power laws. Instead, we shall use the term   BH SXTs to cover only these four sources and the term ``exponential-like'' to describe their similar secular decay behaviors. Other common features of the outburst light curves are  a typical secondary maximum seen in the decay phase about two months after the main maximum and a bump seen at the end of the decay phase, after the secondary maximum. The secondary maximum is exhibited by all four BH SXTs above. The bump is seen in the outburst light curves of A 0620-00 and  GS/GRS~1124-68. It is not very clear in the light curve of GS~2000+25 and absent in that of GRO~J0422+32 (Tanaka $\\&$ Shibazaki 1996).  Three physical effects are likely to determine the behavior of SXTs. These effects are (i) viscous evolution of the accretion disk, (ii) transitions between the two stable states of the disk through hydrogen ionization, as in disk instability models, and (iii) irradiation of the disk by X-rays from the inner disk (also from the neutron star surface in the case of NS SXTs). In this paper we present a comprehensive model for the black hole SXTs based on an interplay between all three effects.  The disk instability models (DIMs) which are successful in reproducing the general characteristics of dwarf nova (DN) light curves (Osaki 1974, Hoshi 1979, Meyer $\\&$ Meyer-Hofmeister 1981) are also suggested to be the possible mechanism for the SXT outbursts (Cannizzo et al. 1985, Huang $\\&$ Wheeler 1989, Mineshige $\\&$ Wheeler 1989). According to DIMs the equilibrium solutions on the accretion rate $\\dot{M}$ - surface density $\\Sigma$ plane form an \"S\" shaped curve for a given radial distance R from the center of the disk. The upper and lower branches of the curve are viscously stable, whereas the middle branch is unstable. The unstable branch appears due to partial ionization of hydrogen and corresponds roughly to the temperature range $\\Teff \\sim 6\\times 10^3 - 10^4$ K. According to DIMs all the disk is in the cold stable branch during the quiescent state. In the cold stable regime, surface densities increase with time at all radii due to the low accretion efficiency. The critical maximum surface density of the cold branch is eventually exceeded first at some particular radius R, and the instability propagates as a density wave towards the smaller and the larger radii. If all or most of the disk jumps to the hot stable branch the increasing accretion rate leads to an outburst. For the disk to return to the quiescent phase the surface density should decrease below the critical minimum surface density of the hot stable branch at some location in the disk. This is expected to occur first at the outermost regions of the disk, where in the hot state the surface densities are lowest, while the critical surface densities are highest. Consequently, all the disk returns back to the cold stable regime by means of a cooling front propagating from the outer disk towards the inner regions. The observed decay in luminosities subsequent to a burst is expected to be driven by the propagation of this cooling front.  The standard version of DIMs have difficulties in explaining the long recurrence times and the slow decay characteristic of the SXT outbursts with plausible standard thin disk $\\a$ parameters (Shakura $\\&$ Sunyaev 1973) estimated  from DN light curve analyses. These difficulties do not exclude the DIMs as the mechanism for the SXT outbursts, but require a modification of their standard form. Including the evaporation and the consequent disk truncation in the models it is possible to explain the long recurrence time scales (few decades or more) of BH SXTs (Meyer-Hofmeister $\\&$ Meyer 1999, Dubus et al. 2001).  X-ray irradiation of the outer disk and  possibly the evaporation also seem to be responsible for the observed characteristics of the decay curves of BH SXTs. It has long been known that the X-ray irradiation is the dominant source of the observed optical light of persistent LMXBs and in particular SXTs in their outburst states. It was suggested that the X-ray irradiation may have a stabilizing effect on the outer mass flow keeping the outer disk temperatures above the hydrogen ionization temperature (Meyer $\\&$ Meyer-Hofmeister 1984, van Paradijs $\\&$ McClintock 1994). In a previous work, we showed that pure viscous evolution of matter located originally at the outer disk can reproduce the rise, the turnover and the early decay characteristics of the BH SXTs by fitting the model light curves to the observed X-ray photon flux data of the BH SXTs GS/GRS 1124-68 and GS 2000+25 (Ertan $\\&$ Alpar 1998). This also suggests the importance of the X-ray irradiation, since the irradiation  could build up the conditions for a pure viscous decay by preventing the propagation of a cooling front for long times. By analytical treatments, King $\\&$ Ritter (1998) (KR) concluded that the slow decay is produced when the disk is fully ionized by irradiation, for inner accretion rates above a certain critical $\\Mdot_{\\m{in}}$ and the decay seemingly linear in time that is seen at the end of the decay phase is produced below this $\\Mdot_{\\m{in}}$, when a cooling front propagates inwards from the outer disk. Detailed numerical calculations show that a disk fully ionized by strong irradiation gives a much slower decay than the observed decays (Cannizzo 2000). By adjusting the evaporation rate in his model,  Cannizzo (2000) obtained the observed decay behavior for a fully ionized disk. On the other hand, following the idea of KR, but employing a moderate irradiation in models, it is also possible to reproduce the observed decays with a negligible effect of evaporation on the decay curves (Dubus et al. 2001; DHL).  The vertical structure analyses of an X-ray irradiated disk show that the direct illumination of the outer disk by the central X-rays is not possible due to the self screening of the disk (Dubus et al. 1999). On account of the fact that the observations invariably give much higher $L_{\\m{opt}}/L_{\\m{x}}$ ratio than expected from intrinsic dissipation alone, Dubus et al. (1999) suggested that the X-ray illumination is very probably still present, but indirect or the outer disk is warped. Indirect irradiation of the outer disk may take place via scattering. Even if the inner disk is vertically optically thick and geometrically thin, the surface layers where heating dominates over cooling could be thermally unstable and could evaporate to a hot corona  (Shaviv $\\&$ Wehrse 1986). Existence of such a hot corona surrounding the inner disk  will be our main assumption regarding the mechanism of indirect irradiation.  Since the optical flux is X-ray irradiation dominated, a change in the intrinsic dissipation at the outer disk can hardly be the reason for the observed enhancement in optical flux at the secondary maximum and the bump. Irradiation must be incorporated in the explanation of these features. In an alternative scenario tidal instabilities due to the 3:1 resonance (Frank et al. 1992 ) at the outer disk could modify the optical light of the BH SXTs, but only if the disk size increases considerably, intercepting more of the X-rays coming from the inner disk (Haswell et al. 2000). However, it is unlikely that the consequent changes in the mass inflow rate in the outer disk can lead  to the observed secondary maximum or the bump.  A further restriction on models comes from the comparison of the optical and X-ray observations. Detailed analyses show that whenever it is possible to make a clear comparison between the optical and the X-ray observations of the main and the minor maxima of BH SXTs, it is seen that the optical rise precedes the X-ray rise. The X-ray delays with respect to the optical light are $\\sim$ 4-6 days for the main and the secondary maximum and around two weeks for the bump (Kuulkers 1998, Ebisawa et al. 1994, Orosz et al. 1997). Considering these observational constraints, we suggest an explanation for the overall outburst phase of BH SXTs including the secondary maximum (Sect. 2) and the bump (Sect. 3). The details of the numerical model testing these explanations is  presented in Sect. 4. In Sect. 5 we summarize and discuss the results. ",
        "conclusions": "In general, soft X-ray transients (SXTs) show a variety of outburst light curves (Chen et al. 1997). In the present work we concentrate on the outburst light curves having a fast rise and a long lasting ($\\sim 250$ days) decay behavior exhibited by the BH SXTs A0620-00, GS/GRS~1124-68, GS~2000+25 and GRO~J0422+32.  We presented a new scenario to explain these outburst light curves together with the  characteristic secondary maxima and the bumps seen on their decay phases. The explanations are based on the disk instability models including the effect of X-ray irradiation, and can be summarized as follows: At the beginning of the main outburst, the accumulated cold matter is released by a thermal-viscous instability. This matter expands to both lower and larger radii. Part of the outermost disk that makes a transition back to the cold state before the efficient X-ray irradiation has started, initially remains shielded by the hot inner disk. The surface densities of the hot inner disk decrease faster than the surface densities  of the shielded  disk at $R>R_{\\m{h}}$. Finally, the shielding is removed after the pressure scale height of the disk at $R~ \\la ~R_{\\m{h}}$~ has decreased enough for the central X-rays to illuminate the cold disk at  $R > R_{\\m{h}}$. The consequently triggered instability outside $R_{\\m{h}}$ leads to enhanced mass flow resulting in the secondary maximum. The hot disk radius $R_{\\m{h}}$ first increases with the triggering, and later gradually decreases, governed by the decreasing irradiation flux through the outer disk. When the inner accretion rate $\\Mdot_{\\m{in}}$ becomes comparable to the evaporation rate, the temperature of the corona  increases due to decreasing cooling rate. Consequently, the corona which is assumed to be the source of indirect illumination heats up and expands leading to more efficient X-ray irradiation and a new thermal-viscous instability beyond the present position of $R_{\\m{h}}$. The subsequent enhancement of the mass transfer rate from the outer disk results in the observed bump.  We tested this scenario by using a one dimensional numerical model (Sect. 4). The model giving the X-ray photon flux curve given in Fig. 1 works briefly as follows:  At the beginning of the main outburst, the accumulated matter is represented by a Gaussian mass distribution at the outer disk. The center of the Gaussian is taken to be $R=R_0 =9\\times 10^{10}$ cm. Its maximum density  $\\Sigma_{0}=1.1\\times 10^3$ g cm$^{-2}$ is comparable to the maximum critical surface density of the cold state for the chosen parameters $M_1=6\\Msun,~ \\a_{\\m{c}}=0.033 ~ and~ R_0 =9\\times 10^{10}$ cm (Cannizzo et al 1988, Dubus et al 2001). We start with the hot state viscosities ($\\a=\\a_{\\m{h}}=0.1$). Until the X-ray irradiation is switched on at $t=t_0$~d, we set $\\a=\\a_{\\m{c}}$ for the grid points with the effective temperatures $\\Teff$ decreasing below 10000 K, and the grid points having higher $\\Teff$ are kept in the hot state. At $t=t_0=3.0$~d, the outermost cold disk region with $\\a=\\a_{\\m{c}}$ ($R>R_{\\m{h}}$) is taken to be shielded. After $t=t_{0}$ d, we use the critical surface densities obtained for irradiated and unsteady disks (Dubus et al. 2001). From $t=t_0$ to $t=t_1$, the disk evolves with a hot inner region ($R < R_{\\m{h}}$) and a cold shielded outer region  ($R > R_{\\m{h}}$). The hot disk radius $R_{\\m{h}}$ is free to move, e.g. inwards, if the surface densities at the grid points inside $R_{\\m{h}}$ decrease below the critical minimum surface densities of the present irradiated conditions. We find that  $R_{\\m{h}}$ remains constant until the onset of the secondary maximum ($t=t_1$) for the irradiation strength ($C=2.3\\times 10^{-4}$) that gives the best fit to the data until the onset of the bump ($t=t_2$). Our 1-D numerical model does not give the irradiated vertical disk structure, so we remove the shielding at a  parametrized time $t=t_1$ whose value, $t_1 =68$ d, is determined from the best fits. We assume that all the disk has become irradiated at $t=t_1$, and the part of the outer disk which has so far remained shielded and cold makes a transition to the hot state at $t=t_1$. After $t=t_1$, $R_{\\m{h}}$ first increases by the removal of the shielding and then decreases gradually governed by the decreasing irradiation strength. To check for an indication that the shielding is removed in the 1-D model, we plot the disk thickness profiles calculated by the midplane temperatures at different times of the main decay phase. We see that the shielding of the outer disk ($R >R_{\\m{h}}$) which is present during the main decay is removed at $t \\sim t_1$ for an inner scattering region with a size $l~ \\la ~10^9$ cm. The confirmation of the initial settling and the removal of the shielding in a self-consistent model will be the subject of future work. The removal of the shielding at $t=t_1$ exposes the outer disk to irradiation. This leads to the secondary maximum, first in the optical, and then in the X-rays, as the enhanced mass flow reaches the inner disk. With $t_1 =68$~d, the model reproduces the secondary maximum.  Increasing irradiation efficiency through an expanding corona, rather than removal of shielding, is taken to be the cause of the bump. At $t=t_2=155$ d,  the irradiation temperatures are increased by a constant ($\\sim 1.6$) factor chosen to fit the amplitude of the bump. In our model, the transitions to the hot state at $t=t_1$ and later at $t=t_2$ are taken to occur simultaneously in the newly irradiated regions, rather than by the propagation of thermal fronts. The tidal forces at outer disk and the accretion from the companion are neglected. We obtained the outer disk radius $R_{\\m{out}}=2\\times 10^{11}$ cm from our fits. This is about the  truncation radius $R_{\\m{tr}}\\simeq 0.9 R_{\\m{L}_{1}}$. The matter going beyond this radius is assumed to be lost from the system. In Sec. 4.6 we showed that the simplifications of our numerical model does not significantly affect the results.  The maximum of the outburst in X-rays is reached at $\\sim 7$ d in the model. The time interval between the triggering and the beginning of the rise given by the model is about 5 days for the secondary maximum and 10 days for the bump. These results are in good agreement  with the reported delays of the X-ray maxima with respect to the optical which are $\\sim 4-6$ days for the main outburst and the secondary maximum, and about two weeks for the bump (Kuulkers 1998, Ebisawa et al. 1994, Orosz et al. 1997).  We have thus shown, through a numerical model with plausible parameters (Fig. 1), that the rise and the decay of BH SXT outbursts including the secondary maximum and the bump are a simple interactive history of the effects of viscous diffusion, irradiation and hot-cold state transitions. Both the secondary maximum  and the bump are events taking place, only once, in the evolution of the disk after the main outburst.  The observed fluctuations shown during the bump could be due to the small variations of the inner disk radius when the evaporation rate becomes comparable to the inner disk accretion rate. The evaporation is expected to be strongest at the inner disk and steeply decrease with increasing radius (Meyer et al. 2000, Dubus et al 2001). Then the large density gradients due to an evacuation of the inner disk by the evaporation could result in the variations of the inner accretion rate, which could easily modify the X-ray flux. In our numerical model, we do not address these fluctuations. A reasonable fit to what might plausibly be the mean behavior of the bump is produced for GRS/GS 1124-68.  In our numerical model we do not include possible mass losses from the disk surface due to evaporation (Meyer et al. 2000, Dubus et al. 2001, Shaviv et al. 1999). The X-ray irradiation is expected to increase the wind losses from the outer disk (de Kool $\\&$ Wickramasinghe 1999). Although there is no general agreement on the physics of evaporation, it has probably no significant effect on the outburst light curves (Dubus et al. 2001). However, there is a  consensus that evaporation is important in the quiescent states, and the inner disk is probably truncated because of the resultant mass losses. Meyer-Hofmeister $\\&$ Meyer (1999) and Dubus et al. (2001) obtain the long outburst recurrence times of SXTs ($\\sim$ few ten years or more)  by including evaporation in their numerical model.  ~~~  {\\it \\bf Acknowledgments} We acknowledge support from the BDP Doctoral Research program of T{\\\"U}B{\\.I}TAK (The Scientific and Technical Research Council of Turkey) and through The High Energy Astrophysics Research Group TBAG-\\c{C}-4 of T{\\\"U}B{\\.I}TAK. We thank S. Kitamoto for providing GINGA ASM outburst data for the source GS/GRS 1124-68. Part of this research was done during the Summer Research Semester in Astrophysics (2001) of the Feza G{\\\"u}rsey Institute of T{\\\"U}B{\\.I}TAK. MAA acknowledges support from the Turkish Academy of Sciences. We thank the referee for useful criticism."
    },
    "0207/astro-ph0207091_arXiv.txt": {
        "abstract": "Hubble Space Telescope (HST) ultraviolet STIS imaging and spectroscopy of the low luminosity AGN (LLAGN) NGC 4303 have identified the previously detected UV-bright nucleus of this galaxy, as a compact, massive and luminous stellar cluster. The cluster with a size (FWHM) of 3.1 pc, and an ultraviolet luminosity log L$_{1500\\AA}$(erg s$^{-1}$ \\AA$^{-1}$)= 38.33 is identified as a nuclear super star cluster (SSC) like those detected in the circumnuclear regions of spirals and starburst galaxies. The UV spectrum showing the characteristic broad P Cygni lines produced by the winds of massive young stars, is best fitted by the spectral energy distribution of a massive cluster of 10$^{5}$ M$_{\\odot}$ (for a Salpeter IMF law with lower-mass cutoff of 1 M$_{\\odot}$) generated in an instantaneous burst 4 Myr ago. The ionizing energy produced by this cluster exceeds the flux needed to explain the nuclear H$\\alpha$ luminosity. No evidence for an additional non-thermal ionizing source associated with an accreting black hole is detected in the ultraviolet. These new HST/STIS results show unambiguously the presence of a compact, super star cluster in the nucleus of a low luminosity AGN, that is also its dominant ionizing source.  We hypothesize that at least some LLAGNs in spirals could be understood as the result of the combined ionizing radiation emitted by an evolving SSC (i.e. determined by the mass and age) and a black hole (BH) accreting with low radiative efficiency (i.e. radiating at low sub-Eddington luminosities), coexisting in the inner few parsecs region. Complementary multifrequency studies give the first hints of the very complex structure of the central 10 pc of NGC 4303 where a young SSC apparently coexists with a low efficiency accreting black hole and with an intermediate/old compact star cluster, and where in addition an evolved starburst could also be present. If structures as such detected in NGC 4303 are common in the nuclei of spirals, the modeling of the different stellar components, and their contribution to the dynamical mass, has to be established accurately before deriving any firm conclusion about the mass of central black holes of few to several million solar masses. ",
        "introduction": "Low luminosity Active Galactic Nuclei (LLAGNs) including low-luminosity Seyferts, classical LINERs, weak$-$[\\ion{O}{1}] LINERs, and LINER/HII transition-like objects are the most common type of galaxies showing nuclear activity. LINERs alone make up 50\\%$-$70\\% of AGNs and 20\\%$-$30\\% of all galaxies in surveys of nearby bright galaxies (Ho, Filippenko \\& Sargent 1997). It is therefore of fundamental importance to identify unambiguously the nature of the energy source in LLAGNs, and to quantify the contribution of stars and accreting black-holes to their energy output.  Recent X-ray observations of LLAGNs confirm that some LINERs are the low luminosity end of the luminous AGN phenomenon identified in Seyfert 1 and QSOs where accreting black-holes are the dominant energy source (Ho et al. 2001; Eracleous et al. 2002), while in others the energy output seems to be dominated by powerful starbursts (Terashima et al. 2000; Eracleous et al. 2002). HST ultraviolet imaging and spectroscopy has shown the presence of young massive stars at various scales in AGNs. Seyfert 2 galaxies known to have bright kpc-size star-forming rings show that the AGN core is barely detected in the UV, and that the massive stars dominate the observed circumnuclear UV emission (Colina et al. 1997a). The detection of stellar winds and photospheric absorption lines in the ultraviolet spectra of Seyfert 2 galaxies (Heckman et al. 1997; Gonz\\'alez Delgado et al. 1998) has proven unequivocally the presence of clusters of massive young stars in the circumnuclear regions of these galaxies. These nuclear starbursts are dusty star-forming regions of a few hundred parsecs in size and with an average age of 3 to 6 Myrs (Gonz\\'alez Delgado et al. 1998). Optical and UV studies of Seyfert 2 galaxies have detected the presence of young massive starbursts within 300 pc of the nucleus in 30\\% to 50\\% of the galaxies investigated (Cid Fernandes et al. 2001 and references therein). Finally, the compact ultraviolet sources detected in some UV bright LINERs are nuclear star clusters with sizes of about 10-15 pc like in NGC 4569 (Maoz et al. 1998; Barth et al. 1998).  In summary, recent evidence collected mainly with HST and {\\it Chandra} indicates the presence of massive stellar clusters in the circumnuclear regions (few to several hundred parsecs) in a large fraction of LINERs and Seyfert 2 galaxies, that substantially contribute to their energy output.  However, the fundamental question of whether or not massive stellar clusters exist in the central few pc region, i.e. nucleus, of LLAGNs contributing substantially, or even dominating, the energy output, requires detailed investigations in selected nearby galaxies like NGC 4303.  NGC 4303 (M 61) is a barred spiral classified as SAB(rs)bc (de Vaucouleurs et al. 1991) and located in the Virgo cluster (adopted distance of 16.1 Mpc hereinafter; see also Colina et al. 1997b). Multi-wavelength HST images of NGC 4303 have unveiled the presence of a nuclear stellar bar of 250 pc in size centered on a bright optical and near-IR nucleus\\footnote{throughout the paper we distinguish between nucleus, nuclear and circumnuclear regions each indicating different physical sizes of about 0.01, 0.1 and 1 kpc, respectively}, itself connected with a nuclear star-forming spiral of about 250 pc in radius. (Colina et al. 1997b; Colina \\& Arribas 1999; Colina \\& Wada 2000). The UV luminosity of the spiral dominates the observed integrated UV output. The brightest knots delineating the spiral have observed UV luminosities L(2200\\AA) $\\sim$ 2 $\\times$ 10$^{37}$ erg s$^{-1}$ \\AA$^{-1}$ similar to that of the R136 cluster in 30 Doradus (Vacca et al. 1995), and in the high UV luminosity end of the distribution function of compact stellar clusters detected in star-forming rings (Maoz et al. 1996).  The 2D velocity field of the warm ionized gas and cold molecular gas shows that the nucleus is also the dynamical center of a disk of 300 pc in radius, and in which the star-forming spiral is embedded (Colina \\& Arribas 1999; Schinnerer et al. 2002; Colina et al. 2002, unpublished). Ground-based spectroscopy reveals that the optical emission lines emitted by the nuclear region ($\\sim$ 100 pc size) have ratios in the borderline of Seyfert 2 and LINER nuclei (Colina \\& Arribas 1999). These observations confirm that the bright nucleus of NGC 4303 is a low luminosity AGN and is indeed located at the true center of the galaxy. ",
        "conclusions": "\\subsection{NGC 4303 nucleus: A LLAGN powered by  super star cluster}  The nuclear cluster characterized by its size of 3.1 pc, its mass of 10$^5$ M$\\_{\\odot}$, and its ultraviolet luminosity log L$_{1500\\AA}$(erg s$^{-1}$ \\AA$^{-1}$)= 38.33, or log L$_{1500\\AA}$(erg s$^{-1}$)= 41.51 assuming $\\nu$ $\\times$ f$_{\\nu}$ $\\times$ 4$\\pi$D$^2$ as the monochromatic luminosity, belongs to the class of luminous super star clusters (SSCs) found at the heart of spirals (Carollo et al. 1997), and in nuclear starburst galaxies (Meurer et al. 1995). At the derived age of 4 Myrs, the SSC in the nucleus of NGC 4303 is extremely luminous with a bolometric luminosity of about 10$^8$ L$_{\\sun}$, and an ionizing flux capable of producing an H$\\alpha$ luminosity of up to about 1.7 $\\times$ 10$^{39}$ erg s$^{-1}$, if all ionizing photons are absorbed by the surrounding interstellar medium.  The H$\\alpha$ flux in the 1\\farcs0 $\\times$ 1\\farcs0 nuclear region along PA220 corresponds to a luminosity of 1.2 $\\times$ 10$^{39}$ erg s$^{-1}$, after correcting for an stellar absorption of 1.8\\AA~ (equivalent width), and after slit/seeing aperture correction effects have been taken into account. The ratio of the predicted H$\\alpha$ flux emitted by the SSC to the measured nuclear H$\\alpha$ flux is 1.4.  Therefore no additional ionizing source other than the SSC itself is required to explain the ionized gas luminosity.   \\subsection{NGC 4303 nucleus: coexisting star clusters and AGN}  As mentioned above, no evidence for a second energy source is present in the ultraviolet spectrum of the nucleus. However, its absolute optical magnitude M$_{F606W}$= $-$14.2 and very red nuclear colors m$_{F606W}-$m$_{F160W}$= $+$3.5 obtained from HST 0\\farcs2 radius aperture measurements with filters WFPC2/F606W and NIC2/F160W\\footnote{Although filters F606W and F160W have a broadband profile different from the standard ground-based V and H filters, their effective wavelengths are very similar and consequently we adopt here the same names for simplicity.} (see also Colina \\& Wada 2000), can not be explained by the emission due to the unobscured UV-bright SSC.  According to the STARBURST99 models (Leitherer et al. 1999), the 4 Myrs, 10$^5$ M$_{\\sun}$ UV-bright cluster should have an absolute visual magnitude of about $-$13, i.e. three times fainter than measured, and an optical $-$ near-infrared color V$-$H $\\sim +$0.2, i.e. much bluer than measured. Optical emission lines could affect the measured F606W magnitude, and therefore the observed V$-$H color could not represent the intrinsic continuum value, that would be even redder than observed.  The WFPC2/F606W filter is an extreme broad-band filter with a bandpass of almost 1600\\AA~ that includes all the optical emission lines in the 4800\\AA~ to 7000\\AA~ spectral range. The predicted equivalent widths of the Balmer lines produced by a 4 Myr cluster are about 100\\AA~ and 400\\AA~ for H$\\beta$ and H$\\alpha$, respectively (Leitherer et al. 1999).  However, the ionized region centered on the nucleus is extended over about 1\\farcs5, as traced by the angular size of the Ly$\\alpha$ emission line region in our STIS long-slit spectrum. Moreover, the combined equivalent width of the nuclear [NII]+H$\\alpha$ line complex, as measured in our WHT spectrum of the nuclear region, corresponds to about 20\\AA~ indicating dilution by the bulge starlight contribution, that dominates the optical continuum in our ground-based spectrum. In summary, the contribution of the optical emission lines to the measured nuclear F606W flux is not expected to be more than a few percent.  The nuclear m$_{F606W}-$m$_{F160W}$ color is also redder than the average bulge's color in face-on spirals (V$-$H= 2.71 $\\pm$ 0.33; de Jong, \\& van der Kruit 1994) and therefore, the presence of an additional red and luminous source has to be invoked. This additional source could either be an intermediate/old stellar population associated with the bulge, an evolved starburst dominated by red supergiants (about 10 Myr), or an accreting black hole. In the following these alternatives are discussed.  The first possibility would be that the bright near-infrared source traces the presence of an intermediate/old nuclear star cluster. Optical HST imaging has shown that many nearby spirals harbor nuclear star clusters, with a small fraction (about 5\\%) being unresolved (Carollo et al. 1997). Our long-slit optical continuum integrated over 1\\farcs0 $\\times$ 1\\farcs0 seems indeed to be dominated by an evolved stellar population around 1-5 Gyr old, with a total initial mass around 10$^8$ M$_\\odot$ (Salpeter IMF normalized between 1 and 100 M$_\\odot$). Even allowing for some reddening (E(B-V) of 0.4 and 0.1 for both ages, respectively), the V$-$H color of this population wouldn't be redder than around 2.5, not being able to explain completely the observed infrared luminosity.  The excess H-band luminosity could also originate, at least partially, from a population of red supergiant stars formed in a previous starburst episode around 10 Myr ago (Cervi\\~no \\& Mas-Hesse 1994; Leitherer et al. 1999), so that a two-stage starburst consisting of 4 + 10 Myr old stars within a few pc could coexist in the nucleus of NGC 4303.  This scenario of a two-stage starburst is reminiscent of the star formation observed in the central 0.5 pc of the Milky Way and in NGC 1569, where a recent 4-8 Myr starburst, fully accounting for the ionizing and bolometric luminosities, coexist with an older cluster (Krabbe et al. 1995; Gonzalez Delgado et al. 1997).  If a 10 Myr cluster were the dominating source of the luminous H-band point source (M$_{F160W}$= $-$17.7) in NGC 4303, it would have had a mass of about 2 $\\times$ 10$^6$ M$_{\\odot}$.  If we include the H-band luminosity associated with the very intermediate/old ($>$ 1 Gyr) stellar population, a starburst of about 5 $\\times$ 10$^5$ M$_{\\odot}$ would be required. In either case, such a massive, unobscured cluster would produce an ultraviolet flux in excess of what is observed, and would produce a much bluer optical continuum than observed.  Therefore, if the 10 Myr old starburst were present, it should be very significantly obscured both in the optical and ultraviolet ranges, with its associated light emerging only in the infrared, if at all.  Finally, the nuclear, luminous near-infrared source could alternatively indicate the presence of an AGN. The radiation emitted by the accreting black hole would dominate the energy output in the near and mid-infrared, as in most Seyfert 2 galaxies (Alonso-Herrero et al. 2001). Recent near-infrared imaging with HST has shown that all surveyed Seyfert 1 and 50\\% of the Seyfert 2 galaxies contain a luminous unresolved continuum source at 1.6 $\\mu$m (Quillen et al. 2001).  The H-band luminosity of the NGC 4303 nucleus (M$_{F160W}$= $-$17.7) agrees with the average value obtained by Quillen and collaborators for the subsample of Seyfert 2 galaxies with unresolved nuclear sources. There are additional independent indications that an accreting black hole exist in the nucleus of NGC 4303. The ground-based optical spectrum of the nuclear region ($\\sim$ 100 pc) have line ratios anywhere between weak$-$[\\ion{O}{1}] LINERs and Seyfert 2 galaxies (see Table 1). However, the strongest evidence for an accreting black hole might come from recent {\\it Chandra} images (Jim\\'enez-Bail\\'on et al.  2002, in preparation) that reveal the presence of an unresolved hard X-ray (2$-$10 keV) power-law source astrometrically coincident with the UV-bright nucleus, and similar in luminosity to other LLAGNs recently studied with {\\it Chandra} (Ho et al. 2001). Following recent determinations of the black hole mass to velocity dispersion relationship (Ferrarese \\& Merritt 2000; Tremaine et al. 2002), the central velocity dispersion of 74 km s$^{-1}$ (Heraudeau \\& Simien 1998) implies a mass of 1.2 to 2.5 $\\times$ 10$^6$ M$_{\\odot}$ for the central black hole in NGC 4303. Given this mass, the X-ray accreting black hole would radiate very unefficiently, at extremely low sub-Eddington luminosities as in other LLAGNs (Terashima et al. 2000).   In summary, the high spatial resolution multifrequency studies done so far, give the first hints of the very complex structure of the central 10 pc of NGC 4303 where a young, luminous SSC apparently coexists with a low efficiency accreting black hole and with an intermediate/old star cluster. The young SSC is the dominant ionizing source, the accreting black hole is a minor contributor to the overall ionization and the old cluster contributes substantially to the optical and near-ir flux. Some additional red supergiant stars associated with an evolved starburst could also contribute to the near-infrared continuum.     \\subsection{Implications for LLAGNs: the SSC$-$AGN connection}  Low-luminosity AGNs as the one identified in the nucleus of NGC 4303, make up the vast majority of the AGN population (Ho, Filippenko \\& Sargent 1997). The empirical evidence obtained so far indicates that a powerful SSC seems to coexists with an accreting black hole within the central 3 pc of NGC 4303. SSCs as the one detected in NGC 4303 are a common phenomenon in the nuclear regions of early and late-type spirals (Carollo et al. 1997; Carollo, Stiavelli \\& Mack 1998; Carollo et al. 2002; Boeker et al. 2002), and in galaxies with nuclear starbursts and circumnuclear star-forming rings (Meurer et al. 1995; Maoz et al. 1996). Therefore SSCs are a natural consequence of the star formation processes in the nuclear regions of spirals. On the other hand, the tightness of the black hole mass and stellar velocity dispersion relation (Ferrarese \\& Merritt 2000 and references; Tremaine et al. 2002) implies a link between massive black holes (MBH) and bulge formation in galaxies, and therefore nuclear MBHs should also be a natural consequence of the physical processes that formed present-day galaxies. So, it is reasonable to hypothesize that, as detected in the LLAGN NGC 4303, powerful SSCs could coexist with MBHs in the nucleus, i.e. inner few parsecs, of a large fraction of spirals. Under this hypothesis, at least some types of LLAGNs could be understood primarily as a consequence on one hand of the age and mass of the SSC, and on the other hand of the accretion rate and mass of the BH. Massive SSCs with masses of 10$^5$ to 10$^6$ M$_{\\odot}$ would have peak bolometric luminosities of 0.2 to 2 $\\times$ 10$^9$ L$_{\\sun}$ at an age of 3 Myrs. The associated, unobscured, H$\\alpha$ luminosities produced in the ionized interstellar medium surrounding the SSC would be in the 0.14 to 1.4 $\\times$ 10$^{40}$ erg s$^{-1}$ range. On the other hand, low mass black holes such as the one in NGC 4303 (see $\\S$4.2) surrounded by accretion disks radiating at extremely low sub-Eddington luminosities ($\\sim$ 10$^{-4} - 10^{-5}$ L$_E$; Terashima et al. 2000), would emit bolometric luminosities of 10$^6$ to 10$^7$ L$_{\\odot}$, i.e. a factor of hundred less than a young SSC.  Therefore some types of LLAGNs could be understood as the result of the combined ionizing radiation emitted by an evolving super star cluster (i.e. determined by the mass and age) and an accreting black hole (i.e. radiating at sub-Eddington luminosities), coexisting in the inner few parsecs region. The SSC would have an ionizing spectral energy distribution peaking in the soft X-rays/far-UV while the accreting black hole would have a harder radiation field with substantial flux beyond 1 keV.  The nucleus of NGC 4303 with a central young (3-4 Myr old) SSC dominating its ionizing output in the inner few parsecs region, could be a prototype of the class of weak$-$[\\ion{O}{1}] LINERs. The LINER/HII nucleus in NGC 4569 could be another example of SSC-dominated LLAGN (Maoz et al. 1998; Barth \\& Shields 2000; Grabel \\& Bruhweiler 2002). On the other hand, LLAGNs like classical LINERs, or even low luminosity Seyfert 2 nuclei, could still host an evolved, i.e. 10 Myr or older, nuclear stellar cluster that will have a minor contribution to the ionizing luminosity, dominated therefore by an accreting  black hole.    \\subsection{Implications for LLAGNs: measuring the low-end of the black hole mass function}  The complementary multifrequency study of the nucleus of NGC 4303 has shown that stellar clusters of different ages seem to coexist with a black hole in the central few parsecs region of this LLAGN. Even the best 0\\farcs1 spatial resolution available with HST STIS spectrograph represents a linear resolution of 10 pc at a distance of 20 Mpc, not enough to resolve the sphere of influence of nuclear black holes with masses of less than 10$^7$ M$_{\\odot}$. Therefore, the mass contribution of massive young SSCs as the one detected in NGC 4303, and of massive, compact, intermediate/old nuclear stellar clusters in spirals with central low mass black holes (masses of a few million solar masses as in the Milky Way, or as the estimated in NGC 4303, see section $\\S$4.2) could not be negligible and has to be taken into account. Black hole mass measurements are generally done under the assumption that the mass to light ratio of the stars is the same for all spirals, and spatially constant over the region used for the measurement (Sarzi et al. 2001; Sarzi et al. 2002). This might not be a bad assumption for ellipticals, but spirals could have large differences, in M/L, or even gradients within same galaxy, if there are young nuclear clusters of different ages. NGC 4303 is a clear example of a spiral where the assumption of constant M/L would not be valid. Therefore, kinematical studies based on the analysis of optical emission lines alone can not provide the mass contribution of nuclear clusters, and therefore additional detailed multifrequency modeling and spectroscopy with HST would be required before deriving any reliable mass for central black holes with masses of a few to several million solar masses."
    },
    "0207/astro-ph0207058_arXiv.txt": {
        "abstract": "{ The 7 and 15~$\\mu$m observations of selected fields in the Galactic Plane obtained with ISOCAM during  the ISOGAL program offer an unique possibility to search for previously unknown YSOs,  undetected by IRAS because of lower sensitivity or confusion problems.  In a previous paper (Felli et al.~\\cite{FCTOS00}) we established criteria of general validity to select YSOs from the much larger population of Post Main Sequence (Post-MS) stars present in the  ISOGAL fields by comparing radio and IR observations of five fields located at l$\\sim$+45$^{\\circ}$.  The selection was based primarily on  the position of the point sources in the [15] - [7]--[15] diagram, which involves only ISOGAL data and allows to find possible YSOs using the survey data alone. In the present work we revise the adopted criteria by comparing radio-identified UC HII regions and ISOGAL observations over a much larger region.  The main indications of the previous analysis are confirmed, but the criteria for selecting YSO candidates had to be revised to select only bright objects, in order to limit the contamination of the sample by Post-MS stars. The revised criteria ([15]$\\le$4.5, [7]--[15]$\\ge$1.8) are then used to extract YSO candidates from the ISOGAL Point Source Catalogue in preparation. We select a total of 715 YSO candidates, corresponding to $\\sim$2\\% of the sources with good detections at 7 and 15~$\\mu$m.  The results are presented in a table form that provides an unique input list of small diameter, $\\le 6$\\arcsec, Galactic YSO candidates. The  global properties of the sample of YSO candidates are briefly discussed. ",
        "introduction": "The first steps of star formation are the pre-stellar and proto-stellar phases, corresponding to the fragmentation and the gravitational collapse of a dense core in a molecular cloud, before a formed star appears.  The IR spectrum of pre- and proto-stellar cloud cores originates from a very cool dust envelope, and  they are observed in absorption against the diffuse background at 7 and 15~$\\mu$m (Molinari et al.~\\cite{MTBCP}, Bacmann et al.~\\cite{BAPABW}). Some of the infrared dark clouds detected by the ISOGAL and MSX surveys may indeed be the same type of objects (P\\'erault et al.~\\cite{Pea96}; Carey et al.~\\cite{Cea98}).  The next step is commonly referred to as Young Stellar Object (YSO), to indicate the phase when the formed star is deeply embedded in a thick dusty envelope, or when it is hidden by an optically thick disk, remnants of the molecular cloud from which it was formed.  This envelope or disk absorbs all the stellar radiation, making the YSO undetectable in the visible range, and re-emits in the IR, thus making it shine as a bright IR source.  YSOs may have widely different luminosities and masses, ranging from a fraction of a solar mass to 100 solar masses. The YSOs associated with the earliest spectral type stars (earlier than B3) can be searched both in the radio continuum, where the ionised gas of the Ultra Compact HII region (UC HII) produces free-free emission, and in the Middle and Far IR (MIR and FIR, respectively), where the dust emits.  Very bright UC HII regions have been extensively studied (see e.g. Churchwell~\\cite{C91} for a review) and models have been developed to explain the Spectral Energy Distribution (SED) and the spatial morphology of the emission at different wavelengths (Scoville \\& Kwan~\\cite{SK76}, Rowan-Robinson~\\cite{RR80}, Churchwell et al.~\\cite{CWW90}, Ivezi\\'c \\& Elitzur~\\cite{IE97}, Faison et al.~\\cite{FCHHLR98}, Miroshnichenko et al.~\\cite{MIVE}, Feldt et al.~\\cite{Fea99}).  YSOs associated with later spectral type stars (later than B3) can only be detected in the FIR (and sub-mm) thanks to dust emission, as the radio emission decreases sharply since the star doesn't supply enough Lyman continuum photons. When the YSO becomes visible in the optical range, with a luminosity in the range $\\sim$10 to 10$^3$ L$_{\\odot}$, it is called Herbig Ae/Be star. The IR excess comes from a disk, an envelope or a combination of the two (see e.g. Berrilli et al.~\\cite{BCILNS92}, Hillenbrand et al.~\\cite{HSVK92}, Pezzuto et al.~\\cite{PSL97}). At lower luminosity (of the order of one L$_{\\odot}$) an evolutionary track is now well established, going from Class I objects, with considerable IR emission, to class II, when a T Tauri star becomes visible in the optical range but is still surrounded by a disk, and finally to class III when the stellar photosphere becomes visible (Andr\\'e et al.~\\cite{AWTB93}; Lada \\& Wilking~\\cite{LW84}; Lada~\\cite{LADA87}; Lada~\\cite{Lada99} and Natta~\\cite{Natta99} for recent review papers).  Finally, YSOs  with masses in the brown dwarf range have now been detected in the MIR (Olofsson et al.~\\cite{Olofsson99}; Persi et al.~\\cite{Pea00}; Comer\\'on et al.~\\cite{CNK00}). The IR properties of these objects are similar to those of more massive YSOs, suggesting that also in these cases the MIR emission is produced by a circumstellar disk (Natta \\& Testi~\\cite{NT01}, Testi et al.~\\cite{Tea02}). Indeed, even though the spread of luminosities is very large, all these types of YSOs show rather similar SEDs in the IR range, because both the emission mechanism (re-radiation at lower temperature of the stellar emission absorbed by the dust) and the geometry of the dust components are similar. Thus the colours depend mainly on the overall optical depth in the dust but not on the stellar luminosity (Ivezi\\'c \\& Elitzur~\\cite{IE97}). In particular, the systematic study of nearby star forming regions in the frame of the ISOCAM guaranteed time program (Nordh et al.~\\cite{Nordh96}, Bontemps et al.~\\cite{Bont01} and references therein) have shown that, especially for low mass YSOs, there is a clear cut between dusty YSOs of class I and II and the young stars without much dust (class III) at Log$_{10}$[S$_{15}$/S$_7$]~=~--0.2, i.e. [7]--[15]~=~1.1.   The ISOCAM observations of the Galactic Plane carried out during the ISOGAL program offer an unique opportunity for an unbiased search of YSOs. However, their identification requires to be able to separate them from the much larger population of Post-MS  stars in the Galactic Plane, which also have some IR excess due to dust in a circumstellar envelope produced by mass loss, but are in an entirely different evolutionary stage.  In a previous work (Felli et al.~\\cite{FCTOS00}, hereafter Paper~II) a comparison between Very Large Array 3.6 and 6 cm radio continuum observations (Testi et al.~\\cite{TFT99}, hereafter Paper~I) and ISOGAL observations of five Galactic fields  at l $\\sim$ +45$^{\\circ}$, in the two broad band filters LW2 (5.5-8.5~$\\mu$m) and LW3 (12-18~$\\mu$m) was used to establish general criteria that allow the  identification of  YSOs. These criteria were then used to extract the YSOs from the preliminary lists of ISOGAL sources in those fields.  In the present paper we extend the comparison between UC HII regions (or massive YSOs) identified in the radio continuum and ISOGAL observations to a much larger region of the Galactic Plane covered uniformly  with the Very Large Array at 6 cm by the BWHZ survey (Becker et al.~\\cite{bec94}). The much larger sample of radio-identified YSOs allows a better refinement of the identification criteria. In Section~\\ref{scat} the revised criteria are then used to extract the YSOs from the ISOGAL Point Source Catalogue (Omont et al. in preparation, Schuller et al. in preparation). Finally, the galactic distribution and global properties of this sample are briefly discussed. ",
        "conclusions": "In this work we have extended and brought to its conclusion the problem, already approached in Paper~II, of identifying YSO candidates from the much larger population of MIR sources (predominantly Post-MS stars) found during the ISOGAL mapping at 7 and 15~$\\mu$m of selected regions of the Galactic Plane with ISOCAM.  The selection criteria proposed in Paper~II, and verified there from the coincidence of ISOGAL selected bright YSOs with thermal radio continuum sources, have now been tested against a much larger sample of radio-IRAS identified YSOs in the Galactic Plane, by cross-correlating ISOGAL sources detected at 7 and 15~$\\mu$m with the list of UC~HII regions from the VLA 5~GHz Galactic Plane survey of BWHZ. A statistical simulation has been implemented to establish the reliability of the radio-ISOGAL cross-correlation as a function of various parameters. The results confirm that ISOGAL sources with an associated radio-loud YSO occupy a well defined region of the ([7]--[15],[15]) colour-magnitude diagram and that this region is relatively well separated from that occupied by the much larger population of Post-MS stars. A similar segregation of the radio identified YSOs detected also at K-band occurs in the (K--[7],[7]) colour-magnitude diagram. However, the near infrared criteria cannot be used for the entire ISOGAL catalogue since the DENIS observations only cover the fields located at $\\delta\\le2^\\circ$. Additionally, most of the K-band sources associated with radio sources are probably fake associations (see sect.~\\ref{sselc} and Appendix~\\ref{app_rnd}), and even for reliable associations the interpretation of the near-infrared data is complicated by the effect of extinction. Therefore we retained the near infrared information only as ``confidence criteria'' for the sources for which there are available data.  With the aim of providing a more reliable list of YSO candidates throughout the Galactic Plane observations of ISOGAL, following the indications of the comparison with the large sample of radio identified YSOs, we have revised the selection criteria adopted in Paper~II, which have been restricted to [15]$\\le$4.5 and [7]--[15]$\\ge1.8$, with the additional confidence criteria of [7]$\\le$6 and K--[7]$\\ge$4. This choice is motivated by the fact that for larger [15] magnitudes, toward the inner regions of the Galactic Plane, the contamination from reddened Post-MS stars becomes too large to obtain a reasonable list of candidates. The application of those criteria to the entire ISOGAL catalogue (only fields observed both at 7 and 15~$\\mu$m were considered) has produced a list of 715 YSO candidates, which represent $\\sim$2\\% of the total number of ISOGAL point sources detected at both wavelengths.  All the three bright, point-like YSOs identified in Paper~II are reselected using our new criteria (and the revised ISOGAL-PSC).  The galactic distribution of the selected sources is strongly peaked on the Galactic Plane, as expected for very young sources, and show a clear peak close to the Galactic Centre. Our results confirm previous suggestions by BWHZ that the IRAS sample of massive YSOs selected using the WC89 criteria is severely limited by confusion in the inner regions of the Galaxy.  Since the separation between YSOs and Post-MS stars in the colour-magnitude plane uses by necessity a sharp function, we expect a contamination of non-YSOs in our list of candidates of at least 20\\%, as well as we expect that we may have missed some YSOs.  Using our selection criteria we recover $\\sim$70\\% of the radio-identified YSOs. The limitation to bright YSOs ([15]$\\le$4.5) is particularly needed in the ISOGAL fields close to the Galactic Centre, where large line-of-sight extinction is expected, and the contamination from reddened Post-MS stars at high [15] magnitudes is particularly high.  This restriction could be released in less extincted regions, such as the l=+45$^{\\circ}$ fields discussed in Paper~II. In high-extinction fields, a lower contamination fraction could be obtained by increasing the [7]--[15] cutoff limit, at the price of a lower YSO selection efficiency."
    },
    "0207/astro-ph0207434_arXiv.txt": {
        "abstract": "Carbon-12 and carbon-13 abundances have been measured in eleven bright giant members of the globular cluster \\omegacen\\ via observations of the first-overtone CO bands near 2.3 $\\mu$m.  The stars in this sample were selected to span a substantial fraction of the range of iron abundances found in this cluster.  In addition, the sample covers a range of [O/Fe], [Na/Fe] and [Al/Fe] abundance ratios derived in previous studies.  In all $\\omega$ Cen giants the $\\coratio$ abundance ratio is found to be quite low, indicating deep mixing in these red giants.  The mean value for the entire sample is $\\langle\\coratio\\rangle = 4.3 \\pm 0.4$ ($\\sigma = 1.3$), with nine stars equal, within the errors, to the equilibrium ratio $\\coratio = 3.5$ and two stars having slightly higher values.  There is no correlation between the $\\coratio$ and the abundance of iron.  In addition, no correlation of $\\coratio$ with [$^{12}$C/Fe] is found (all giants are deeply mixed), although the derived abundances of [$^{12}$C/Fe] show a positive correlation with [O/Fe], and an anticorrelation with [Na/Fe] (with the oxygen and sodium abundances taken from previous studies in the literature).  A comparison of the isotopic carbon ratios in $\\omega$ Cen with those from other globular clusters (M4, M71, NGC6752, and 47 Tuc), and with literature oxygen abundances, may reveal a slight trend of decreasing $\\coratio$ ratios with decreasing [O/Fe] in the entire globular cluster sample of red giants.  A comparison between $^{12}$C/$^{13}$C and both [Na/Fe] and [Al/Fe], however, reveals no trend. ",
        "introduction": "In the last two decades, it has become clear that evolved giants in metal-poor globular clusters often exhibit a wide range in the abundances of light elements (Paltoglou \\& Norris 1995; Norris \\& Da Costa 1995b; Pilachowski, Sneden, \\& Kraft 1996; Shetrone 1996a, b; Kraft et al. 1997; Sneden et al. 1997; Gonzalez \\& Wallerstein 1998; Ivans et al. 1999; Smith et al. 2000). These variations are generally agreed to arise from proton capture chains that convert C and O into N, Ne into Na, and Mg to Al in the hydrogen-burning layers of red giants.  There has been considerable discussion, however, as to the origin of these abundance patterns (e.g., Cottrell \\& Da Costa 1981; Kraft 1994; Norris \\& Da Costa 1995a; Sneden et al. 1997).  In the ``primordial'' scenario, the abundance patterns reflect nucleosynthesis in a prior generation of massive stars; support for this comes from variations in the C, N, O, Na, Mg, and Al abundances among both main-sequence and turn-off stars in various globular clusters, as found by Bell, Hesser, \\& Cannon (1983), Briley, Hesser, \\& Bell (1991), Suntzeff \\& Smith (1991), Briley et al. (1996), Cannon et al. (1998) or Gratton et al. (2001). In contrast, the ``evolutionary'' scenario envisions that the products of internal proton capture are brought to the stellar surface through mixing in the giants now observed in the clusters. Evidence in favor of this comes from abundance patterns that depend on the evolutionary state of the stars in a cluster.  For example, several systems show a decline in the overall carbon abundance with increasing stellar luminosity and a corresponding rise in the abundance of nitrogen, a sharp decline in the ratio of $\\coratio$ from the presumed primordial value near 90 (i.e., the solar value) to values as low as the nuclear equilibrium value of 3.5.  This equilibrium ratio of $^{12}$C/$^{13}$C= 3.5, which is set largely by the ratio of reaction rates $^{12}$C(p,$\\gamma$)$^{13}$N($\\beta$$^{+}$,$\\nu$)$^{13}$C and $^{13}$C(p,$\\gamma$)$^{14}$N, is insensitive to temperature and is a well-determined astrophysical limit (see the review by Wallerstein et al. 1997).  So-called ``standard'' stellar models evolving up the first ascent of the red giant branch, e.g., Iben (1964) or Charbonnel (1994), predict values to go only as low as $^{12}$C/$^{13}$C$\\sim$ 20: the differences between the observed (low) isotopic carbon ratios and the higher values predicted by the standard models of stellar evolution has been a longstanding piece of evidence used to invoke extra mixing (or deeper mixing) in red giants. Quite possibly, in the globular cluster stars, both primordial and evolutionary chemical evolution occurs, whereby deep mixing in the current giants is superimposed on preexisting abundance patterns (Briley et al. 1994; Ivans et al. 1999).  Mixing signatures are seen also in metal-poor field giants (Gratton et al. 2000, and references therein), with some interesting differences found between the field stars and globular cluster giants of the same overall metallicity.  The O--Na and Mg--Al anticorrelations found in some of the globular clusters are not observed in any of the field red giants.  One effect noted in the field giants by Gratton et al. (2000), and also predicted by current non-standard stellar models (Charbonnel 1994; Wasserburg, Boothroyd, \\& Sackmann 1995; Charbonnel, Brown, \\& Wallerstein 1998), is that the efficacy of first giant branch mixing will increase as the metallicity decreases. The globular cluster $\\omega$ Centauri is a site where the various mixing trends (as a function of [Fe/H] and as a function of Na and Al variations) can be studied in one object.  The most massive globular in the Milky Way, $\\omega$ Cen has a wide spread of [Fe/H]\\footnote{ Throughout this paper, we will use the normal spectroscopic notation that ${\\rm [A/B]} \\equiv \\log_{10}(N_A / N_B)_{\\rm star} - \\log_{10}(N_A / N_B)_\\odot$, for elements A and B.} (Suntzeff \\& Kraft 1996; Norris, Freeman, \\& Mighell 1996, and references therein), unique among globulars, with the possible exception of M22 (which may show an abundance spread, but much smaller than in $\\omega$ Cen).  In $\\omega$ Cen the distribution of [Fe/H] has a floor near --2.0 to --1.8, presumably representing the initial metallicity of the gas out of which the cluster was formed, and shows an extended tail to higher metal abundances.   The abundances of most elements increase with [Fe/H] and exhibit relatively small scatter, with the exception of O, Na, Al and some other light elements (Brown \\& Wallerstein 1993; Norris \\& Da Costa 1995b; Smith et al. 2000).  As in other globular clusters, the giants of $\\omega$ Cen have [Na/Fe] and [Al/Fe] abundances which are positively correlated with each other, but anticorrelated with [O/Fe] in a way that is most easily (but not definitively!) explained as a product of deep mixing. The abundance patterns of elements heavier than Fe suggest nucleosynthesis from AGB stars between 1 and 3$M_\\odot$, but whether this occurred during formation of the cluster that involved self-enrichment, or from mergers of fragments with different chemical abundances, is not yet determined.  To date, only a few measurements of the $\\coratio$ isotopic ratio exist for $\\omega$ Cen giants (3 from Brown \\& Wallerstein 1993). Their 3 values were all quite low, with $^{12}$C/$^{13}$C= 6, 4, and 4; however, the metallicities sampled by these 3 giants was limited ([Fe/H]= -1.36, -1.34, and -1.26). The goal of this work is to derive additional values of $\\coratio$ in a larger sample of $\\omega$ Cen giants, in order to investigate if there are any trends with [Fe/H], or if this somewhat unique globular cluster differs in its $\\coratio$ ratios when compared to other clusters. ",
        "conclusions": "The $^{12}$C/$^{13}$C ratios in all of the $\\omega$ Cen giants are quite low, with a mean $\\langle$$^{12}$C/$^{13}$C$\\rangle$ = 4.3$\\pm$0.4: this indicates extensive and deep mixing in these low-mass, low-metallicity giants.  This result is in general agreement with previous studies of globular cluster giants, as well as low-metallicity field giants.  In particular for the $\\omega$ Cen giants, all of which have M$_{\\rm V}$$\\le$ -1.8, there is no meaurable change in $^{12}$C/$^{13}$C over the metallicity range of [Fe/H]= -1.7 to -0.7.  It is found that [$^{12}$C/Fe] correlates well with [O/Fe] and anticorrelates with [Na/Fe], and that the overall trends among these various ratios agrees very well between $\\omega$ Cen and other globular clusters.  The $^{12}$C/$^{13}$C ratios themselves exhibit no measurable trends with [Na/Fe] or [Al/Fe], but may hint at a positive correlation with [O/Fe].  \\medskip We wish to thank the staff of the Cerro Tololo Inter-American Observatory, especially M.\\ Fern\\'andez and H.\\ Tirado, for their excellent assistance with the observations.  Partial support for this project came from the National Science Foundation under grants AST-9157038 and INT-9215844 to The Ohio State University Research Foundation. VVS acknowledges support from the National Science Foundation through grant AST-9987374."
    },
    "0207/astro-ph0207602_arXiv.txt": {
        "abstract": "We observed the compact central object CXOU J085201.4--461753 in the supernova remnant G266.2--1.2 (RX J0852.0--4622) with the \\chan\\/ ACIS detector in timing mode. The spectrum of this object can be described by a blackbody model with the temperature $kT=404\\pm 5$ eV and radius of the emitting region $R=0.28\\pm0.01$ km, at a distance of 1~kpc. Power-law and thermal plasma models do not fit the source spectrum. The spectrum shows a marginally significant feature at 1.68 keV. Search for periodicity yields two candidate periods, about 301 ms and 33 ms, both significant at a 2.1$\\sigma$ level; the corresponding pulsed fractions are 13\\% and 9\\%, respectively. We find no evidence for long-term variability of the source flux, nor do we find extended emission around the central object. We suggest that CXOU J085201.4--461753 is similar to CXOU J232327.9+584842, the central source of the supernova remnant Cas A. It could be either a neutron star with a low or regular magnetic field, slowly accreting from a fossil disk, or, more likely, an isolated neutron star with a superstrong magnetic field. In either case, a conservative upper limit on surface temperature of a 10 km radius neutron star is about 90 eV, which suggests accelerated cooling for a reasonable age of a few thousand years. ",
        "introduction": "The shell-like supernova remnant (SNR) G266.2--1.2 (also known as RX~J0852.0--4622, or ``Vela Junior'') at the south-east corner of the Vela SNR was discovered by Aschenbach (1998) in the {\\sl ROSAT} All-Sky Survey data. Possible detection of the 1.156 MeV $\\gamma$-ray line of the radioactive isotope $^{44}$Ti (half-life $\\sim 90$ yr) with the Compton Gamma-Ray Observatory (Iyudin et al. 1998) may imply a very young SNR age of $\\sim 680$ yr, at a distance of $\\sim 200$ pc. Aschenbach, Iyudin, \\& Sch\\\"onfelder (1999) estimated upper limits of 1100 yr for the age, and 500 pc for the distance. Observations with {\\sl ASCA} (Tsunemi et al. 2000; Slane et al. 2001) demonstrate that the X-ray spectra of the SNR shell are nonthermal. Fits of these spectra with a power-law (PL) model yield a hydrogen column density substantially higher than that for the Vela SNR, implying a plausible distance to the remnant of 1--2 kpc, and an age of a few thousand years.   Aschenbach (1998) suggests that G266.2--1.2 was created by a core-collapse supernova that left a compact remnant --- a neutron star (NS) or a black hole (BH). Three compact remnant candidates have been reported from the observations with {\\sl ROSAT} (Aschenbach 1998; Aschenbach et al.\\ 1999), {\\sl ASCA} (Slane et al.\\ 2001), and {\\sl Beppo-SAX} (Mereghetti 2001). Pavlov et al.\\ (2001) observed G266.2--1.2 with the \\chan\\ Advanced CCD Imaging Spectrometer (ACIS) and found only one bright X-ray source, CXOU\\, J085201.4--461753 (J0852 hereafter), close to the SNR center. They measured the source position with accuracy better than $2''$ and proved that J0852 is not an X-ray counterpart of bright optical stars in the field. Follow-up optical observations (Pavlov et al.\\ 2001; Mereghetti, Pelizzoni, \\& De Luca 2002a) revealed an object located only $2\\farcs4$ south-west of the J0852. The colors of the optical source  are consistent with those of a main sequence star at a distance of 1.5--2.5 kpc; most likely, this is a  field star. The limiting optical magnitudes at the position of the X-ray source ($B>22.5$, $R>21$ --- Pavlov et al.\\ 2001; $B>23$, $R>22.5$ --- Mereghetti et al.\\ 2002a) rule out the possibility that the X-ray source is an AGN. The lack of variability combined with the X-ray spectral properties makes a cataclysmic variable interpretation also implausible. The nature of the source remains elusive, although an isolated cooling NS or a NS with a ``fallback'' disk seem to be possible interpretations.  The large frame time, 3.24 s, of the previous snapshot (3 ks) ACIS observation made it impossible to search for short periods and led to strong saturation (pile-up) of the source image, precluding an accurate spectral analysis. To search for  pulsations from the compact source and obtain a more accurate spectrum, we observed J0852 with  {\\sl Chandra} ACIS with a time resolution of 2.85 ms. We present the results of this observation in \\S2 and discuss the nature of the source in \\S3. ",
        "conclusions": "The X-ray data and optical limits indicate that J0852 is the compact remnant (NS or BH) of the supernova explosion. The X-ray spectral properties and the lack of radio emission (Duncan \\& Green 2000) suggest that J0852 is not an active pulsar. Furthermore, the {\\sl Chandra} observations show no sign of a pulsar-wind nebula (PWN) around the point source. From the 3 ks observation in Timed Exposure mode (Pavlov et al.\\ 2001), the $3\\sigma$ upper limit on the PWN brightness (in counts arcsec$^{-2}$) can be estimated as $3(b/A)^{1/2}$, where $b=0.029$ counts arcsec$^{-2}$ is the background surface brightness, and $A$ is the (unknown) PWN area. Scaling the area as $A=1000 A_3$ arcsec$^2$ (which corresponds to the transverse size of about $5\\times 10^{17} A_3^{1/2}$ cm) and assuming a PL spectrum with a photon index $\\gamma=1.5$--2, we obtain an upper limit of (1.3--$2.0)\\times 10^{30} A_3^{1/2} d_1^2$ erg s$^{-1}$ on the PWN luminosity in the 0.2--10 kev band, for $n_{\\rm H, 21}$ in the range of 1.4--5.3.  The observational properties of J0852 strongly resemble those of the other radio-quiet central compact objects (CCOs) in SNRs (see Pavlov et al.\\ 2002a for a review), particularly the CCO in the SNR Cas A (Murray et al.\\ 2002, and references therein). At least one of these sources, 1E 1207.4--5209, has been proven to be a NS rotating with a period of 424 ms (Zavlin et al.\\ 2000; Pavlov et 2002b). A number of possible interpretations of CCOs have been recently discussed by several authors (e.g. Pavlov et al.\\ 2000, 2001, 2002a; Chakrabarty et al.\\ 2001). The limits on X-ray-to-optical flux ratio for J0852 and the Cas A CCO virtually rule out models which involve accretion onto a NS or a BH from a binary companion. If these are accreting objects, a more plausible source of accreting matter might be a ``fossil disk'', left over after the SN explosion (van Paradjis, Taam, \\& van den Heuvel 1995). Alternatively, thermal emission from an isolated, cooling NS could explain the observational results. We discuss these two options below.  \\subsection{Accretion-powered X-ray pulsar?}  If J0852 is an accreting NS, the observed luminosity, $L_{\\rm x}\\sim 2\\times 10^{32} d_1^2$ erg s$^{-1}$, could be due to a rather low accretion rate, $\\dot{m} \\sim 1.5\\times 10^{12} R_6 M_1^{-1} d_1^2$ g~s$^{-1}$, where $R_6=R_{\\rm NS}/(10^6\\, {\\rm cm})$, $M_{1}=M/M_\\odot$. The accreting matter could be supplied from a fossil (``fallback'') disk. The formation of such a disk from the ejecta produced by a SN explosion was discussed by a number of authors (e.g. Marsden, Lingenfelter, \\& Rothschild 2001, and references therein). Some models suggest that a fossil disk can be formed several days after the SN explosion (``prompt'' disk) and range from $0.001M_{\\odot}$ to $0.1M_{\\odot}$, while others suggest that the disk can be formed later, years after the SN explosion (``delayed'' disk). The details of the formation mechanism and the disk properties are highly uncertain, and, consequently, the accretion rate $\\dot{m}$ is also poorly constrained, but the required value of $\\sim 10^{12}$ g s$^{-1}$ is low enough not to exhaust the disk at any reasonable age of J0852.  The accretion onto a NS can proceed in two different regimes (e.g., Frank, King, \\& Raine 1992), depending on the relation between the corotation radius, $R_c=1.5\\times 10^8 P^{2/3} M_1^{1/3}$ cm, and the magnetospheric radius, $R_M=3.5\\times10^{9}B_{12}^{4/7}\\dot{m}_{12}^{-2/7} M_1^{-1/7}R_6^{12/7}$ cm, where $P$ is the NS spin period, $B=10^{12} B_{12}$ G is the magnetic field at the NS surface, and $\\dot{m}_{12}=\\dot{m}/(10^{12}\\, {\\rm g}\\, {\\rm s}^{-1})$. If $R_{M}>R_{c}$, the infalling material is stopped at the magnetospheric radius and  expelled as a wind due to centrifugal force. In this ``propeller regime'' (Illarionov \\& Sunyaev 1975), X-ray emission is mainly due to optically thin thermal bremsstrahlung produced in the flow (Wang \\& Robertson 1985). Since the thermal bremsstrahlung model does not fit the observed spectrum, we consider this case unlikely.   If $R_{M}<R_{c}$ (i.e., $P\\gtrsim 10^2 B_{12}^{6/7} \\dot{m}_{12}^{-3/7} M_1^{-5/7} R_6^{18/7}$ s, or $B\\lesssim 4\\times 10^9 P^{7/6} \\dot{m}_{12}^{1/2} M_1^{5/6}R_6^{-3}$ G), the accreting matter is able to reach the NS surface. At extremely low magnetic fields, $B \\lesssim 6\\times 10^5 \\dot{m}^{1/2} M_1^{1/4} R_6^{-5/4}$ G, when the magnetospheric radius is smaller than the NS radius, a hot layer is formed at the boundary between the accretion disk and the NS surface (e.g., Frank et al.\\ 1992). Since this boundary layer is expected to be optically thin at $\\dot{m}\\ll 10^{16}$ g~s$^{-1}$ (Inogamov \\& Sunyaev 1999), its radiation cannot explain the observed BB spectrum. At reasonable magnetic fields, $B\\gg 10^6 \\dot{m}^{1/2} M_1^{1/4} R_6^{-5/4}$ G ($R_M\\gg R_{\\rm NS}$), the accretion flow is channeled onto the NS poles, producing hot spots of radius $a\\sim R_{\\rm NS}^{3/2}/R_{M}^{1/2} \\sim 0.17 B_{12}^{-2/7} \\dot{m}_{12}^{1/7} M_1^{1/14} R_6^{9/14}$ km. The observed size and temperature of the BB-like radiation are consistent with being emitted from such a cap at $B\\sim 10^{11}$ G. Such an estimate requires a pulsar period $P\\gtrsim 10$ s, much longer than our candidate periods. If we assume $P=301$ ms, the condition $R_M<R_c$ requires $B\\lesssim 10^9\\, \\dot{m}_{12}^{1/2} M_1^{5/6} R_6^{-3}$ G and $a\\gtrsim 1.2\\, \\dot{m}_{12}^{1/7} M_1^{1/14} R_6^{9/14}$ km, considerably larger than the size of emitting region, $R\\approx 0.3 d_1$ km, inferred from the BB fit. However, given the crudeness of the polar cap size estimate, which can be much smaller than adopted above (see, e.g., Frank et al.\\ 1992, and references therein), we cannot rule out the candidate period of 301 ms based on the apparent inconsistency between $a$ and $R$. Thus, in the accretion hypothesis, J0852 could be a low-luminosity X-ray pulsar, presumably with a magnetic field much lower than those of binary X-ray pulsars, slowly accreting from a fossil disk. An argument against this interpretation is a lack of nonperiodic variability in the radiation from J0852, which is commonly observed from accreting sources (at least, X-ray binaries). On the other hand, variability could be found in further observations of this source. A direct confirmation of the accreting hypothesis would be detection of an accretion disk, which would require deep IR-optical observations with high angular resolution.   \\subsection{Isolated cooling neutron star?} One can also assume that J0852 is an isolated (non-accreting) NS emitting thermal radiation from its surface. The ``standard'' NS cooling models predict a luminosity of $\\sim (0.5$--$2)\\times 10^{34}$ erg s$^{-1}$ for a NS of 0.1--10 kyr age (e.g., Tsuruta 1998). The lower observed luminosity of J0852 could be interpreted as due to an accelerated cooling mechanism, but applicability of the cooling models to J0852 is questionable because the models assume a uniformly heated NS surface while the size of the emitting region obtained from the BB fit is only $\\approx 0.3\\, d_{1}$ km.  Apparent sizes of the emitting regions much smaller than the canonical NS radius have been observed from other isolated NSs (Pavlov et al.\\ 2002a,c). In particular, the Cas A CCO shows a (blackbody) size of 0.3 km with a temperature of 0.6 keV (Pavlov et al.\\ 2000), which hints that it is an object similar to J0852, with a higher temperature possibly due to its younger age. Pavlov et al.\\ (2000) suggested a two-component thermal model for the Cas~A CCO, in which the observed X-rays are emitted from hydrogen polar caps of about 1 km radius and 0.24 keV effective temperature, while the rest of the NS surface is iron at a temperature of 0.15 keV, too cold to be observable because of strong interstellar absorption. In this model, the polar caps are hotter because of the higher thermal conductivity of hydrogen. Weaker ISM absorption for J0852 allowed us to find a lower temperature limit for the cold component, $<90$ eV, too low to explain the temperature difference by different chemical compositions. It should be mentioned that this limit is a factor of 1.4 lower than the temperature predicted by the so-called standard (slow) cooling model for a $10^{3}$ yrs old NS (see, e.g., Fig.\\ 4 in Slane, Helfand, \\& Murray 2002). This may indicate that if J0852 is a NS, it undergoes fast cooling, perhaps associated with direct Urca processes in the NS core (e.g., Yakovlev et al.\\ 2002).   Hot spots on the NS surface could also be associated with a very strong magnetic field, $B\\gg 10^{13}$ G. Due to anisotropic heat conductivity of the NS crust, the surface temperature is higher at the magnetic poles (Greenstein \\& Hartke 1983; Shibanov \\& Yakovlev 1996). To produce small hot spots, the surface magnetic field should be strongly nonuniform (e.g., an offset dipole or a quadrupole --- Page \\& Sarmiento 1996). Fast decay of a superstrong magnetic field ($B\\gtrsim 10^{14}$ G) could provide an additional source of polar cap heating (Thompson \\& Duncan 1996; Colpi, Geppert \\& Page 2000). In such strong magnetic fields,  electron-positron pair creation should be suppressed due to photon splitting (Baring \\& Harding 2001), which is consistent with the apparent lack of pulsar activity in J0852.   One can crudely estimate the magnetic field assuming that one of the two candidate periods, 33 ms or 301 ms, is the true period. If the initial period of the pulsar was much shorter than the current period, then the period derivative, rotation energy loss rate, and ``canonical'' magnetic field [$B\\equiv 3.2\\times 10^{19}(P\\dot{P})^{1/2}$ G], can be estimated as $\\dot{P}=3.2\\times 10^{-11} P[(n-1) t_3]^{-1}$, $\\dot{E}=1.25\\times 10^{36} P^{-2} [(n-1)t_3]^{-1}$ erg~s$^{-1}$, and $B=1.8\\times 10^{14} P [(n-1) t_3]^{-1/2}$ G, where $t=10^3 t_3$ yr is the NS age, and $n$ is the braking index. Assuming $n=2.5$, (close to that observed in young pulsars), we obtain, for $P=33$ ms, $\\dot{P}=7.0\\times 10^{-13} t_3^{-1}$, $\\dot{E}=8.6\\times 10^{38} t_3^{-1}$ erg~s$^{-1}$, and  $B= 4.9\\times 10^{12} t_3^{-1/2}$ G --- parameters typical for a young, active pulsar, in apparent contradiction with observations. On the other hand, for the 301 ms period, we obtain $\\dot{P}=6.4\\times 10^{-12} t_3^{-1}$, $\\dot{E}=9.2\\times 10^{36} t_3^{-1}$ erg~s$^{-1}$, and $B= 4.4\\times 10^{13} t_3^{-1/2}$ G. Since the local magnetic field can be much higher than the canonical value (e.g., for an offset dipole), one can speculate that, for $P=301$ ms, it is high enough to explain the hot spot(s) and the lack of radio-pulsar activity. If this hypothesis is correct, the J0852 could be a very young Anomalous X-ray Pulsar (AXP) whose period will become of order 6--12 s (as observed in AXPs) when it grows older by a factor of 20--40.  However, there are considerable differences between the properties of AXPs and J0852. Contrary to AXPs, whose spectra contain both the BB and PL components of comparable luminosities (Mereghetti et al.\\ 2002b), the spectrum of J0852 fits well with a single BB model. The size of the emitting region in J0852 is substantially smaller (0.3 km vs.\\ 0.7--5 km), and the temperature somewhat lower (0.4 keV vs.\\ 0.4--0.6 keV), than those of AXPs. These differences (particularly, the lack of a PL component in J0852) hint at different NS parameters. For instance, it is quite possible that none of the candidate periods is correct, and the true period is even longer than the AXP periods. In this case, the magnetic field could be even higher than those adopted in the magnetar interpretation of AXPs --- e.g., $B=2.4\\times 10^{15} (P/20\\,{\\rm s}) t_3^{-1}$ G. Such a strong field can inhibit not only the pair cascade, but also the emission of primary particles from the NS surface, which might explain the lack of particles in the NS magnetosphere (hence, the lack of nonthermal radiation) in J0852. If the NS rotates sufficiently slow, $P \\gtrsim 0.5 B_{15}^{4/15}(Z/26)^{-12/15}$ s, the critical parallel electric field required to pull out electrons from the NS surface, $E_{\\parallel,\\rm{crit}}\\approx 2.7 \\times 10^{12} (Z/26)^{6/5} B_{15}^{3/5}$ V cm$^{-1}$ (Usov \\& Melrose 1995), is higher than the maximum parallel electric field at the surface, $E_{\\parallel,\\rm{max}}\\approx 1 \\times 10^{10} B_{15}(P/20\\,{\\rm s})^{-3/2}$ V cm$^{-1}$. On the other hand, the surface temperature, $kT_e\\approx 0.5 (Z/26)^{4/5}B_{15}^{2/5}$ keV, above which the thermoionic emission of electrons becomes efficient (Usov \\& Melrose 1995), grows with increasing magnetic field. (These estimates assume that the NS has no light-element [e.g., hydrogen] atmosphere.) Therefore, a long period and a superstrong magnetic field might explain the lack of the PL tail in the spectrum of J0852 and other enigmatic CCOs (e.g., in the Cas A and Pup A SNRs;\\  Pavlov et al. 2002a).  In summary, the observations of J0852 can be explained assuming it is a NS. Given the deep limiting optical magnitudes and the lack of nonperiodic variability, we consider the interpretation in terms of an isolated NS with a very strong magnetic field somewhat more plausible than the accretion models. Further observations are required to confirm or reject this hypothesis. Particularly important would be X-ray timing observations to measure the period unequivocally, high-resolution X-ray spectral observations to look for spectral features, and IR-optical observations to search for a NS counterpart (e.g., a fossil disk)."
    },
    "0207/astro-ph0207373_arXiv.txt": {
        "abstract": "An imaging survey of the {\\em IRAS} 1-Jy sample of 118 ultraluminous infrared galaxies was conducted at optical ($R$) and near-infrared ($K^\\prime$) wavelengths using the University of Hawaii 2.2m telescope. The methods of observation and data reduction are described. An $R$ and $K^\\prime$ atlas of the entire sample is presented along with some of the basic astrometric and photometric parameters derived from these images. A more detailed analysis of these data is presented in a companion paper (Veilleux, Kim, \\& Sanders 2002). ",
        "introduction": "Ultraluminous infrared galaxies (ULIGs; log [$L_{\\rm IR}$/$L_\\odot$] $\\ge$ 12 by definition; $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and $q_0$ = 0) are an important class of extragalactic objects that appear to be powered by a mixture of starburst and AGN activity, both of which are fueled by an enormous supply of molecular gas that has been funneled into the nuclear region during the merger of gas-rich spirals. ULIGs rival quasi-stellar objects (QSOs) in bolometric luminosity, and there is speculation that ULIGs may indeed represent an important stage in the formation of QSOs as well as powerful radio galaxies, and that they may also represent a primary stage in the formation of giant ellipticals (see Sanders \\& Mirabel 1996 for a review).  A complete sample of 118 ULIGs ($f_{60} > 1$) was compiled by Kim (1995; see also Kim \\& Sanders 1998) from a redshift survey of objects in the {\\em IRAS} Faint Source Catalog (FSC: Moshir et al. 1992). As the nearest (median redshift = 0.145) and brightest ULIGs, this so-called 1-Jy sample provides the best list of objects for detailed multiwavelength studies. Optical spectroscopy for 108 objects in the sample has been published by Veilleux, Kim, \\& Sanders (1999a), and near-infrared spectra have been published for 60\\% of the sample by Veilleux, Sanders, \\& Kim (1999b). We now present the results from our optical and near-infrared ground-based imaging study of this sample.  In this first paper, we discuss the methods of observation and data reduction (\\S 2) and then present an $R$ and $K^\\prime$ atlas of the 1-Jy sample (\\S 3.1) along with some of the basic astrometric (\\S 3.2) and photometric (\\S 3.3) parameters derived from these images. In Paper II (Veilleux, Kim, \\& Sanders 2002, we examine the data in more detail and combine the results of the image analysis with those from our spectroscopic survey and other published studies to address the important issue of the origin of ULIGs and their possible evolutionary link with QSOs, radio galaxies, and ellipticals.  For both this and the companion papers, we adopt $H_0$ = 75 km s$^{-1}$, $q_0$ = 0.0, $M^\\ast_R = -21.2$ mag. and $M^\\ast_{K^\\prime} = -24.1$ mag. The values for $M^\\ast$ are justified in \\S 1 of Paper II. ",
        "conclusions": "We have carried out an optical ($R$) and near-infrared ($K^\\prime$) imaging survey of the IRAS 1-Jy sample of 118 ULIGs.  In this first paper, we present an atlas of the sample galaxies and list a few basic astrometric and photometric parameters derived from these images. Nearly all ULIGs show signs of tidal interaction, in agreement with previous studies. The 1-Jy ULIGs show a broad range of luminosities with median $R$-band ($K^\\prime$-band) integrated luminosities of about 2 (3) $L^\\ast$. Their nuclear and, to a lesser extent, global $R - K^\\prime$ colors are significantly redder than those of normal galaxies, a result which is attributed to dust reddening and emission in the inner 4 kpc of these galaxies. This is also consistent with the fact that the brightness distributions of 1-Jy ULIGs are more compact at $K^\\prime$ than at $R$. The generally smaller nuclear offsets between $R$ and $K^\\prime$ images among Seyfert 1s may indicate lower dust extinction in these objects, but this result is only marginally significant.  A more detailed analysis of these data is presented in a companion paper (Veilleux, Kim, \\& Sanders 2002; Paper II).   \\clearpage"
    },
    "0207/astro-ph0207659_arXiv.txt": {
        "abstract": "Using the StarTrack population synthesis code we compute the distribution of masses of merging compact object (black hole or neutron star) binaries. The shape of the mass distribution is sensitive to some of the parameters governing the stellar binary evolution. We discuss the possibility of constraining stellar evolution models using mass measurements obtained from the detection of compact object inspiral with the upcoming gravitational-wave observatories. ",
        "introduction": "The  operational phase of the  gravitational wave detectors  LIGO\\cite{1992Sci...256..325A} and VIRGO \\cite{1990brada} is quickly approaching, and in a few years we anticipate the operation  of the more sensitive advanced LIGO detector. Mergers of compact objects in binaries are the most promising sources of gravitational waves to be detected by these detectors. The expected rates of such mergers have been a subject of intense research. Recent  analyses \\cite{2001ApJ...556..340K,2002ApJ...572..407B,KKL} have shown there is a little chance  that LIGO should see any mergers of compact object binaries, while LIGO~II could detect a substantial number, up to hundreds or thousands of such events. These estimated rates however, do vary strongly with the model of stellar binary evolution.  Once we begin to detect such events, the data analysis should yield information about each system that merged. The merging of compact object binaries proceeds in three phases: inspiral, merger, and ringdown. It is the inspiral phase that will be used for detection of the stellar masses involved.  The analysis of the frequency change with time of the event should yield the masses of the merging objects\\cite{will}. However, determining their type, i.e. a neutron star or a black hole will be difficult\\cite{blanchet} with the measurement during this phase. The modifications of the gravitational wave signal due to the mass distribution of the neutron star are small and may amount to one in $10^4$ oscillations.  The analysis of the observation of merger phase may lead to determination of the type of the compact object i.e. will be different for mergers involving neutron star, and may even lead to constraints on the neutron star equation of state. Also, if the neutron stars are rotating there is a chance of exciting stellar oscillations by the orbital motion which may modify the gravitational waveform\\cite{1999MNRAS.308..153H}.  It is then important to ask what is the expected two dimensional distribution of masses of compact object binaries that merge, and how does this distribution depend on the assumed model of stellar binary evolution. In \\S\\,2 we present the dependence of the mass ratio distribution on the particular stellar evolution model. Next, \\S\\,3 is devoted to the potential observations. The main observable derived from the gravitational-wave signal is the chirp mass. We show the dependence of the distribution of the observed chirp masses on the stellar evolution parameters and estimate the possibility of distinguishing between the models using the amount of data that we expect to be obtained. Finally in \\S\\,4 we summarize our results.   \\begin{figure}[t] \\centerline{ \\includegraphics[width=0.9\\textwidth]{ryson.ps}} \\caption{The distributions of intrinsic mass ratio in compact objects produced in the standard model simulation (left panel), and the contributions of different types of binaries (BH-BH thick line, NS-NS thin line , and BH-NS dotted line)to this distribution. The distributions  have been calculated using a total of 106307 compact object binaries and the  bin  width is $dq=0.01$. Each  distribution is normalized to unity. } \\label{rysone} \\end{figure} ",
        "conclusions": "Using the StarTrack population synthesis code we show that  the distribution of masses of merging compact objects carries  useful information about the ways massive stars evolve. In particular the stellar wind strength, common envelope efficiency,  loss of angular momentum efficiency, and the fraction of mass accreted in mass transfer events leave distinct tracks on the  distribution of observed chirp masses. The Monte Carlo simulations show that analysis of this distribution shall yield  useful  constraints on the evolutionary parameters of binaries containing massive stars. The advanced LIGO configuration will  be able to gather sufficient number of  merging events to perform such analysis."
    },
    "0207/astro-ph0207145_arXiv.txt": {
        "abstract": "We have developed a model in which the diffuse synchrotron emission from radio mini--halos, observed in some cooling flow clusters, is due to a relic population of relativistic electrons reaccelerated by MHD turbulence \\textit{via} Fermi--like processes. In this model the energetics is supplied by the cooling flow itself. Here, the model (successfully applied to the Perseus cluster, A426) is preliminarily applied to the possible mini-halo candidate A2626, for which we present VLA data. ",
        "introduction": "Several clusters of galaxies show extended ($\\sim$ Mpc size) synchrotron emission not directly associated with the galaxies but rather diffused into the intracluster medium (ICM): these radio sources are called radio halos. In some cooling flow clusters with a central dominant galaxy, the diffuse radio emission is extended on a smaller scale, forming the so--called \\textit{mini--halos} (Feretti \\& Giovannini 1996). \\\\ The radiative life--time of an ensemble of relativistic electrons losing energy by synchrotron emission and Inverse Compton (IC) scattering off the CMB photons is given by $ \\tau (\\mbox{yr}) = 24.3/[(B^2 + B_{CMB}^2) \\, \\gamma]$, where $B$ is the magnetic field intensity (in G), $\\gamma$ is the Lorentz factor and $B_{CMB} = 3.18 (1+z)^2 \\mu$G. In a cooling flow region (i.e. for distances $r < r_c$, the cooling radius) the compression of the thermal ICM is expected to produce a significant increase of the strength of the frozen--in intracluster magnetic field: $B \\propto r^{-2}$ for radial compression (Soker \\& Sarazin 1990) or $B \\propto r^{-0.8}$ for isotropic compression (Tribble 1993). Therefore, in absence of a reacceleration or continuous injection mechanisms, relativistic electrons injected at a given time in these intense fields (of order of some tens of $\\mu$G, e.g. Ge \\& Owen 1993) should already have lost most of their energy and the radio emission would not be observable for more than $\\sim 10^{7\\div8}$ yr. This short lifetime contrasts with the diffuse radio emission observed in mini--halos, hence it seems plausible that the electrons have been reaccelerated. ",
        "conclusions": ""
    },
    "0207/astro-ph0207235_arXiv.txt": {
        "abstract": "{ We present results of our diffraction-limited mid-infrared imaging of the massive star-forming region W3(OH) with SpectroCam--10 on the 5-m Hale telescope at wavelengths of 8.8, 11.7, and 17.9\\,\\micron{}.  The thermal emission from heated dust grains associated with the ultracompact \\hii{} region W3(OH) is resolved and has a spatial extent of $\\sim$2\\arcsec{} in the N band. We did not detect the hot core source W3(H$_2$O) which implies the presence of at least 12 mag of extinction at 11.7\\,\\micron{} towards this source. These results together with other data were used to constrain the properties of W3(OH) and W3(H$_2$O) and their envelopes by modelling the thermal dust emission. ",
        "introduction": "The ``hot cores'' revealed in recent years by molecular line investigations are small ($\\lesssim 0.1$\\,pc), very dense ($n\\gtrsim 10^7$cm$^{-3}$), and hot ($> 100$\\,K) entities of giant molecular clouds (e.g., Cesaroni et al. \\cite{cesa98}).  They are considered to be the likely birthplaces of massive stars (e.g., Garay \\& Lizano \\cite{garay99}, Kurtz et al. \\cite{kurtz00}). Hot cores are frequently associated with ultracompact \\hii{} regions (\\uchii s) which are more evolved and accessible to near-infrared (NIR) studies of their stellar population (e.g., Feldt et al. \\cite{feldt98}, Henning et al. \\cite{henning01}). Although the temperatures and sizes of hot cores suggest that they might be conspicuous objects in the infrared sky, extinction in the mid-infrared (MIR, N (10\\,\\micron{}) and Q (20\\,\\micron{}) bands) caused by the large column densities ($N$(H$_2)\\gtrsim 10^{23}$\\,cm$^{-2}$) must not be neglected.  Conclusions concerning the heating and the stellar content of hot cores have to be based on the knowledge of their luminosity. This quantity is difficult to estimate since the immediate neighbourhood of \\uchii s often leads to source confusion, especially in the far-infrared (FIR) range where these objects emit most of their energy and the angular resultion of spaceborn observations is as yet poor. Radio  interferometry at mm/submm wavelengths allow to  separate the dust continuum emission of the hot core from the free-free radiation of the adjacent \\uchii. Such measurements were used to constrain models of hot cores (Osorio et al. \\cite{osorio99}). High-resolution ground-based MIR observations, on the other hand, provide information on the spectral energy distribution (SED) shortward of the peak flux. A corresponding study of the IRc2 source in the Orion BN/KL complex (Gezari et al. \\cite{gezari98}) illustrates their importance. More recently, \\cite{2002ApJ...564L.101D} were able to detect MIR emission from the hot core G29.96$-$0.02 with a morphology similar to that of the warm ammonia. We performed high resolution MIR imaging of hot cores, including W3(H$_2$O) and the neighbouring \\uchii{} W3(OH), in order to measure their flux densities or to provide at least upper limits. While results for G10.47+0.03 will be the subject of a forthcoming paper (Pascucci et al., in prep.), we present here our findings for W3(H$_2$O) and the \\uchii{} W3(OH).   The \\uchii{} W3(OH) is very well-studied in the radio domain by continuum and molecular line observations. It is located at the distance of 2.2\\,kpc (Humphreys \\cite{hum78}) and harbours numerous OH masers. The hot core W3(H$_2$O), also known as Turner-Welch object (TW, Turner \\& Welch \\cite{tw84}), is situated $\\sim6$\\,\\arcsec{} east of W3(OH). % This enigmatic source shows an outflow traced by the proper motion of H$_2$O masers (Alcolea et al. \\cite{alco92}) and is associated with a double-sided radio continuum jet, presumably of synchrotron nature (Reid et al. \\cite{reid95}). Recent interferometric imaging at 220\\,GHz by Wyrowski et al. (\\cite{wyrowski99}) revealed another % source in the immediate neighbourhood of W3(H$_2$O), suggesting that the region harbours a cluster of protostars. The thermal infrared emission from W3(OH) has been studied by Keto et al. (\\cite{keto92}) using one of the first MIR array cameras (Arens et al. \\cite{arens87}). Keto et al. claimed the detection of W3(H$_2$O) at the wavelength of 12.2\\,\\micron{} with a flux density of $45 \\pm 10$\\,mJy. W3(H$_2$O) and W3(OH) were the subject of a continuum and molecular line study of \\cite{2000ApJ...537..283V}. Since their results are based on single-dish data which do not resolve the two objects, they cannot be directly compared to our model presented below. ",
        "conclusions": "\\subsection{The \\uchii{} W3(OH)} The FWHM of the MIR emission from W3(OH) was derived from Gaussian fits, taking the size of the diffraction-limited beam into account. The sizes and flux densities of W3(OH) are listed in Tab. \\ref{prop}. The fluxes are based on an aperture of 4\\arcsec{} diameter. The extent of the MIR emission originating from the thermal emission of heated dust grains exceeds that of the 8.4\\,GHz radio continuum (FWHM 1\\farcs52), indicating that the warm dust is more extended than the ionized gas. The variation of the angular size in dependence on wavelength can be approximated as FWHM($\\lambda$)\\,$\\sim \\lambda^{0.6\\pm0.2}$ and results from the decline of the temperature with increasing distance from the heating star(s).  The comparison of our flux densities of W3(OH) with other estimates allows conclusions on the influence of different beam sizes. For this purpose, we retrieved the IRAS-LRS spectrum, identified W3(OH) in the MSX point source catalog (Egan et al. \\cite{egan99}), and retrieved an ISO-LWS spectrum from the data archive.   The LRS spectrum was integrated according to the applied passbands. It is obvious that the 8.8\\,\\micron{} fluxes given in Tab.\\ref{prop} considerably exceed our value. This can be explained by ubiquitous emission attributed to Polycyclic Aromatic Hydrocarbons (PAHs) surrounding the \\uchii{} which strongly contributes to the flux in the large apertures of IRAS and MSX. Pronounced 7.7 and 8.6\\,\\micron{} PAH bands can be misleading in ground-based derivations of the optical depth of the 9.7\\,\\micron{} silicate feature (e.g., Roelfsema et al. \\cite{roelf96}).     \\subsection{The hot core W3(H$_2$O)} Our attempt to detect  W3(H$_2$O) in the IR was stimulated by the presence of an outflow. Generally, IR emission can escape in outflow lobes primarily due to scattering  (e.g., Fischer et al. \\cite{fisch96}). An example is NGC6334~I(N), a presumed high-mass Class 0 object, for which Sandell (Sandell \\cite{sandell00}) rendered the detection of IR emission impossible because of the extremely high extinction derived from mm/submm maps. However, this source is associated with NIR emission (Tapia et al. \\cite{tapia96}, Megeath \\& Tieftrunk \\cite{megeath99}) obviously originating from the blue-shifted lobe of its outflow.  The flux densities from W3(H$_2$O) in the absence of any intervening absorbing matter can be estimated from the temperature map given by Wyrowski et al. (\\cite{wyrowski97}). The expected peak surface brightness amounts to 2170\\,Jy/$\\sq$\\arcsec{} at 11.7\\,{\\micron}. Together with our 3$\\sigma$ sensitivity, this yields a lower limit to the extinction at this wavelength of 12\\,mag. Our failure to detect this source is consistent with the high molecular hydrogen column densities of 1\\dots3.5$\\times$10$^{24}$\\,cm$^{-2}$ inferred from molecular line and continuum investigations (Turner \\& Welch \\cite{tw84}, Wyrowski et al. \\cite{wyrowski97}). It suggests that the molecular outflow of W3(H$_2$O) is very young, i.e. did not fully penetrate the hot core yet, and, in addition, might be in the plane of the sky. In fact, Fig.\\,1 from Wyrowski et al. (\\cite{wyrowski99}) shows that the jet is confined to the region of the hot core. The moderate expansion velocity of the H$_2$O masers of 20\\,km\\,s$^{-1}$ (Alcolea et al. \\cite{alco92})  implies a dynamical timescale of only 500\\,yr which is consistent with the upper limit on proper motions of the radio jet of 150\\,km\\,s$^{-1}$ (Wilner et al. \\cite{wilner99}). These velocities are low compared to those of thermal radio jets (Anglada \\cite{anglada96}) and indicate that the outflow is presumably slowed-down by the high-density environment.         \\begin{enumerate} \\item The thermal infrared counterpart to the hot core source W3(H$_2$O) was not detected at any wavelength of our observations. This revises the finding of Keto et al. (\\cite{keto92}) which most probably resulted from confusing W3(H$_2$O) with the \\uchii{} northeast of W3(OH). We derived a lower limit for the extinction at 11.7\\,\\micron{} towards W3(H$_2$O) of 12\\,mag. \\item Our diffraction-limited imaging resolved the thermal emission from W3(OH) and clearly indicates a wavelength dependence of its apparent size. \\item The comparison of our 8.8\\,\\micron{} flux to those measured with IRAS and MSX led to the conclusions that PAHs surround W3(OH). \\item In accordance with our thermal infrared imaging and mm/submm studies, the ISO-LWS spectrum of W3(OH) might be decomposed in two components with W3(OH) being the hotter, more evolved, object while W3(H$_2$O) dominating at mid/far-infrared wavelengths. \\end{enumerate}"
    },
    "0207/astro-ph0207003_arXiv.txt": {
        "abstract": "We present the results of a long-look monitoring of 3C~273 with $RXTE$ between 1996 and 2000. A total of 230 observations amounts to a net exposure of 845~ksec, with this spectral and  variability analysis of 3C~273 covering the longest observation period available at hard X-ray energies. Flux variations by a factor of 4 have been detected over 4~years, whereas less than 30$\\%$ flux variations have been observed for individual flares on time-scales of $\\sim$\\,3~days. Two temporal methods, the power spectrum density (PSD) and the structure function (SF), have been used to study the variability characteristics of 3C~273. The hard X-ray photon spectra generally show a power-law shape with a differential photon index of $\\Gamma$$\\simeq$ 1.6$\\pm$0.1. In 10 of 261 data segments, exceptions to power-law behaviour have been found: (i) an additional soft excess below 4~keV, and (ii) a broad Fe fluorescent line feature with $EW$ $\\sim$ 100$-$200 eV. Our new observations of these previously reported X-ray features may imply that 3C~273 is a unique object whose hard X-ray emission occasionally contains a component which is not related to a beamed emission (Seyfert like), but most hard X-rays are likely to originate in inverse Compton radiation from the relativistic jet (blazar like). Multi-frequency spectra from radio to $\\gamma$-ray are presented in addition to our $RXTE$ results. The X-ray time variability and spectral evolution are discussed in the framework of beamed, synchrotron self-Compton picture. We consider the ``power balance'' (both radiative and kinetic) between the accretion disk, sub-pc-scale jet, and the 10~kpc-scale jet. ",
        "introduction": "As the brightest and nearest ($z$ = 0.158) quasar, 3C~273 is the ideal laboratory for studying active galactic nuclei (AGN). Studies of this source are relevant to all AGN physics, as 3C~273 displays significant flux variations, has a well-measured wide-band spectral energy distribution, and has a relativistic jet originating in its central core (see Courvoisier 1998 for a review). VLBI radio observations of the pc-scale jet have shown a number of jet components moving away from the core at velocities apparently faster than the speed of the light (e.g., Pearson et al.\\ 1981; Vermeulen \\& Cohen 1994). The collimated jet structure extends up to $\\sim$\\,50~kpc from the core. Since 3C~273 is bright at all wavelengths and on various scale sizes, it provides a valuable opportunity to probe the most inner part of the accretion disk ($\\sim$\\,10$^{-4}$~pc) as well as the large scale jet ($\\sim$\\,10$^4$~pc) at the same time.  3C~273 is generally classified as a blazar and is also a prominent $\\gamma$-ray source. It was the only extra-galactic source of gamma-rays identified in COS-B observations (Swanenburg et al. 1978), and was subsequently detected at energies in the 0.05\\,MeV to 10\\,MeV range with OSSE (McNaron-Brown et al. 1995), in the 0.75\\,MeV to 30\\,MeV range with COMPTEL (Sch\\\"onfelder et al. 2000), and above 100\\,MeV in numerous EGRET observations (Hartman et al. 1999). EGRET observations helped establish that the overall spectra of blazars (plotted as $\\nu$$F_{\\nu}$) have two pronounced continuum components: one peaking between IR and X-rays, and the other in the $\\gamma$-ray regime (e.g., Mukherjee et al. 1997). The low energy component is believed to be produced by synchrotron radiation from relativistic electrons in magnetic fields, while inverse-Compton scattering by the same electrons is thought to be dominant process responsible for the high energy $\\gamma$-ray emission (Ulrich, Maraschi, \\& Urry 1997). The radiation is emitted from a relativistic jet, directed close to our line of sight (e.g., Urry \\& Padovani 1995).  3C~273 is no exception to this picture, but an additional $Big$ $Blue$ $Bump$ (hereafter $BBB$\\,) dominates the optical--soft-X-ray emission (see Paltani, Courvoisier, \\& Walter 1998 ; Robson 1996 for a review). Interestingly, although similar excesses have been observed in the optical-UV region of Seyfert galaxies, they have not been reported for any blazar other than 3C~273. Although its origin is still far from being understood, it has been proposed that $BBB$ may be due to thermal emission from the surface of a standard accretion disk (Shields 1978), including optically thin parts of the disk and a corona (see Czerny 1994 for a review). Various models have been suggested (e.g., Courvoisier \\& Clavel 1991), but it is generally agreed that the $BBB$ is nearly isotropic emission from the vicinity of the central black hole, presumably from the accretion disk.  The presence of a fluorescent emission line at 6.4\\,keV is a signature of X-ray reprocessing by cold material. There is some evidence for the presence of a weak line at this energy in the X-ray spectrum of 3C~273. One of the $Ginga$  observations showed evidence for the line at 99$\\%$  level, but the other observations provided only upper limits (Turner et al. 1990). A line detection is also reported by Grandi et al. (1997) and Haardt et al. (1998) in a $BeppoSAX$ observation in 1996. $ASCA$ provided only upper limits in a 1993 observation (Yaqoob et al. 1994), whereas a broad line feature was clearly detected in 1996 observations. This somewhat confused situation probably means that both thermal (as for Seyferts) and non-thermal (as for blazars) emission processes are taking place in this particular object. However, it is completely unknown (i) which process dominates the radiation and (ii) how often the Fe line is ``visible'' in the photon spectrum.  Many $\\gamma$-ray blazars, including 3C~273, have shown large flux variations either, on time-scales as short as a day for some objects (e.g., von Montigny et al. 1997; Mukherjee et al. 1997). However, there are only a few blazars whose variability characteristics are well studied in the high energy bands (both X-ray and $\\gamma$-ray bands). For example, recent X-ray studies of Mrk 421, the proto-typical ``TeV emitting'' blazar,  have revealed that (i) the variability time scale is $\\sim$ 1~day, and (ii) the flux variation in the X-ray and the TeV $\\gamma$-ray bands is well correlated on time-scales of a day to years (see Takahashi et al. 2000; Kataoka  et al. 2001). Clearly, these observations provide important clues to understanding jet physics, and potentially for discriminating between various emission models for blazars. Although 3C~273 is a bright object and particularly well-sampled, it has not yet been possible to undertake such studies since data from previous X-ray satellites were too sparse (e.g., the 13 observations spanning over 5 years of Turner et al. 1990).  In this paper, we analyze the archival hard X-ray data obtained with $RXTE$ between 1996 and 2000, with a total exposure of 845~ksec. This report of both the temporal and spectral variability of 3C~273 is thus based on the highest quality and most densely sampled data in this energy band. The observation and data reduction are described in $\\S$2. Two temporal methods are introduced; the power spectrum density ($\\S$3.1) and the structure function ($\\S$3.2). The X-ray spectral evolutions are summarized in $\\S$$\\S$4.1 and 4.2. Multi-frequency spectra are presented in $\\S$4.3. In $\\S$5, we discuss scenarios which systematically account for the hard X-ray variability and spectral evolution of 3C~273. Throughout the paper, we discuss the energetics between the central engine and the relativistic jet, taking into account observations of the large scale jet by $Einstein$, $ROSAT$ and $Chandra$. Finally, in $\\S$6, we present our conclusions.  Throughout this paper, we adopt $H_0 = 75$~km\\,s$^{-1}$\\,Mpc$^{-1}$ and $q_0 = 0.5$. The luminosity distance to the source is $d_{\\rm L} = 2.02 \\times 10^{27}$~cm. ",
        "conclusions": "We have analyzed the archival $RXTE$ data available for 3C~273 between 1996 and 2000. A total of 230 observations amounts to the net exposure of 845~ksec over the 4 years.  This is the longest, and most densely sampled, exposure for this object in the hard X-ray band. Both the PSD and the SF show a roll-over with a time-scale on the order of $\\sim$\\,3~days, although the lower-frequency (i.e., longer time-scale) variability is still unclear. We found that the variability time-scale of 3C~273 is similar to those observed in TeV gamma-ray emitting blazars, whereas the variability amplitude is an order of magnitude smaller. Considering that the hard X-ray spectra of 3C~273 generally maintains a constant power-law shape with $\\Gamma \\simeq 1.6 \\pm 0.1$, beamed, inverse Compton emission inside the jet is the most likely the origin of the X-ray/$\\gamma$-ray emission. Two kinds of exceptions have been found, which may be interpreted either as the hard-tail of the $BBB$ or the emission from the accretion disk $occasionally$ superposed on the jet emission. From a multi-frequency analysis, we constrain the physical quantities relevant for the jet emission. We argue that, (i) the kinetic power carried by $relativistic$ electrons corresponds to only 5$-$10\\,$\\%$ of the disk luminosity, and (ii) the various powers in the sub-pc-scale jet are ranked $L_{\\rm kin} \\ge L_{\\rm rad} \\ge L_{B}$. The connection between the sub-pc-scale jet and the 10~kpc scale jet remains uncertain, but our work suggests that the most of the jet power might be $hidden$ at the base of the sub-pc-scale jet, and effectively released at the 10~kpc scale via a completely different mechanism of energy dissipation."
    },
    "0207/astro-ph0207529_arXiv.txt": {
        "abstract": "The quenching factor for proton recoils in a stilbene scintillator was measured with a $^{252}$Cf neutron source and was found to be 0.1 -- 0.17 in the recoil energy range between 300~keV and 3~MeV. It was confirmed that the light yield depends on the direction of the recoil proton.  The directional anisotropy of the quenching factor could be used to detect the wind of the WIMPs caused by the motion of the earth around the galactic center. ",
        "introduction": "There is substantial evidence that most of the matter in our Galaxy must be dark matter and exist in the form of Weakly Interacting Massive Particles (WIMPs) \\cite{balyon}. WIMPs can be directly detected through elastic scattering with nuclei in a detector. The typical recoil energy is of order of 10~keV. The detector is required to measure about 10 keV nuclear recoil energies and discriminate between nuclear recoils and $\\gamma$-rays to reject background events.  For reliable detection of WIMPs, conventional searches depend mainly on observing the annual modulation of the counting rate arising from the orbital motion of the earth around the sun and the rotation of the solar system itself around the galactic center. However, this approach is subject to a small modulation amplitude and a large systematic error induced by long-term changes of the condition of the detector and seasonal variations of environmental radiations. For an alternative method, it might be useful to observe the wind of the WIMPs caused by the motion of the earth around the galactic center \\cite{dR/dE,diurnal} (Fig. \\ref{wimp2d}). The effect of the wind is large ($\\simeq$ 244~km/s) and a small systematic error can be achieved by means of periodically rotating the detector. To detect the wind of the WIMPs, a directional detector is required.   It is known that organic crystal scintillators such as anthracene and stilbene have directional anisotropies in their scintillation responses to heavy charged particles \\cite{anthHC,zp2}. Protons and carbon ions are produced internally by WIMPs scattering and the scintillation responses depend on the direction of their motion. These anisotropies can be used to obtain directional information of the WIMPs \\cite{nonisotropic}.  In using scintillation detectors to observe WIMPs, one needs to measure the scintillation efficiency or the quenching factor of the nuclear recoil. For organic crystal scintillators, the anisotropies of the scintillation responses are expressed in the directional variations of the quenching factor.  Measurements of the quenching factor for nuclear recoils have usually been performed with several MeV mono-energetic neutrons produced by nuclear reactions using accelerators \\cite{qF-NaI,qF-CaF2}. The nuclear recoil was caused by elastic scattering with neutrons instead of WIMPs. We performed the measurement employing a $^{252}$Cf neutron source which is easier to handle.  In this paper, we report on the measurement of the quenching factor of a stilbene scintillator as a function of the recoil energy and its dependence on the direction of the proton recoil. ",
        "conclusions": "The quenching factor for proton recoils in a stilbene scintillator was measured with a $^{252}$Cf neutron source and was found to range from 0.10 to 0.17 with recoil energy from 300~keV to 3~MeV. The measurement was performed for two recoil directions and we confirmed that the quenching factor depends on the recoil direction. It was smaller for the recoils perpendicular to the cleavage plane than for the recoils parallel to it. A measurement of the scintillation responses for $\\alpha$-particles was reported in Ref. \\cite{zp2}, where the responses for the two directions are qualitatively consistent with our result.  For the carbon recoil, definite data were not obtained because of the low $E_{\\rm visible}$ and the large background at low energies which were induced by accidental coincidences from $\\gamma$-rays. To measure the quenching factor in the low $E_{\\rm visible}$ region, we would need to improve the signal-to-background ratio.  For the dark matter search, one can obtain directional information through comparing the data collected with different orientations of the detector. We are planning to measure two energy spectra directing the crystal axes of maximum and minimum scintillation responses toward the direction of the motion of the earth around the galactic center."
    },
    "0207/nucl-th0207067_arXiv.txt": {
        "abstract": "Recent observations by the Chandra observatory suggest that some neutron stars may cool rapidly, perhaps by the direct URCA process which requires a high proton fraction. The proton fraction is determined by the nuclear symmetry energy whose density dependence may be constrained by measuring the neutron radius of a heavy nucleus, such as $^{208}$Pb. Such a measurement is necessary for a reliable extrapolation of the proton fraction to the higher densities present in the neutron star. A large neutron radius in $^{208}$Pb implies a stiff symmetry energy that grows rapidly with density, thereby favoring a high proton fraction and allowing direct URCA cooling. Predictions for the neutron radius in $^{208}$Pb are correlated to the proton fraction in dense matter by using a variety of relativistic effective field-theory models. Models that predict a neutron ($R_n$) minus proton ($R_p$) root-mean-square radius in $^{208}$Pb to be $R_n\\!-\\!R_p\\alt0.20$~fm have proton fractions too small to allow the direct URCA cooling of 1.4~$M_\\odot$ neutron stars. Conversely, if $R_n\\!-\\!R_p\\agt0.25$~fm, the direct URCA process is allowed (by all models) to cool down a 1.4~$M_\\odot$ neutron star. The Parity Radius Experiment at Jefferson Laboratory aims to measure the neutron radius in $^{208}$Pb accurately and model independently via parity-violating electron scattering. Such a measurement would greatly enhance our ability to either confirm or dismiss the direct URCA cooling of neutron stars. ",
        "introduction": "Neutron stars are created with very high temperatures in supernova explosions. Indeed, neutrinos observed from SN1987A indicate a neutrinosphere temperature that could be as high as 5 MeV~\\cite{raf}. Neutrons stars then cool, primarily by neutrino emission~\\cite{Pe92}. In the standard scenario, the modified URCA reaction, \\begin{equation} n + n \\rightarrow n + p + e^- + \\bar\\nu_{e}, \\label{modified} \\end{equation} emits neutrinos from the volume of the star. This process, however, is relatively slow as a second nucleon is necessary to conserve momentum.  Recent X-ray observations of the neutron star in 3C58~\\cite{3c58}, Vela~\\cite{vela}, and Geminga~\\cite{geminga} indicate low surface temperatures. Moreover, the low quiescent luminosity in the transiently accreting binaries KS 1731-260~\\cite{wijnands} and Cen X-4~\\cite{page} suggest rapid cooling. As X-ray observatories progress and our knowledge of neutron-star atmospheres and ages improve, additional ``cold'' neutron stars may be discovered. Such low surface temperatures appear to require enhanced cooling from reactions that proceed faster than the modified URCA process of Eq.~(\\ref{modified}). Enhanced cooling may occur via the weak decay of additional hadrons such as pion or kaon condensates~\\cite{kaon}, hyperons~\\cite{Pr92}, or quark matter~\\cite{quarks}. Yet perhaps the most conservative enhanced-cooling mechanism is the direct URCA process~\\cite{urca,leinson} of neutron beta decay followed by electron capture: \\begin{subequations} \\begin{eqnarray} && n \\rightarrow p + e^{-} + \\bar\\nu_{e}\\;, \\label{urca1}\\\\ && e^{-} + p \\rightarrow n + \\nu_{e}\\;.     \\label{urca2} \\end{eqnarray} \\end{subequations} This mechanism is not ``exotic'' as it only requires protons, neutrons, and electrons---constituents known to be present in dense matter. However, to conserve momentum in Eq.~(\\ref{urca1}) the sum of the Fermi momenta of the protons plus that of the electrons must be greater than (or equal to) the neutron Fermi momentum. This requires a relatively large proton fraction.  Yakovlev and collaborators~\\cite{yakovlev} are able to reproduce measured neutron-star temperatures using a relativistic mean-field equation of state that allows direct URCA for neutron stars with masses above $1.358~M_\\odot$ ($M_{\\odot}\\!=\\!$~solar mass). In contrast, Tsuruta and collaborators~\\cite{t5} rely on pion condensation to reproduce the measured temperatures. They argue that microscopic calculations of neutron-rich matter~\\cite{fp} using nonrelativistic nucleon-nucleon interactions yield too small a proton fraction for the URCA process to operate. Unfortunately, these microscopic calculations depend on a poorly known three-nucleon force and on relativistic effects that could end up increasing the proton fraction at high densities.  Superconductivity and superfluidity can greatly influence neutron-star cooling~\\cite{Ya99,Pa00}.  For temperatures much lower than the pairing gap, pairing correlations suppress exponentially the rate of many cooling reactions. Yet for temperatures of the order of the pairing gap, the thermal breaking and subsequent reformation of nucleon ``Cooper'' pairs promotes an additional neutrino-emission mechanism that rapidly cools the star~\\cite{pairing}. However, it has been argued in Ref.~\\cite{yakovlev} that this mechanism alone is unlikely to explain the low temperature of some neutron stars. This is because for a large enough neutron-pairing gap, pair breaking would rapidly cool all neutron stars at a rate almost independent of the mass of the star.  This would disagree with observations of some warm neutron stars.  Tsuruta and collaborators have claimed that microscopic calculations with a high proton concentration show a small proton pairing gap~\\cite{tpairing}. If so, a direct URCA process (one not controlled by pairing correlations) will cool a star so quickly that thermal radiation would become invisible~\\cite{t5}. However, we caution that drawing definitive conclusions from microscopic calculations of pairing gaps may be premature, as significant uncertainties remain in the interactions, equation of state, composition, and phases of high-density matter.  Although the precise mechanism remains unknown, some kind of enhanced cooling appears to be required to explain the recent observations of cold neutron stars. While the need for exotic phases of matter is appealing, more conventional cooling scenarios, such as the direct URCA process, can not be dismissed on purely theoretical grounds. Moreover, neutron-star observations alone may not be able to resolve the detailed mechanism of enhanced cooling. Thus, we consider complementary laboratory experiments that could help us confirm (or possibly dismiss) the direct URCA process. This can be achieved by constraining the symmetry energy of dense matter. The symmetry energy describes how the energy of (asymmetric) nuclear matter increases as one departs from equal numbers of neutrons and protons. The proton fraction $Y_{p}\\!=\\!Z/A$ of nuclear matter in beta-equilibrium is sensitive to the symmetry energy~\\cite{urca}. A large symmetry energy imposes a stiff penalty on the system for upsetting the $N=Z$ balance hereby forcing it to retain a large proton fraction.  Energetic heavy-ion collisions probe the symmetry energy at high nuclear densities~\\cite{li}. Possible observables include the ratio of $\\pi^-$-to-$\\pi^+$ production and the neutron-proton differential collective flow. However, these reactions may suffer from important uncertainties associated with the complex strong interactions of the heavy-ion collisions. Thus, we rely on a purely electroweak reaction that can be unambiguously interpreted. Parity violating elastic electron scattering from a heavy nucleus is sensitive to the neutron density. This is because the weak charge of a neutron is much larger than the weak charge of a proton. The Parity Radius Experiment at the Jefferson Laboratory aims to measure the neutron radius in $^{208}$Pb to a 1\\% accuracy ($\\pm 0.05$~fm)~\\cite{prex}. This measurement can be both accurate and model independent~\\cite{bigpaper}. In Ref.~\\cite{brown} Brown showed that the neutron radius of $^{208}$Pb determines the pressure of neutron-rich matter at normal densities which, in turn, is related to the density dependence of the symmetry energy~\\cite{prakash}. In an earlier work we showed how the neutron radius of $^{208}$Pb determines properties of the neutron-star surface, such as the transition density from a solid crust to a liquid interior~\\cite{prl}. Furthermore, we argued that by comparing the neutron radius of $^{208}$Pb (a low-density observable) to the radius of a neutron star (a high and low-density observable) evidence may be provided in support of a phase transition in dense matter~\\cite{radii}.  In the present work we show how the neutron radius of a heavy nucleus (such as $^{208}$Pb) controls the density dependence of the symmetry energy. Unfortunately, the density dependence of the symmetry energy ($da_{\\rm sym}/d\\rho$) is poorly known.  Thus a measurement of the neutron radius of $^{208}$Pb seems vital, as it will constrain the density dependence of the symmetry energy at low density. This, in turn, will allow a more reliable extrapolation of the symmetry energy, and thus a more reliable determination of the proton fraction at the higher densities required in the study of neutron-star structure. While in principle collective modes of nuclei, such as the giant-dipole or isovector-monopole resonances, are sensitive to $da_{\\rm sym}/d\\rho$, in practice this sensitivity is small. Moreover, the parameter sets used in the calculations (see various tables) have been adjusted so that well known ground-state properties remain fixed while changing the neutron radius. This shows that existing ground-state information, such as charge densities or binding energies, do not determine the neutron radius uniquely. Thus the need for a new measurement---such as the neutron radius in $^{208}$Pb---that will provide important information on $da_{\\rm sym}/d\\rho$.  The paper has been organized as follows. In Sec.~\\ref{sec:formalism}, relativistic effective-field theories for both dense matter and finite nuclei are discussed. A large number of parameter sets are considered so that the density-dependence of the symmetry energy may be changed while reproducing existing ground-state data. In Sec.~\\ref{sec:results}, results for the equilibrium proton fraction as a function of baryon density are presented using interactions that predict different neutron radii in $^{208}$Pb. Our summary and conclusions are offered in Sec.~\\ref{sec:conclusions}. In particular, we conclude that for models with a large neutron skin in $^{208}$Pb ($R_{n}\\!-\\!R_{p}\\agt0.25$~fm) the symmetry energy rises rapidly with density and the direct URCA cooling of a $1.4~M_\\odot$ neutron star is likely. Conversely, if the neutron radius is small ($R_{n}\\!-\\!R_{p}\\alt 0.20$~fm) it is unlikely that the direct URCA process occurs. In this case, the enhanced cooling of neutron stars may indeed require the presence of exotic states of matter, such as meson condensates, hyperonic, and/or quark matter. ",
        "conclusions": "\\label{sec:conclusions} Recent X-ray observations suggest that some neutron stars cool quickly. This enhanced cooling could arise from the direct URCA process---that requires a high proton fraction---or from the beta decay of additional hadrons in dense matter, such as pions, kaons, hyperons, or quarks. Yet, it seems unlikely that the X-ray observations alone will determine the origin of the enhanced cooling.  In this work we propose to use a laboratory experiment to constrain the direct URCA process in neutron stars. The Parity Radius Experiment at the Jefferson Laboratory~\\cite{prex,bigpaper} aims to measure the neutron radius of $^{208}$Pb accurately and model independently via parity-violating electron scattering.  For the direct URCA process to be realized, the equilibrium proton fraction in the star must be large. The equilibrium proton fraction is determined by the symmetry energy, whose density dependence can be strongly constrained through a measurement of the neutron radius in $^{208}$Pb. Such a measurement could provide a reliable extrapolation of the proton fraction to higher densities. Thus, predictions for the neutron radius in $^{208}$Pb have been correlated to the proton fraction in dense neutron rich matter by using a wide range of relativistic effective-field theory models.  We find that models with a neutron skin in $^{208}$Pb of $R_n\\!-\\!R_p\\alt0.20$~fm generate proton fractions that are too small to allow the direct URCA process in 1.4~$M_\\odot$ neutron stars. Conversely, if $R_n\\!-\\!R_p\\agt0.25$~fm, then all models predict the URCA cooling of 1.4~$M_\\odot$ stars.  While this paper has focused on relativistic effective field-theory models, we expect our conclusions to be general and applicable to other approaches, both relativistic and nonrelativistic. For example, the nonrelativistic equation of state of Friedman and Pandharipande~\\cite{fp} predicts too small a proton fraction for URCA cooling to be possible. Moreover, this equation of state yields a neutron skin in $^{208}$Pb of only $R_n\\!-\\!R_p\\!=\\!0.16\\pm0.02$~fm~\\cite{brown}. Thus, these results are fully consistent with Fig.~\\ref{Figure3} that predicts no URCA cooling for such a small value of $R_n\\!-\\!R_p$.  The equation of state considered in this work consists of matter composed of neutrons, protons, electrons, and muons in beta equilibrium; no exotic component was invoked. Further, no explicit proton or neutron pairing was considered. Nucleon superfluidity is an accepted phenomenon in nuclear physics and superfluid gaps are important for the cooling of neutron stars~\\cite{yakovlev}. Thus, the study of pairing gaps in relativistic effective-field theories is an important area of future work; first steps in this direction have been taken in Ref.~\\cite{super}. In particular, the proton pairing gap in matter with a high proton concentration must be computed~\\cite{t5}.  In summary, the feasibility of enhanced cooling of neutron stars via the direct URCA process was studied by correlating the proton fraction in dense, neutron-rich matter to the neutron skin of $^{208}$Pb. Thus, a measurement of the neutron radius in $^{208}$Pb may become vital for confirming (or dismissing) the direct URCA cooling of neutron stars. If direct URCA cooling is ruled out, then observations of enhanced cooling may provide strong evidence in support of exotic states of matter, such as meson condensates and quark matter, at the core of neutron stars."
    },
    "0207/astro-ph0207286_arXiv.txt": {
        "abstract": "The recent measurements of the power spectrum of Cosmic Microwave Background anisotropies are consistent with the simplest inflationary scenario and big bang nucleosynthesis constraints. However, these results rely on the assumption of a class of models based on primordial adiabatic perturbations, cold dark matter and a cosmological constant. In this paper we investigate the need for deviations from the $\\Lambda$-CDM scenario by first characterizing the spectrum using a phenomenological function in a $15$ dimensional parameter space. Using a Monte Carlo Markov chain approach to Bayesian inference and a low curvature model template we then check for the presence of new physics and/or systematics in the CMB data. We find an almost perfect consistency between the phenomenological fits and the standard $\\Lambda$-CDM models. The curvature of the secondary peaks is weakly constrained by the present data, but they are well located. The improved spectral resolution expected from future satellite experiments is warranted for a definitive test of the scenario. ",
        "introduction": "The recent observations of the cosmic microwave background (CMB) anisotropies power spectrum (\\cite{toco},\\cite{b97}, \\cite{Netterfield},\\cite{halverson},\\cite{lee}, \\cite{cbi}, \\cite{vsa}, \\cite{archeops}, \\cite{acbar}, \\cite{ruhl},\\cite{vsae}) have presented cosmologists with the possibility of studying the large scale properties of our universe with unprecedented precision. As is well known (see e.g. \\cite{review}), the structure of the theoretical CMB spectrum, given mainly by the relative positions and amplitude of the so-called acoustic peaks, is sensitive to several cosmological parameters. The existing CMB data sets are therefore being analyzed with increasing sophistication (see \\cite{koso} and \\cite{sko} for important advancements) in an attempt to measure the undetermined cosmological quantities. The most recent analyses of this kind (\\cite{debe2001}, \\cite{pryke}, \\cite{stompor}, \\cite{wang}, \\cite{cbit}, \\cite{vsat},\\cite{bean},\\cite{saralewis},\\cite{mesilk}, \\cite{archeops2}, \\cite{slosar},\\cite{wang2}) have revealed an outstanding agreement between the data and the inflationary predictions of a flat universe and of a primordial scale invariant spectrum of adiabatic density perturbations. Furthermore, the CMB constraint on the amount of matter density in baryons $\\omega_b$ is now in very good agreement with the independent constraints from standard big bang nucleosynthesis (BBN) obtained from primordial deuterium (see e.g. \\cite{burles}, \\cite{hansen}) and consistent within $2$-$\\sigma$ with those derived from the combined analysis of $^4He$ and $^7Li$ (\\cite{cyburt}). Finally, the detection of power around the expected third peak, on arc-minutes scales, provides a new and independent evidence for the presence of non-baryonic dark matter (\\cite{mesilk}).  The data therefore suggests that our present cosmological model represents a beautiful and elegant theory able to explain most of the observations.  However, the CMB result relies on the assumption of a particular class of models, based on adiabatic, {\\it passive} and {\\it coherent} (see \\cite{andy}) primordial fluctuations, and cold dark matter. In the following we refer to this class of models as $\\Lambda$-Cold Dark Matter ($\\Lambda$-CDM).  This weak point, shared by most of the current studies, should not be overlooked: it has been recently shown, for example, that the very legitimate inclusion of gravity waves (see e.g. \\cite{efstathiou}, \\cite{gw}) or isocurvature modes (\\cite{kxm}, \\cite{trotta}, \\cite{amendola}) into the analysis can completely erase most of the constraints derived from CMB alone.  Furthermore, since even more exotic modifications like quintessence (\\cite{caldwell}), topological defects (\\cite{bouchet},\\cite{dkm}), broken primordial scale invariances (\\cite{alexandra}, \\cite{bend}, \\cite{covi}), extra dimensions (\\cite{bisilk}) or unknown systematics (just to name a few) can be in principle considered, one should be extremely cautious in making any definitive conclusion from the present CMB observations.  It is therefore timely to investigate if the present CMB data are in complete agreement with the $\\Lambda$-CDM scenario or if we are losing relevant scientific informations by restricting the current analysis to a subset of models (see e.g. \\cite{tegza}).  In the present {\\it paper} we check to what extent modifications to the standard $\\Lambda$-CDM scenario are {\\it needed} by current CMB observations with two complementary approaches: First, we provide a model-independent analysis by fitting the data with a phenomenological function and characterizing the observed multiple peaks. Phenomenological fits have been extensively used in the past and recent CMB analyses (\\cite{rocha}, \\cite{page}, \\cite{miller2k2}, \\cite{podariu}, \\cite{boghdan}, \\cite{douspis}). Our analysis differs in two ways: we include the latest CMB data from the Boomerang (\\cite{ruhl}), VSAE (\\cite{vsae}), ACBAR (\\cite{acbar}), and Archeops (\\cite{archeops}) experiments and we make use of a Monte Carlo Markov Chain (MCMC) algorithm, which allows us to investigate a large number of parameter simultaneously ($15$ in our case).  We then compare the position, relative amplitude and width of the peaks with the same features expected in a $4$-parameters model template of $\\Lambda$-CDM spectra. By doing a peak-by-peak comparison between the theory and the phenomenological fit which is based on a much wider set of parameters, we then verify in a systematic way the agreement with the standard theoretical expectations.  As a by-product of the analysis, we present a set of cosmological diagrams that directly translate, under the assumption of $\\Lambda$-CDM, the constraints on the features in the spectrum into bounds on several cosmological parameters. These diagrams offer the opportunity of quick, by-eye, data to model comparison.  Our paper is organized as follows: In section II we discuss the phenomenological representation of the power spectrum, the analysis method we used and the MCMC algorithm. In section III we present our results. Finally, in section IV, we discuss our conclusions.  \\medskip ",
        "conclusions": "\\medskip  In this {\\it paper} we investigated the consistency of the most recent CMB data with a class of $\\Lambda$-CDM adiabatic inflationary models. First we characterized the positions, amplitudes and widths of the peaks by fitting the data with simple phenomenological functions composed by several gaussians. The detection of the peak amplitudes and positions is quite robust and stable between different data sets. We found that all the features are consistent with those expected by the standard theory. We also examined where the data contains the most information in the power-angular scale plane. We found that the low frequency experiments provide good constraints at small angular scales, consistent with the expected damping tail, whereas high frequency experiments provide strong limits on the power at large and intermediate scales. We observe that HF experiments and LF experiments yield very consistent results, although LF data seems to provide evidence for higher secondary oscillations. Overall, the power spectrum is now well determined until $\\ell \\sim 1500$. The inclusion of older data does not affect our conclusions as they do not measure the power beyond $\\ell \\sim 400$.  Furthermore, we related the features in the spectrum with several cosmological parameters by introducing cosmological diagrams that can be used for quick, by-eye, parameter estimations.  The relative amplitude of the first and second peak, in particular, of about $\\sim 1.56$ is consistent with the baryon density expected from BBN and suggests a value of $n_S$ lower than one in the case of negligible reionization. The amplitudes of the third peak relative to the first and to the second, $\\Delta T_1/\\Delta T_3\\sim 1.6$ and $\\Delta T_2/\\Delta T_3 \\sim 1$ strongly suggest the presence of cold dark matter but also limits time its contribution to values $\\omega_{cdm} <0.2$. The relative positions of the peaks, $\\ell_2-\\ell_1 \\sim 330$ and $\\ell_3-\\ell_1 \\sim 610$ is pointing towards the presence of a cosmological constant, a Hubble parameter on the low side of the value allowed by the recent HST measurements ($h \\sim 0.65$) and to an age of the universe $t_0 \\sim 14.5$ Gyrs consistent with the measurements of the oldest globular clusters.  It is reassuring that all those conclusions, obtained by just drawing few lines in the diagrams presented in Figs. $5-9$, are in agreement with the results obtained by a more careful standard analysis. Within the models considered in our database we found (at $68 \\%$ c.l.): $n_s=0.96\\pm 0.03$, $\\omega_b=0.022\\pm 0.003$, $\\omega_{cdm}=0.12 \\pm 0.03$, $\\Omega_{\\Lambda}=0.63\\pm0.16$, and $t_0= 14.2 \\pm 0.7$ Gyrs.  The results obtained here show no need for modifications to the standard model, like gravity waves, quintessence, isocurvature modes, or extra-backgrounds of relativistic particles. Furthermore, possible systematic effects due to unknown foregrounds or calibration and beam uncertainties are not immediately suggested, since the different data sets are consistent with the theory.  Even if the width of the gaussians is poorly constrained, we found supporting evidence for multiple oscillations in the data between $430 < \\ell < 910$. Beyond that, the newest experimental results show a damping of the power.  It is the duty of future satellite CMB experiments to point out discrepancies that might place the possibility of new physics in a more favorable light.  \\medskip  \\textit{Acknowledgements} It is a pleasure to thank Ruth Durrer, Anthony Lewis, Ruediger Kneissl, Roya Mohayaee, Lyman Page, Joseph Silk and Anze Slosar for useful comments. We acknowledge the use of CMBFAST~\\cite{sz}. CJO is supported by the Leenaards Foundation, the Acube Fund, an Isaac Newton Studentship and a Girton College Scholarship. AM is supported by PPARC."
    },
    "0207/astro-ph0207079_arXiv.txt": {
        "abstract": "In a paper by Sanwal et al. (2002), it is supposed to be very difficult to interpret the absorption features in term of cyclotron lines. However, we would like to address here that the possibility of the absorption being cyclotron resonance can not be ruled out. We propose that the isolate neutron star, 1E 1207.4-5209 in the center of supernova remnant PKS 1209-51/52, has a debris disk and is in a propeller phase, with an accretion rate $\\sim 6\\times 10^{-11}M_\\odot$/year. In this scenario, 1E 1207.4-5209 could also be a bare strange star.  \\vspace{0.4cm} % \\noindent % {\\em PACS:} 97.60.Gb, 97.60.Jd, 97.60.Sm ",
        "introduction": "Strange (quark) star is composed of nearly equal number of up, down, and strange quarks, and a few electrons for keeping neutralization of matter. It has important implications for studying the phase diagram of strong interaction system whether this kind of quark stars exist. Recently, Xu (2002) suggests that a featureless thermal spectrum could be a probe for identifying strange stars, since no bound charged particle is in discrete quantum states on the quark surface without strong magnetic field. Nonetheless, it is worth noting that discrete Landau levels appear for charged particles in strong fields.  Two absorption lines, at $\\sim 0.7$ and $\\sim 1.4$ keV, are detected from an isolate neutron star (1E 1207.4-5209) with {\\it Chandra} by Sanwal et al. (2002), and are then confirmed with {\\it XMM-Newton} by Mereghetti et al. (2002). Certainly 1E 1207.4-5209 can not be a bare strange star if those two lines are atomic-transition originated, although the stellar mass $M$ and radius $R$ may be derived by obtaining the gravitational redshift (as $M/R$) and the pressure broadening (as $M/R^2$) of the lines. However, if these double lines are caused by the Landau-level transition of electrons, 1E 1207.4-5209 could also be a bare strange star since no atom might be on the stellar surface. Sanwal et al. (2002) addressed that the features are associated with atomic transition of once-ionized helium, and thought that it is hard to interpret the absorption features in term of cyclotron lines. However, we will find in the next section that the possibility of the absorption being cyclotron resonance can not be ruled out. We will present, in this {\\em Letter}, a short note to confute the interpretation of the recently discovered lines in the X-ray spectrum of 1E1207-52, that these spectral lines are testifying the presence of an atmosphere on the star which can absolutely not be a bare strange star. ",
        "conclusions": "Certainly Sanwal et al's discovery is very important in both possibility: the mass and radius may be derived if the absorption are atomic transition originated, or the accretion rate in the propeller phase could be estimated, for the first time, in case of two cyclotron lines (point 1).  Mereghetti et al. (2002) found that the absorption features are phase-dependent: the $\\sim 1.4$ keV line prefers to appear during the minimum and the rising the parts, rather than at the peak, of the pulse profile. This observational property may reflect the geometry of resonant cyclotron emission: there is almost only one fundamental line for an observer along the magnetic fields, while more harmonic lines appear if the line-of-sight is perpendicular to the fields (Fig.3 of Freeman et al. 1999). An effort to fit the observed spectrum of 1E 1207.4-5209 was tried by Hailey \\& Mori (2002) who presumed that the star has an atmosphere with He-like Oxygen or Neon in not too high a field; whereas an elaborate model calculation (being prepared), to fit in term of cyclotron resonance lines, is also necessary in order to know the details of the source.  A very interesting and important question is: why is 1E 1207.4-5209 the only one in which significant absorption features have been detected so far? To answer this question, Mereghetti et al. (2002) suggested 1E 1207.4-5209 has a metal atmosphere, which is not old enough to accrete a hydrogen layer. However, this question may naturally be answered by the selective effect in observations, since maybe only a few sources have magnetic fields being suitable for creating cyclotron lines with energies in the detector energy range. The fundamental electron cyclotron resonance lies at $\\Delta E = 11.6 B_{12}\\sqrt{1-R_{\\rm s}/R}$ keV, where $B_{12}$ is the polar magnetic field in $10^{12}$ G, $R_{\\rm s}\\equiv 2GM/c^2$ is the Schwarzschild radius, and $M$ and $R$ are the stellar mass and radius, respectively. For a bare strange star with certain mass $M$, one can obtain its radius $R$ by integrating numerically the TOV equation, with the inclusion of the equation of state for strange matter: $P=(\\rho-4B)/3$ ($B$ is the bag constant). The fundamental resonance energy as a function of magnetic field for strange stars with different masses is shown in Fig.1. We can see that, for detectors ({\\it Chandra} or {\\it XMM-Newton}) from $\\sim 0.1$ to $\\sim 10$ keV, the sensitivity fields in which electrons can absorb resonantly photons within that energy range are from $9\\times 10^9$ G to $1\\times 10^{12}$G. It is well known that pulsars tend to have a magnetic field of $\\sim 10^{12}$ G (normal pulsars) or of $\\sim 10^8$ G (millisecond pulsars); it is thus not a surprise that only few sources are observed to show spectral lines. {\\em No source} listed in the table of Xu (2002) has definitely a suitable field.  Recently, a 5 keV absorption feature has been detected and confirmed in the bursts of a soft-gamma-repeater SGR 1806-20 (Ibrahim et al. 2002a, 2002b), which is believed to be the feature as one of the proton cyclotron lines in superstrong magnetic field ($\\sim 10^{15}$ G) by the authors. However there are some difficulties in this explanation. 1, due to the high mass-energy ($\\sim 1$ GeV) of a proton, the ratio of the oscillator strength of the first harmonic to that of fundamental in $10^{15}$ G is {\\em only} $\\sim 10^{-6}$! It is not reasonable to detect the first and the {\\em even} higher harmonics. In fact, numerical spectrum simulations of atmospheres with protons in superstrong fields have never show more than two proton absorption lines (Ho \\& Lai 2001, 2002). 2, a better and more reasonable model for the continuum spectrum component is needed in order to identify such absorption features in reality. Motivated by these flaws, we suggest that the possible absorption lines at $\\sim 5$, $\\sim 11.2$, and $\\sim 17.5$ keV could be interpreted as electron cyclotron lines, while the $\\sim 7.5$ keV absorption might be caused by other effects (e.g., can the accreting plasma with irons absorb at $\\sim 7.5$ keV?). The much small ratio, $\\sim 10^{-7}$, of oscillation strength can also not large enough to produce a second harmonic of $\\alpha$ particle at $\\sim 7.5$ keV. SGR 1806-20 may have an ordinary magnetic field, $\\sim 5\\times 10^{11}$ G, which should be another pulsar-like compact stars with suitable magnetic fields for the detectors in the sky.  Strange stars could exist; the exotic surface of a bare strange star might eventually result in the identification of them, especially the most probable one RX J1856 (e.g., Drake et al. 2002, Xu 2002). Although each of the observed phenomena from pulsar-like stars may be interpreted under the regime of traditional neutron star with unusual or artificial physical properties, it might be a natural way to understand the observations by updating ``neutron'' stars with (bare) strange stars.  {\\it Acknowledgments.}~~ We would like to thank Dr. Bing Zhang for his valuable suggestions. RXX wishes to thank Dr. Jianrong Shi for his valuable discussions about cyclotron line formation.  \\vspace{1cm}"
    },
    "0207/astro-ph0207553_arXiv.txt": {
        "abstract": "Measurements of optical properties in media enclosing \\v{C}erenkov neutrino telescopes are important not only at the moment of the selection of an adequate site, but also for the continuous characterization of the medium as a function of time. Over the two last decades, the Baikal collaboration has been measuring the optical properties of the deep water in Lake Baikal (Siberia) where, since April 1998, the neutrino telescope NT-200 is in operation. Measurements have been made with custom devices. The NEMO Collaboration, aiming at the construction of a km$^3$ \\v{C}erenkov neutrino detector in the Mediterranean Sea, has developed an experimental setup for the measurement of oceanographic and optical properties of deep sea water. This setup is based on a commercial transmissometer. During a joint campaign of the two collaborations in March and April 2001, light absorption, scattering and attenuation in water have been measured. The results are compatible with previous ones reported by the Baikal Collaboration and show convincing agreement between the two experimental techniques. ",
        "introduction": "After a long period of experimental work, large \\v{C}erenkov detectors for high energy neutrinos are going to open a new observational window to the sky. Their main goal is to extend the volume of the explored Universe by neutrinos, to obtain a complementary view of astronomical objects and to learn about the origin of high energy cosmic rays. They are the successors of underground neutrino detectors which have turned out to be too small to detect the faint fluxes of neutrinos from cosmic accelerators.  The new detectors are large, expandable arrays of photomultipliers constructed in open water or ice. The photomultipliers span a three-dimensional coarse grid and map the \\v{C}erenkov light of secondary particles produced in  neutrino interactions. Actually, the basic idea for this detection method goes back to the early 60's \\cite{Markov1961}. Pioneering attempts towards its realization have been made in the course of the DUMAND project \\cite{dumand}. In 1993, the Baikal Collaboration \\cite{Astroparticle-97} succeeded to built the first deep underwater \\v{C}erenkov neutrino detector, which has been stepwise upgraded to its present stage, NT-200. The AMANDA Collaboration \\cite{Andres2000} has built a \\v{C}erenkov detector in the South Pole ice. Other collaborations (ANTARES \\cite{ANTARES}, NESTOR \\cite{Resvanis1993}) are constructing underwater neutrino detectors of similar size. Since a few years, the NEMO Collaboration \\cite{Capone1999} is performing an intensive R\\&D program  aiming at the construction of a km$^3$ \\v{C}erenkov neutrino telescope in the Mediterranean Sea. Another cubic kilometer detector, IceCube \\cite{Spiering} is planned at the South Pole. The cubic kilometer scale is set by various predictions on the extremely low fluxes of high energy neutrinos expected from astrophysical sources.  In underwater \\v{C}erenkov neutrino telescopes, water acts not only as a target but also as radiator of \\v{C}erenkov photons produced by relativistic charged particles. The detection volume, as well as the angular and energy resolutions strongly depend on the water transparency.  The transparency of water as a function of photon wavelength $\\lambda$, is described by the so called inherent optical properties, like the coefficients for absorption $a(\\lambda)$, for scattering $b(\\lambda)$, for attenuation $c(\\lambda) = a(\\lambda) + b(\\lambda)$, and by the phase scattering function $\\beta(\\lambda , \\vartheta)$ (also referred to as volume scattering function) which represents, for a photon, the probability to be diffused at an angle $\\vartheta$ \\cite{Mobley1994}. Another parameter commonly used in literature is the effective scattering coefficient $b^{eff}(\\lambda) = b(\\lambda)(1-\\overline{cos(\\lambda,\\vartheta)})$, where $\\overline{cos(\\lambda,\\vartheta)} = \\int _{0}^{\\pi} cos(\\vartheta) \\beta(\\lambda,\\vartheta) d\\vartheta / \\int_ {0}^{\\pi} \\beta(\\lambda,\\vartheta) d\\vartheta$ is the average cosine of the phase scattering function at a given $\\lambda$. The optical properties of natural water have to be measured {\\it in-situ} in order to allow an unbiased knowledge of light transmission properties in the medium.  The Baikal collaboration has been investigating the fresh water deep in Lake Baikal since 1980. The inherent optical properties have been measured with a series of specially designed devices. It was shown that the water transparency at depths between 900 m and 1200 m is adequate to operate a neutrino telescope. Put into operation at April 6$^{th}$, 1998, the neutrino telescope NT-200 incorporates a long-term monitoring system which performs continuous measurements of the water parameters. This information serves as input for Monte-Carlo simulations of the detector response to atmospheric muons which represent a well-known calibration source for neutrino telescopes. The muon fluxes measured with NT-200 are in very good agreement with simulation results. This fact confirms that the custom-made devices and the methods to extract the relevant information on optical parameters yield reliable results.  The NEMO collaboration has been investigating oceanographic and optical properties of several deep sea marine sites close to the Italian coast, with the aim to select the optimal site for the construction of a km$^3$ detector in the Mediterranean Sea. Absorption and attenuation coefficients for light in the wavelength region between 412$\\div$715 nm \\cite{Capone2001} have been measured with a set-up based on commercial devices.  Optical measurements in deep water are extremely difficult, and possible systematic errors related to these measurements suggest careful cross checks of results by complementary methods. For these reasons, during March - April 2001, the NEMO and Baikal Collaborations have started a joint campaign to measure the optical properties of deep water in Lake Baikal using two different devices. One set-up is based on the transmissometer {\\it AC9}, operated by the NEMO group, the other device, {\\it ASP-15} (Absorption, Scattering and Phase function meter), was developed and operated by the Baikal Collaboration. The cross check of experimental results has been crucial for both devices, since both have an excellent sensitivity in measuring water optical properties, however, they can be affected by different sources of systematic errors which could deteriorate the absolute accuracy. The measurements reported in the following sections have been carried out during March - April 2001, from the ice camp above the neutrino telescope NT-200. ",
        "conclusions": "Measurements of the optical water properties in Lake Baikal confirm that the NT-200 telescope is located at optimal depth, where light absorption and attenuation processes are the smallest. Data have been collected with two instruments, which use different measurement principles and have different sources of systematic errors. Data show that, at a depth of 1000 m, the highest transparency is observed for $\\lambda=$ 488 nm. The measured values for absorption length $L_a$, scattering length $L_b$ and attenuation length $L_c$ at 1000 m depth are: $L_a(488)=27.9 \\pm 0.7$ m, $L_c(488)=18.3 \\pm 0.3$ m as measured with {\\it AC9} and $L_a(488)=28.3 \\pm 1.0$ m, $L_b(488)=58.8 \\pm 3.5$ m as measured with {\\it ASP-15}. The depth profile of the absorption coefficient measured by {\\it AC9} (see figure \\ref{fig:Tempacbaikal}) shows the effect of biologically active substances and mineral particulate suspended in water. This effect is very conspicuous in the depth range 0 $\\div$ 400 m (above the boundary depth of penetration of solar radiation), and starts to be visible again for depth higher than 1150 m, near the lake bed.\\\\ The obtained results demonstrate that the systematic errors are rather small for both instruments and validate the use of both devices to characterize {\\it in situ} the inherent optical properties of underwater sites."
    },
    "0207/astro-ph0207309_arXiv.txt": {
        "abstract": "We show that the stellar masses implied by our predictions of the wind properties of massive stars are in agreement with masses derived from evolution theory and from direct measurements using spectroscopic binaries, contrary to previous attempts to derive masses from wind theory. ",
        "introduction": "The stellar winds of early-type stars are thought to be driven by radiation pressure on spectral lines. Predictions of wind properties are usually tested by comparing them to \\mdot\\ \\& \\vinf\\ values derived for a set of about 30 well studied Galactic and Magellanic Cloud O and early-B stars (Puls et al. 1996). Although it has proven challenging enough to match predictions with observations, these stars provide only a limited test. To better constrain wind theory one should confront a wider range of predictions with observations. Meaningful new tests include extending the comparisons of \\mdot\\ and \\vinf\\ to: {\\em   i)} extremely luminous Of and WN5h-6h stars; {\\em  ii)} stars at (both sides of) the bistability jump, an abrupt discontinuity in \\vinf\\ found to occur at spectral type B1, and {\\em iii)} Luminous Blue Variables. We have applied our predictions of mass loss to all the above cases with excellent results (de Koter et al. 1997; Vink et al. 1999, 2000, 2002).  As the properties of stellar winds depend also on stellar mass, an alternative test would be to compare masses derived from line-driven wind theory with those of independent methods. This is a relevant issue in view of current problems with masses based on mass-loss rates and terminal wind velocities. ",
        "conclusions": ""
    },
    "0207/astro-ph0207623_arXiv.txt": {
        "abstract": "We assess the impact of the trace element \\neon on the cooling and seismology of a liquid C/O white dwarf (WD).  Due to this elements' neutron excess, it sinks towards the interior as the liquid WD cools. The subsequent gravitational energy released slows the cooling of the WD by 0.25--1.6 Gyrs by the time it has completely crystallized, depending on the WD mass and the adopted sedimentation rate.  The effects will make massive WDs or those in metal rich clusters (such as NGC 6791) appear younger than their true age.  Our diffusion calculations show that the \\neon mass fraction in the crystallized core actually increases outwards.  The stability of this configuration has not yet been determined.  In the liquid state, the settled \\neon enhances the internal buoyancy of the interior and changes the periods of the high radial order $g$-modes by $\\approx$ 1\\%.  Though a small adjustment, this level of change far exceeds the accuracy of the period measurements.  A full assessment and comparison of mode frequencies for specific WDs should help constrain the still uncertain \\neon diffusion coefficient for the liquid interior. ",
        "introduction": "} After \\carb and \\oxy, the most abundant nucleus in a $M<M_\\sun$ white dwarf (WD) interior is \\neon.  The reason for this starts with the physics of the CNO cycle in the hydrogen burning phase of a star's life. The slowest step in the CNO cycle is the proton capture onto \\nitr and almost all of the CNO catalyst nuclei end up as \\nitr at the completion of \\h burning. The \\nitr$(\\alpha\\,,\\gamma)$\\flor$(\\beta^+)$\\oxyet$(\\alpha,\\gamma)$\\neon reaction sequence during the subsequent helium burning then processes all the \\nitr into \\neon.  This results in a \\Ne mass fraction of $X_{22}\\approx Z_{\\mathrm{CNO}}\\approx 0.02$ for recently formed WDs of $M<M_\\sun$ (see, for example, \\citet{um99}) made from stars of initial mass $\\lesssim 6 \\msun$.   The possibility that \\Ne could play a role in the energetics of a cooling WD was first noted by \\citet{isern91}. They discussed the possibility of gravitational energy release if \\Ne phase separates at core crystallization. This was followed by more detailed studies \\citep{xv92, ogata93, seg94, seg96} which differed in their conclusions about the final state once crystallization was complete; ranging from all the \\Ne in the stellar center to most of it having an unchanged profile. \\citet{bild01}  (hereafter BH01) examined another aspect of \\Ne which had been previously overlooked (except for a brief mention in Bravo et al. 1992), namely its ability to sink through the liquid WD interior.  As discussed in BH01, there is an upward pointing electric field of magnitude $eE\\approx 2m_p g$ in the degenerate WD interior, where $g$ is the local gravitational acceleration and $m_p$ is the proton mass.  The net force on \\Ne is therefore $\\vec{F}=- 22m_p g \\rhat + 10eE \\rhat=- 2m_pg \\rhat$, biasing the diffusion of \\neon inward. It is the excess neutrons of the \\Ne nucleus (relative to the predominant $A=2Z$ nuclei, $A$ being the atomic mass number and $Z$ the charge) that make it special in this regard.  BH01 found that the gravitational energy released by sedimenting \\neon is comparable to the thermal content of the WD at an age of a few Gyrs.  This raised the question of whether \\neon could settle fast enough to significantly affect the WD's cooling significantly.  The rate at which \\neon sinks in the liquid C/O WD depends on the \\neon diffusion coefficient, $D$. The conditions there are very non-ideal, as the average Coulomb energy for the ions is comparable to the thermal energy, or $\\Gamma\\equiv (Ze)^2/ a kT>1$, where $a^3=3Am_p/4\\pi \\rho$. As discussed by \\citet{paq86}, there is substantial (factors of many) uncertainty in the diffusion coefficients in these liquid regimes, as the familiar notions of mean free path lose their meaning.  In the absence of a definitive calculation of $D$ for this situation, BH01 proceeded by estimating $D$ for \\neon in a C/O plasma by the self-diffusion coefficient, $D_s$, of the classical one-component plasma (OCP). With this, BH01 then estimated the power released, $L_g$, by \\neon sinking for a fixed \\neon profile. This calculation suggested that \\neon sedimentation might release sufficient energy to impact WD cooling.  We now take the next step in addressing this question by performing a self-consistent evolution of both the \\neon density, $\\rho_{22}$, and the WD core temperature, $T_c$, in WDs composed of a single dominant ion species.  We find that \\neon heating delays the time it takes a WD to cool to a given luminosity.  \\emph{The total increase in cooling age by the time the WD completely crystallizes ranges from 0.25-1.6 Gyr, depending on the value of $D$  and the WD mass}.  We also investigate the seismological impact of the \\neon abundance profile at the time the WD crosses the ZZ Ceti instability strip.  The gradient in \\neon abundance produces a gradient in the electron mean molecular weight, $\\mu_e$, that provides an additional buoyancy and alters the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency, $N$. \\emph{This contribution alters the pulsation periods of high radial order g-modes by more than the measurement errors}.  Thus $\\mu_e$ gradient contributions from sedimenting \\neon cannot be ignored in precision WD pulsation work such as recent contraints on the interior abundance profiles \\citep{brad01} or \\carb$(\\alpha,\\gamma)$\\oxy reaction rate \\citep{met01}.  An unexpected result of our calculations is the interaction between the infalling \\neon and the outward moving crystal/fluid boundary as the WD cools. We assume that sedimentation halts in the crystalline interior, forcing \\neon to accumulate at the crystal/fluid boundary and elevating the abundance there.  This abundance is then frozen as the crystal front moves outward.  The jump in \\neon abundance at a given location depends on both the rate of \\neon infall and the rate at which the crystal front moves outward.  This leads to an \\neon abundance in the crystal regions of the star that \\emph{increases} outward.  Whether or not this profile is subject to an elastic Rayleigh-Taylor instability remains an open question.  In $\\S$  \\ref{sec:methods}, we cover the details of how we model the evolution of the \\neon density, $\\rho_{22}$,  and the WD core temperature, $T_c$. We also discuss the current uncertainties in the OCP self-diffusion coefficient used to calculate our \\neon flow rates. The details of the numeric evolution of $\\rho_{22}$ are in $\\S$\\ref{sec:diffusion}.  In \\S\\ref{sec:cool} we detail the results of our self-consistent calculations. We first discuss the evolution of $\\rho_{22}$, highlighting the influence of $D$, WD mass, and crystallization on the adundance profiles.  We then turn to the thermal evolution of the WD and present our new cooling curves. With the \\neon abundances in hand, we estimate in \\S\\ref{sec:brunt} the impact \\neon can have on WD pulsations, focusing on high radial order g-modes. We close in \\S\\ref{sec:conclusions} with a summary of our results and a discussion of unanswered questions.  The appendix discusses the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency in the deep interior. ",
        "conclusions": "} With two extra neutrons (over and above the $A = 2 Z$ ratio of the background ions), the diffusion of \\neon is biased inward in the liquid interior of WDs. The impact of this sedimentation on WD cooling was first estimated by BH01.  We have extended their work by calculating the mutual evolution of the \\neon density, $\\rho_{22}$, and core temperature, $T_c$, for WD models composed of a single background ion species.  The heating produced by \\neon increases the cooling age of our models on order of 0.2-1.5 Gyr (see Figure \\ref{fig:delays}) depending on the value of $D$.   We have also performed initial estimates of the effects that \\neon abundance gradient can have on the pulsation modes of WDs.  Although more precise work needs to be done, these initial calculations indicate that the corrections to mode periods due to $\\mu_e$ gradients impact the results at the current level of observational precision.  The uncertainty in the actual value of the diffusion coefficient produces a large uncertainty in the possible impact \\neon may have on WD evolution. The significant impact that \\neon sedimentation may have on ages of recently formed WDs provides clear astrophysical motivation for authoritative calculations for the  $D$ of \\neon through a multicomponent plasma.  Hopefully further theoretical work in this direction will eliminate this source of uncertainty.  Observations can constrain the diffusion rate of \\neon in WD interiors.  For example, say the age of a WD is known because of its membership in a cluster.  Its luminosity will then depend on the rate at which \\neon sinks and on the overall amount of \\neon present in the WD (see equation (\\ref{eq:Lg})). Also the amount by which the brightness of the WD will be affected is very dependent on the WDs mass, as can be seen by the differences between the 0.6 and 1.0 $\\msun$ cooling curves in Figure \\ref{fig:Ltcrys}.  Having an independent measure of these four quantities ($M$, $L$, WD age, and \\xne) can allow meaningful tests of our theory.  The main uncertainty in WD cooling, namely fractionation effects, do not come into play until a large portion of the WD has crystallized and do not seem to be a function of \\neon concentration (at least for $\\xne \\lesssim 0.1$; see \\citet{seg96}).  Increases in WD luminosity over that expected from the canonical cooling theory for times prior to crystallization or which show a strong metallicity dependence are most readily interpreted as the effects of \\neon sedimentation.  The question is whether or not \\neon can affect the WD luminosity to a measurable degree in the regimes where fractionation is unimportant.  One place we might look are massive young WDs. For ages between 1--4 Gyrs, the $L$ of such objects is strongly dependent on $D$. But, on the other hand, these WDs begin to crystallize in this time frame. For $D \\sim D_s$, C/O fractionation effects may dominate over \\neon sedimentation, at least if $\\xne \\sim 0.02$ (\\citet{sal00} and see also Figure \\ref{fig:delays}). For larger values of $D$ though, these WDs will be maintained at a high enough temperature that fractionation becomes significantly less important at these times. The other place we might look for the effects of \\neon are in WDs born from metal rich progenitors.  In this case, the effects on \\neon sedimentation even in the lower mass WDs are more readily apparent and may be observable.  The data required for either of the above two observational programs can be obtained for WDs that are members of open clusters.  In this case, the age of the WD can be inferred from the difference between the cluster's turn-off age and the main sequence lifetime of it progenitor, if the WD mass is known. The cluster's metallicity constrains the \\neon content of the WDs. To date, though, the actual quantity of data for confirmed WD cluster members is rather scarce.  A compilation of such objects was performed by von Hippel in 1998 \\citep{vonhip98}.  In this sample (which is still seemingly complete), there is only one WD (G152 in NGC 2682) with a mass or $g$ determination in a cluster older than 1 Gyr \\citep{lands98}.  The other 27 objects all reside in clusters younger than this.  Of these, there are about 5-6 WDs with masses greater than 0.9 $\\msun$ and several (most notably the two WD members of NGC 2168) whose mass determination are highly suspect \\citep{koes88}.  Obviously, it would be ideal to increase the number of known WDs in clusters with a wider range of ages than we have currently. Such a program should be a priority if for no other reason than the need for a direct observational test of WD cooling theory in general.  In the past seven years or so, the number of WD candidates in open clusters have grown substantially due to the series of deep (limiting magnitudes of 24-26 in the $V$ band) photometric studies of open clusters that several groups have undertaken \\citep{vonhip95, rich98, kali2, kali3, vonhip00,vonhip98b,andr02}.  The goal of these studies is to provide a determination of cluster ages using the age of the oldest WD candidates found in each cluster.  Identification of an object as a WD is based on the object's location in the cluster's color-magnitude diagram and on it having a spatial morphology that is stellar in nature (as opposed to one that is galactic---contamination of the WD sequence with faint blue galaxies in these studies is one of the  difficulties encountered)\\citep{vonhip00, kali2}.  To date, no follow up spectroscopy to confirm these objects aras WDs and to make mass determinations if they are has been done, although the CFHT group has future plans to do so \\citep{kali3}.  Overall the number of new WD candidates added by these studies is around 300 objects.  The clusters studied so far range in age from 0.5-7 Gyrs.  It is worth noting that only one of the 19 clusters in the CFHT survey has an age greater than 1 Gyr \\citep{kali1}.  Follow up spectroscopy on clusters outside of the CFHT sample is thus highly desirable.  As mentioned earlier, the cluster NGC 6791 is an extremely interesting target due to its high metal content.  To our knowledge, there have been no WD candidates identified in this cluster and discovery of a WD sequence in NGC 6791 could provide a ready test of our theory since the effects of \\neon on WD cooling will be greatly amplified there.  We would like to thank Peter H\\\"{o}eflich for providing the pre-WD evolutionary models from which we based our discussion of WD composition, Leandro Althaus for providing the $L-T_c$ relations used in our calculations, and James Liebert for a critical reading of our manuscript. This work was supported by the NSF under Grants PHY99-07949, AST01-96422, and AST02-05956.  L. B. is a Cottrell Scholar of the Research Corporation.  \\appendix"
    },
    "0207/astro-ph0207637_arXiv.txt": {
        "abstract": "We construct a Galactic cosmic ray (CR) diffusion model while considering that CR sources reside predominantly in the Galactic spiral arms. We find that the CR flux (CRF) reaching the solar system should periodically increase each crossing of a Galactic spiral arm. We search for this signal in the CR exposure age record of Iron meteorites and confirm this prediction.  We then check the hypothesis that climate, and in particular the temperature, is affected by the CRF to the extent that glaciations can be induced or completely hindered by possible climatic variations. We find that although the geological evidence for the occurrence of IAEs in the past Eon is not unequivocal, it appears to have a nontrivial correlation with the spiral arm crossings---agreeing in period and phase. Thus, a better timing study of glaciations could either confirm this result as an explanation to the occurrence of IAEs or refute a CRF climatic connection. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207547_arXiv.txt": {
        "abstract": "{ We investigate the role of rotational effects in inducing asymmetry present above $\\sim 5$~GeV in the double-peak lightcurves of the bright EGRET pulsars: Vela, Crab, and Geminga. According to Thompson~(\\cite{thompson2001}), the trailing peak dominates over the leading peak above $\\sim 5$~GeV consistently for all three pulsars, even though this is not the case over the entire energy range of EGRET, i.e. above $\\sim 100$~MeV. We present the results of Monte Carlo simulations of electromagnetic cascades in a pulsar magnetosphere within a single-polar-cap scenario with rotationally-induced propagation effects of the order of $\\beta$ (where $\\beta$ is the dimensionless local corotation velocity). We find that even in the case of nearly aligned rotators with spin periods of $P\\sim 0.1$~s rotation may lead to asymmetric (with respect to the magnetic axis) magnetic photon absorption which in turn leads to asymmetric gamma-ray pulse profiles. The resulting features - softer spectrum of the leading peak and the dominance of the trailing peak above $\\sim 5$ GeV - agree qualitatively with the EGRET data of the bright gamma-ray pulsars. ",
        "introduction": "Good quality gamma-ray data for three pulsars - Vela, Crab, and Geminga - acquired with EGRET aboard the CGRO tempts to analyse the properties of pulsar high-energy radiation as a function of photon energy and phase of rotation. Gamma-ray spectra of pulsed radiation from these sources (as well as from three other EGRET pulsars: B1706-44, B1951+32, and B1055-52) extend up to $\\la 10$~GeV. All three pulsars feature gamma lightcurves characterised by two strong peaks separated by 0.4 to 0.5 in rotational phase. These double-peak pulses are asymmetrical and their profiles change with energy. Above $\\sim 100$~MeV the leading peak (LP) is stronger than the trailing peak (TP) in the case of the Vela and the Crab pulsars, and the opposite is true for Geminga. However, for all three pulsars their  leading peaks exhibit lower energy cutoffs - around $\\sim 5$~GeV - than the trailing peaks (TP). In other words, the trailing peaks dominate over the leading peaks above $\\sim 5$~GeV (Thompson \\cite{thompson2001}). In the case of the Vela pulsar and Geminga, this effect is accompanied by the softening of the spectrum of the leading peak (Fierro et al.~\\cite{fmnt}, Kanbach \\cite{kanbach99}). Potential importance of the double-peak pulse asymmetry in the case of Vela was already acknowledged -- at the time when the COS-B data became available -- by Morini~\\cite{morini} who attempted to explain the asymmetry with a hybrid model, with two different mechanisms responsible for the formation of the leading and the trailing peak.    High-energy cutoffs in pulsar spectra are interpreted within polar cap models as due to one-photon absorption of gamma-rays in strong magnetic field with subsequent $e^\\pm$-pair creation. A piece of observational support for such an interpretation comes from a strong correlation between the inferred `spin-down' magnetic field strength and the position of the high-energy cutoff (Baring \\& Harding \\cite{bh2000}, Baring \\cite{b2001}). This in turn opens a possibility that the observed asymmetry between LP and TP, i.e. the dominance of LP over TP above $\\sim 5$~GeV, is a direct consequence of propagation effects (which eventually lead to stronger magnetic photon absorption for photons forming LP than TP) rather than due to some inherent property of the gamma-ray emission region itself.  The aim of this paper is to investigate the role of pulsar rotation in a built-up of such asymmetry in the double-peak pulse profiles. We consider purely rotational effects: due to presence of rotation-induced electric field $\\vec E$, aberration of photon direction and slippage of magnetosphere under the photon's path. In section 2 we compare them with some other effects which may be responsible for the asymmetry (like various distortions of the magnetic field structure). In section 3 we show that the rotation effects result in an asymmetric pair production rate for the leading and the trailing part of the magnetosphere even in the case when the magnetic field structure and the population of radiating particles  are symmetric around the magnetic pole. In section 4 asymmetric pulse profiles are calculated as a function of photon energy and then the model predictions of the ratio of fluxes in the leading and trailing peaks are compared with the inferred ratio for Vela at different energy bins. In section~5 we address the significance of rotation-driven asymmetry across the pulsar parameter space. Our main results are discussed in Section~6. ",
        "conclusions": "We have shown that pulsar rotation induces an asymmetry in the magnetic absorption rate with respect to the magnetic dipole axis. Its consequences are potentially interesting in constraining the phase-space of parameters in the polar cap models of high-energy radiation, provided that very high quality gamma-ray data (e.g. as expected from GLAST) are at hand. Its magnitude depends mainly on the linear velocity $\\beta$ of the magnetosphere at sites of particle acceleration and magnetic photon absorption. When the region of electron acceleration is placed just above the  neutron star surface rotation does not produce any detectable effects even for relatively fast rotating young gamma-ray pulsars. However, it has been argued that at least in the case of the Vela pulsar, such a situation is difficult to reconcile with the spectral high-energy cutoff at about 10~GeV (e.g. Dyks et al.~\\cite{alic}). We find then that raising the accelerator up to $\\sim 4$ neutron star radii (in the spirit of Harding \\& Muslimov \\cite{hm98}) above its polar cap produces asymmetric gamma-ray pulse profiles even in the case of nearly aligned rotators with a spin period of $P\\sim 0.1$~s. The resulting features - softer spectrum of the leading peak and the dominance of the trailing peak above $\\sim 5$ GeV - do agree qualitatively with the EGRET data of the bright gamma-ray pulsars (Thompson~\\cite{thompson2001}).  We are far from concluding that the rotation effects alone can account for the observed asymmetry in the double peaks of the bright EGRET pulsars. On the contrary - some axial asymmetry intrinsic to the region of electron acceleration is inevitable in order to explain the double-peak properties at $\\> 100$~MeV of Geminga and B1706-44, where the leading peak is weaker than the trailing peak. Strong deviations of the actual magnetic field structure from the pure dipole at the stellar surface (eg.~Gil et al.~\\cite{gmm2002}) might be responsible for maintaining axial asymmetry at the site of electron acceleration (unlike the symmetric initial conditions introduced in Section 2). This in turn would lead to electromagnetic cascades whose properties vary with magnetic azimuth. It is important, however, that the propagation effects due to rotation work in the right direction, i.e.~they explain qualitatively the observed weakening of the leading peak with respect to the trailing peak. We emphasize that this weakening occurs only in the vicinity of the (phase-averaged) high-energy spectral cutoff, where the flux level decreases significantly.   Another consequence of the magnetic absorption of high energy photons is a noticeable change in the separation $\\Delta^{\\rm peak}$ between the two peaks in the pulse, taking place near the high-energy spectral cutoff (Dyks \\& Rudak \\cite{dr2000}). In the  model discussed above, with electrons ejected only from a rim of the polar cap, the higher energy of photons requires higher emission altitudes to avoid absorption. Therefore, a slight increase in $\\Delta^{\\rm peak}$ is visible in the three lowermost pulse profiles in Fig.~\\ref{profiles}\\thinspace b. However, if the emission from the interior of the polar cap were included, just the opposite behaviour would occur: $\\Delta^{\\rm peak}$ would decrease near the high-energy cutoff in the spectrum. This is because in this case of a ``filled polar cap tube\", the highest energy non-absorbed photons are emitted closer to the magnetic dipole axis (see Fig.~2 in Dyks \\& Rudak \\cite{dr2000}). The latter case agrees qualitatively with the marginal decrease in peak separation found in the EGRET data for Vela (Kanbach \\cite{kanbach99}).  Stimulated by high-quality observations of gamma-ray pulsars anticipated with GLAST we analysed in Sect.~5 the importance of rotation-driven asymmetry in magnetic absorption for a broad range of pulsar parameters. A decline in gamma-ray flux at high-energy spectral cutoff should inevitably be accompanied by strong changes in pulse profiles: whereas at lower photon energies the profile is determined by the density distribution of primary electrons over the polar cap and the efficiency of photon emission mechanism, in the vicinity of the cutoff it becomes additionally constrained by likely high values of the asymmetry parameter ${\\cal R}_{\\rm esc}(\\theta_{\\rm pc})$ - the situation anticipated for fastly rotating  ($P < 0.01$~s), and highly inclined ($\\alpha \\ga 45^\\circ$) pulsars."
    },
    "0207/astro-ph0207292_arXiv.txt": {
        "abstract": "Long-baseline optical interferometers can now detect and resolve hot dust emission thought to arise at the inner edge of circumstellar disks around young stellar objects (YSOs).  We argue that the near-infrared sizes being measured are closely related to the radius at which dust is sublimated by the stellar radiation field.  We consider how realistic dust optical properties and gas opacity dramatically affect the predicted location of this dust destruction radius, an exercise routinely done in other contexts but so far neglected in the analysis of near-infrared sizes of YSOs.  We also present the accumulated literature of near-infrared YSO sizes in the form of a ``size-luminosity diagram'' and compare with theoretical expectations.  We find evidence that large ($\\simge$1.0\\,$\\mu$m) dust grains predominate in the inner disks of T~Tauri and Herbig Ae/Be stars, under the assumption that the inner-most gaseous disks are optically-thin at visible wavelengths. ",
        "introduction": "Radically improved infrared (IR) detectors \\citep{rmg1999b} have recently allowed optical interferometers to investigate the inner accretion disks around young stellar objects (YSOs) for the first time \\citep{malbet1998,rmg1999,rmgphd, akeson2000,rmg2001,lkha2001,mwc2001,akeson2002}.  In most cases, the near-IR sizes were found to be significantly larger than expected from the favored disk models of the time \\citep{malbet1991,hillenbrand1992,hartmann1993,cg97}.  These modeling efforts typically did not directly consider the near-IR disk sizes; indeed, much of the model sophistication was directed toward explaining the observed long-wavelength ($\\lambda\\simge10\\mu$m) spectral energy distributions.  Theoretical work has recently been focused again on reproducing the near-IR properties of YSOs, both the near-IR spectrum and characteristic sizes \\citep{natta2001,dullemond2001}.  The new models abandon the optically-thick, spatially-thin gas disk in the inner-most region, incorporating an {\\em optically-thin} central cavity to explain the ``large'' near-IR sizes \\citep[as suggested by][]{lkha2001}. Alternatively, some workers argue for a quasi-spherical envelope of dust grains around Herbig Ae/Be stars which also can explain the large characteristic sizes \\citep{mirosh1997,mirosh1999}.  This paper takes a step back from the increasingly complicated models one encounters in the literature in order to develop a minimalist framework useful for interpreting the new interferometric measurements (e.g., as generic, or {\\em model-independent}, as possible).  While the essential physics are gradually being incorporated in sophisticated, self-consistent (``physical'') models, theoretical guidance is needed now in crafting and interpreting disk size surveys beginning at the Keck Interferometer and the Very Large Telescope Interferometer.  Indeed, most of the free parameters in the physical models do not strongly affect the near-IR sizes and only serve to obsfucate the relevant physics in the context of the recent interferometric measurements.  In this Letter, we explore how basic dust grain properties and gas absorption of the stellar ultraviolet continuum can affect the observed characteristic sizes of YSO disks in the near-IR. ",
        "conclusions": "Physical models of accretion disks around young stars are critical for understanding the level of IR excess, the spectral energy denstribution, and the interferometric disk sizes.  However, such models have a large number of parameters and are subject to many uncertain assumptions.  In this Letter we considered only the problem of understanding the observed characteristic near-IR sizes of YSOs. By limiting our focus to just the inner disk, we have motivated a minimalist framework founded on the argument that the characteristic size of a given source is directly related to the radius of dust sublimation.  We investigated quantitatively how the radius of dust sublimation depends on realistic optical constants and gas absorption of the stellar ultraviolet continuum, and we found that the magnitudes of these effects are strong functions of stellar effective temperature and grain size. Our results showed that the observed sizes of YSOs are consistent with the presence of an optically-thin cavity surrounding the star, only if the near-IR emission arises from relatively large dust grains ($a\\simge 1\\mu$m for Herbig Ae/Be stars, somewhat smaller grains are allowed for T~Tauris) heated to the sublimation temperature ($T_s \\sim1500K$). For the hotter stars (Herbig Be), gas absorption/scattering of the stellar ultraviolet continuum is likely to be non-neglible and may further help explain the observed size-luminosity relations.  Despite success at reproducing the typical sizes, we can not easily explain the large scatter in the size-luminosity diagram for the Herbig Ae/Be stars using this minimalist framework. More complete disk surveys are beginning with the Keck Interferometer and the Very Large Telescope Interferometer to explore these systems in more detail. Further, interferometer arrays, such as IOTA, COAST, and CHARA,  image the disk emission from YSOs to determine whether the ring morphology (a central assumption here) is common to these systems or unique to Lkh$\\alpha$~101."
    },
    "0207/astro-ph0207017_arXiv.txt": {
        "abstract": "% We summarize the results of {\\it Beppo}SAX ToO observation of blazars that were known to be in a high state from observations carried out in the optical or X-ray or TeV bands. In some of the sources observed, two spectral components were detected, which are interpreted as synchrotron and inverse Compton emission, respectively. Fast variability was detected in three sources (ON\\,231, BL\\,Lac and S5\\,0716+714), but always only for the synchrotron component. Most of the triggers are from optical observations, consequently most of the sources observed are LBL or intermediate objects. They were in a high state in the X-ray band, but not in an exceptionally high state. No strong shift in the synchrotron peak frequency are reported. This is in line with the findings that the synchrotron peak frequency is more variable for HBL objects, i.e. sources that have this peak at higher energies. ",
        "introduction": "Determining the continuum production mechanism is critical for understanding the central engine in AGN, a fundamental goal in extragalactic astrophysics. The continuum emission of the Blazar class of Active Galactic Nuclei (AGN) is dominated by non-thermal radiation from the radio to the X ray, up to the MeV, GeV and in some cases TeV energy bands. This emission is often rapidly variable at all frequencies and, in general, it has been observed that the amplitude of flux density variations increases and the time scales decrease as a function of frequency from the radio to the X-ray. A natural explanation is that Blazars are dominated by relativistic jets at small angles to the line of sight (Blandford \\& Rees 1978; Urry \\& Padovani 1995).  The variability behaviour of a blazar in a given band depends also from its Spectral Energy Distribution (SED). It is well known that the Blazar SED is double-peaked (in a $\\nu \\ vs \\ \\nu f_{\\nu}$ representation). This is interpreted as due to non thermal synchrotron self-compton (SSC) emission, with the first component due to synchrotron radiation and peaking at IR to X-ray frequencies and the second one due to inverse Compton scattering and peaking in the GeV to TeV band (e.g. Fossati et al. 1998). The location of the synchrotron peak is used to define different classes of Blazars: HBL (High frequency peak Blazars) and LBL (Low frequency peak Blazars) (Giommi \\& Padovani 1994). In the most studied bands, i.e. radio, optical and X-ray, one expects that different sources have different variability behaviour, depending from where the synchrotron peak is located. Normally one expect that the variability is more enhanced after the synchrotron peak, towards the end part of the synchrotron emission, where the cooling time of the electron is shorter. Longer time scale are expected, and observed, in the radio and far-IR bands. Correspondingly, in the X-ray band we expect fast variability for sources that hare dominated by the synchrotron emission (HBL), while for the sources whose X-ray emission is due to the inverse Compton mechanism we do not expect frequent and rapid variability. Actually, the same source can have the synchrotron peak located at different frequencies, e.g. in  the presence of flare like events, as the one seen in Mkn\\,501. This can be explained as due to the injection of fresh electron in the jet (Pian et al. 1998).  Since Blazars emit over the entire electromagnetic spectrum, a key for understanding blazar variability is the acquisition of several wide band spectra in different luminosity states during major flaring episodes. Coupling spectral and temporal information greatly constrains the jet physics, since different models predict different variability as a function of wavelength. Before the {\\it Beppo}SAX advent, important progress in this respect has been achieved for some of the brightest and most studied blazars, as PKS2155-304 (Urry et al. 1997), BL\\,Lac (Bloom et al. 1997), 3C\\,279 (Wehrle et al. 1998). However, thanks to its good energy resolution and sensitivity over an unprecedented large X-ray energy band, from 0.01 up to 200 keV, {\\it Beppo}SAX immediately provided unique results in the multiwavelength study of Blazar (e.g. PKS\\,2155-304 Giommi et al. 1998, Chiappetti et al. 1999; Mkn\\,501 Pian et al. 1998). Having in mind all these results we successfully used the {\\it Beppo}SAX satellite to perform Target of Opportunity (ToO) observations of blazars, that were known to be in a high state from observations carried out either in the optical or X-ray or TeV bands.   \\begin{figure}[h,b] \\begin{tabular}{cc}  \\epsfysize=8cm % \\hspace{-1.5cm} \\epsfbox{tagliaferri_R_1fig1a.ps}  &  \\epsfysize=8cm \\hspace{-1.cm} % \\epsfbox{tagliaferri_R_1fig1b.ps} \\end{tabular} \\caption[h]{{\\bf Left:} The SED of BL Lac during the flare in the summer 1997 (filled symbols, see Bloom et al. 1997) compared to its `quiescent' level, constructed collecting non--simultaneous data found in literature. \\\\ {\\bf Right:} Two {\\it Beppo}SAX observations of Mkn\\,501 while it was in a high state of activity, compared to previous observations. Note the extremely large shift of the synchrotron peak toward higher energies.}  \\end{figure}  This ToO program was motivated in particular by two spectacular cases. The first is the 1997 multiwavelength flare of BL Lac, that we used as the paradigmatic case for the optical triggering. During this flare a number of ground based telescope as well as satellites (ISO, XTE, ASCA and EGRET) were promptly pointed to BL Lac, triggered by the optical observations (IUAC 6693, 6700) of a brightening of over 1 mag. Data taken within the flaring period are reported in Fig. 1 (left panel, filled symbols), where they are compared to previous data. It is evident the increase of the flux at all wavelengths, especially in the X--ray band and in the $\\gamma$--ray EGRET band, testifying the large increase of the bolometric luminosity. Particularly interesting is the behaviour in the X and $\\gamma$--ray range, where also large spectral variations are evident. This challenges any model: for instance, in the synchrotron self--Compton scenario, an increasing number of emitting electrons leads to a linear increase of the optical (synchrotron) flux and to a quadratic increase of the X and $\\gamma$--ray (Compton) flux, as observed, but this model does not simply account for the flattening of the $\\gamma$--ray spectrum and the corresponding shift of the peak of the Compton component. The second case was provided by the {\\it Beppo}SAX observations of Mkn\\,501. Quiescent for all 1996, as witnessed by the All Sky Monitor onboard RossiXTE, at the beginning of 1997 Mkn 501 entered in an extremely high activity phase. Continuous flaring activity was detected in the TeV band, with flux levels reaching 4--8 times the level of the Crab. BeppoSAX observations were scheduled during one of these flares, leading to the discovery of an unprecedented X--ray emission for this object (Pian et al. 1998), with a synchrotron spectrum peaking at or above 100 keV. Compared to previous observations, the peak shifted by more than two decades in frequency (see Fig. 1). This was used as the best case for a X-ray or TeV trigger.  As part of our ToO program we observed 7 different blazars, some of them more than one time, over a period of 3.5 years. The journal of these observations are given in Table 1, where we report the source name, the observation date, the exposure time and the trigger criteria that started the observation (optical or X-ray, unfortunately we did not have a TeV trigger). We also report other two ToO observation of Blazars that were carried out by {\\it Beppo}SAX, but that were not part of our program. They are Mkn\\,421 (Malizia et al. 2000) and OJ\\,287 (Massaro et al. 2002). We will now give the results of some of these observations in more details.    \\begin{table}[h,t] \\caption{Journal of the {\\it Beppo}SAX Blazars ToO observations} \\begin{center} \\begin{tabular}{|l|c|c|l|} \\hline  Source Name   & Observ. Date & Exposure & Trigger \\\\ \\hline  ON\\,231       & 11 May 1998  & 25\\,ks   & optical \\\\ & 11 Jun 1998  & 32\\,ks   &  \\\\ PKS\\,2005-489 & 01 Nov 1998  & 52\\,ks   & X-ray \\\\ BL\\,Lac       & 05 Jun 1999  & 54\\,ks   & optical+X-ray \\\\ & 05 Dec 1999  & 54\\,ks   &  \\\\ OQ\\,530       & 03 Mar 2000  & 26\\,ks   & optical \\\\ & 26 Mar 2000  & 23\\,ks   &  \\\\ S5\\,0716+714  & 30 Oct 2000  & 43\\,ks   & optical \\\\ MS\\,14588+2249 & 19 Feb 2001 & 48\\,ks   & optical \\\\ 1ES\\,1959+65  & 23 Sep 2001  & ~7\\,ks   & optical \\\\ & 28 Sep 2001  & 48\\,ks   &  \\\\ \\hline Mkn\\,421      & 22 Jun 1998  & 32\\,ks   & X-ray \\\\ OJ\\,287       & 20 Nov 2001  & 40\\,ks   & Optical \\\\  \\hline  \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We presented some of the most important results that were obtained with {\\it Beppo}SAX ToO observations of Blazars. They immediately show the importance of observing in the X-ray band Blazars that are known to be in a high state. With {\\it Beppo}SAX this is even more true: thanks to its large energy range, it has been possible to simultaneously detect both the synchrotron and the Compton components. We detected fast variability events that allowed us to put constraints on the size of the emitting region and to infer the properties of the jets responsible for the Blazars' emission.  But we had also limits in this projects. For instance, it can be seen from Table 1 that the Blazars ToO program has been dominated by optical triggers. Thus, we are probably biased towards sources that have an higher optical variability. These should be the blazars that have the synchrotron peak in the IR-optical band, i.e. sources with the peak on the left side of the bands in which the blazars are usually monitored. These are of course essentially LBL or intermediate blazars. Sources that have the synchrotron peak in the UV--X-ray band are not expected to show strong optical variability. They are probably more easily detected in a high state from systematic monitoring at higher frequencies. Of course this is much more difficult than in the optical and this explain while in Table 1 there are only two X-ray triggers. This could explain the fact that most of the sources observed as ToO observations were in a high state in the X-ray band, but not in an extremely high state. Moreover, we note that from the optical trigger to the actual X-ray observations normally there are delays of a few days and this could also explain while we did not find sources in exceptionally high states.   \\begin{figure}[h,t] \\vspace{-2.0cm}  \\epsfysize=9cm  \\epsfbox{tagliaferri_R_1fig5.ps}  \\caption[h]{Peak frequency vs. luminosity at the peak frequency for a sample of Blazars. Note how the high peak objects seem to show higher variability of the synchrotron peak frequency, while lower peaked sources show a rather steady $\\nu_{peak}$ among different luminosity states (from Costamante et al. 2000). }  \\end{figure}  These observations shows also that the synchrotron peak frequency does not seem to vary a lot in the LBL or intermediate objects. On the contrary, strong shifts have been observed for Mkn\\,501 and Mkn\\,421 that are HBL. This is in line with the findings of Costamante et al (2001), that studied with {\\it Beppo}SAX blazars with extreme synchrotron peak frequencies ($\\nu_{peak} > 1$ keV). They found that these sources seems to be characterised by larger $\\nu_{peak}$ variability, compared with lower $\\nu_{peak}$ objects (see Fig. 5).   We detected strong spectral variability, founding in the same source either two or only one component. We also found fast variability, but only in the synchrotron component of sources showing both the synchrotron and the Compton components. Fast variability is present also in the optical band, but it is less pronounced and there is no one to one correspondance. All this can be intepreted with the presence of a steady Compton component, and the erratic variability of the synchrotron tail emission, coming in and out of the soft X-ray band. The Compton emission we see in the X-ray band is well below the Compton peak and it is produced by low energy electrons scattering low frequency synchrotron photons. The variability seen in the synchrotron part can be obtained by changing the slope of the injected electron distribution, without affecting the total injected power. Time to time there are variability also in the Compton component and this imply a strong modification of the overall blazar SED, as in the case of the 1997 BL\\,Lac flare. All these behaviours can be reproduced by the presence of relativistic jets dominated by shock events produced by colliding shells (e.g. Spada et al. 2001).  In any case these observations show the importance of observing Blazars over such a large X-ary energy band while they are in a high state. It will be crucial to perform these ToO observations also in the future and in particular for Blazars that are detected in a high state in the X-ray or TeV bands, now that new and more sensitive TeV telescope will be operational. In the forseable future these ToO observations, as the one carried out with {\\it Beppo}SAX, will be possible probably either with the combination of simultaneous observations from Integral and other soft X-ray satellites (such as XMM-Newton or Chandra), or with the Swift satellite. Swift will be launched at the end of 2003. Swift is dedicated to the study of Gamma Ray Bursts, but it should also observe other interesting sources, whose emission goes from the soft to the hard X rays, and Blazars are obvious good candidates."
    },
    "0207/astro-ph0207221_arXiv.txt": {
        "abstract": "We present an analysis of optical ($B-R$) and optical-infrared ($R-K_s$) color maps for 47 extremely late-type, edge-on, unwarped, bulgeless disk galaxies spanning a wide range of mass.  The color maps show that the thin disks of these galaxies are embedded within a low surface brightness red envelope.  This component is substantially thicker than the thin disk ($a/b\\,\\sim\\,$4:1, vs $>\\,$8:1), extends to at least 5 vertical disk scale heights above the galaxy midplane, and has a radial scale length that appears to be uncorrelated with that of the embedded thin disk.  The color of the red envelope is similar from galaxy to galaxy, even when the thin disk is extremely blue, and is consistent with a relatively old ($>\\!6\\Gyr$) stellar population that is not particularly metal-poor. The color difference between the embedded thin disk and the red stellar envelope varies systematically with rotation speed, reflecting an increasing age difference between the thin and thick components in lower mass galaxies, driven primarily by changes in the age of the thin disk.  The red stellar envelopes are similar to the thick disk of the Milky Way, having common surface brightnesses, spatial distributions, mean ages, and metallicities.  We argue that the ubiquity of the red stellar envelopes implies that the formation of the thick disk is a nearly universal feature of disk formation and is not necessarily connected to the formation of a bulge.  Our data suggest that the thick disk forms early ($>6\\Gyr$ ago), even within galaxies where the bulk of the stars formed very recently ($<2\\Gyr$). We argue that several aspects of our data and the observed properties of the Milky Way thick disk argue in favor of a merger origin for the thick disk population.  If so, then the age of the thick disk marks the end of the epoch of major merging, and the age difference between the younger thin disk and the older thick disk can become a strong constraint on cosmological constants and models of galaxy and/or structure formation. ",
        "introduction": "Currently, most observational studies of galaxy formation focus on two epochs -- extremely high redshift, where one can observe galaxy formation in progress, and zero redshift, where one can disentangle past history using individual stars within the Milky Way.  While each of these approaches has been essential in shaping our current view of galaxy formation, they each have fundamental limitations.  At high redshifts, it is extremely difficult to match galaxies to their low redshift descendents due to morphological transformation, luminosity evolution, and merging; only changes in the mean galaxy population can be tracked, revealing few details of the physical mechanisms which drive evolution.  In contrast, within the Milky Way a wealth of detail can be extracted from the ages, metallicities, and kinematics of stars, allowing us to trace the formation of the faintest individual components of the Galaxy (the stellar halo, the thick disk, tidal streams, etc).  However, in the end these data address only the formation of the Milky Way, and give no constraints on how galaxy formation proceeds in the mean population, or varies with fundamental parameters (e.g.\\ mass, angular momentum, local density, etc.).   An opportunity to bridge these regimes lies in the realm of nearby galaxies, just beyond the confines of the Local Group.  At moderate distances ($cz\\lesssim5000\\kms$), galaxies of all types are plentiful and are extremely well resolved spatially ($\\Delta\\theta\\lesssim200h^{-1}\\pc$ from the ground, or $\\Delta\\theta\\lesssim20h^{-1}\\pc$ from space), allowing us to trace their morphology and internal dynamics on small spatial scales.  These features can be studied at very high signal-to-noise and/or at low surface brightnesses inaccessible at higher redshifts. To place observational constraints on the process of disk galaxy formation using nearby galaxies, we are engaged in a comprehensive program to study the dynamics, gas content, metallicity, and stellar populations of a population of late-type, bulgeless disk galaxies. This population forms a structurally uniform sample, allowing us to isolate changes in the physical properties of the galaxies (i.e.\\ mass, angular momentum, etc.) independent of changes in morphology, akin to what is possible with the nearly single parameter sequence spanned by elliptical galaxies.  By avoiding systems with bulges, we also limit the degree to which the baryonic component of the galaxy may have been affected by dissipation or angular momentum transport during formation.  This yields a sample which represents the purest endpoint of the disk galaxy formation process.  Details of the sample selection and the optical and infrared imaging can be found in Dalcanton \\& Bernstein 2000 (hereafter ``Paper I'').  In this paper we use the imaging presented in Paper I to undertake an analysis of the color maps of the bulgeless, edge-on disks which comprise our sample.  We focus our attention on the vertical color gradients within the galaxies, probing the stellar populations of the galaxies at many scale heights above the thin disk.  Buried within these low surface brightness components are the remnants of some of the earliest epochs in the assembly of galaxies, namely the thick disk and stellar halo.  Because the typical metallicity of a galaxy tends to increase with time, it has long been recognized that the low metallicity thick disk and stellar halo are fossil records of the very early history of the Milky Way.  Detailed studies of their kinematics and metal abundance have revealed signatures of the processes which led to their formation over 10 billion years ago, even though these components contain only a small fraction of the total stellar mass of the Milky Way.  In general, these two components are thought to be the leftovers from either the monolithic, dissipative collapse of the early galaxy (Eggen et al.\\ 1962, hereafter ``ELS'') or the buildup of the galaxy through hierarchical merging (Searle \\& Zinn 1978). Because of the low surface brightness of stellar halos and thick disks, it is impossible to study their formation directly at high redshift (due to cosmological $(1+z)^4$ dimming), and we are confined to deducing their history from very low redshift data alone.  Almost all of the detailed knowledge of the formation of thick disks and stellar halos comes from evidence within the Milky Way alone (see van den Bergh 1996 for a review), teaching us little about galaxy formation {\\emph{in general}}.  For this reason, astronomers have attempted to identify these faint components in other very nearby galaxies, particularly in the edge-on orientation where the light from the younger thin disk can be minimized (Burstein 1979, Tsikoudi 1979, van der Kruit \\& Searle 1981).  Previous detections of possible halo or thick disk stellar light in external galaxies have been made in a scant handful of nearby edge-on galaxies (e.g.\\ recently Neeser et al.\\ 2000, Fry et al.\\ 1999, Zheng et al.\\ 1999, Morrison et al.\\ 1994, N\\\"aslund \\& J\\\"ors\\\"ater 1997, Morrison et al.\\ 1997, Sackett et al.\\ 1994, van Dokkum et al.\\ 1994, Shaw \\& Gilmore 1990; see \\S\\ref{otherworksec} below for further discussion).  Typically, the presence of a thick disk or stellar halo has been identified by the need for an additional disk component when attempting to fit models of the light distribution in a deep image.  Not all galaxies have required this second component, however.  Instead, thick disks have only been identified in a handful of relatively massive Sc (or earlier) galaxies with substantial bulges (see summary by Morrison 1999).  One limitation of the previous searches for thick disks is that almost all have been based on imaging in a single passband, discriminating between the thick disk and thin disk components through subtle changes in the surface brightness profile perpendicular to the plane.  It is therefore difficult to make a unique decomposition of the thick and thin disks when neither dominates in the region studied, as noted by Morrison et al.\\ (1997).  In this paper, however, we use multi-color imaging to identify thick disks via the systematic changes in broad band colors produced by the variation in the stellar populations of the thick and thin disks.  As we show below, we find unambiguous evidence for stellar envelopes surrounding the majority of the nearly 50 disks in our sample, across all galaxy masses.  The structure of the paper is as follows.  We begin by briefly summarizing the galaxy sample and imaging data in \\S\\ref{datasec}.  We present color maps in \\S\\ref{colormapsec} and discuss the general, qualitative implications for vertical color gradients, radial color gradients, and the presence of dust in the sample.  We further quantify the results in \\S\\ref{extractionsec} and interpret them based on comparison with stellar population models in \\S\\ref{interpsec}.  We show that the color gradients and color maps argue for the presence of old, red stellar envelopes around most, if not all, disk galaxies.  We analyze the shapes of the stellar envelopes in \\S\\ref{isophotesec}.  Color gradients and isophotes are more difficult to interpret in the more massive galaxies for a variety of reasons which we discuss in \\S\\ref{massivegalaxysec}.  In \\S\\S\\ref{interpretationsec} \\& \\ref{formationsec}, we suggest that the stellar envelopes in this sample are analogous to the thick disk of the Milky Way and have properties consistent with those expected by a stochastic merging scenario for the formation of the thick disk.  To conclude, we discuss the general constraints which can be placed on galaxy formation based on the observation of ubiquitous thick disks (\\S\\ref{conclusionsec}). ",
        "conclusions": "\\label{conclusionsec}  In this paper we have presented several lines of evidence which lead us to conclude that thick disks are a common product of disk galaxy formation for galaxies of all masses.  Specifically, we have analyzed the color maps, vertical color gradients, and faint isophote shapes for a large sample of edge-on bulgeless disks.  Our observations are consistent with the conclusion that that all galaxies in the sample are embedded within somewhat flattened ($\\sim$4:1) red stellar envelopes whose properties vary little from galaxy to galaxy although the galaxies themselves span an enormous range in mass and color.  We have used stellar synthesis models to argue that the stellar envelopes are old (at least $6\\Gyr$, and probably older), but not necessarily metal poor ([Fe/H]$<$-1).  We argue that the properties of the red stellar envelopes are consistent with their being close analogs of the MW thick disk.  We find that the evidence in hand, from our sample and the Milky Way, is consistent with a picture for disk galaxy formation which procedes as follows: (1) a thin stellar disk forms at high redshift ($z\\gtrsim1-2$); (2) partial disruption occurs during a significant merger capable of dramatically heating the thin disk, but not necessarily leading to the formation of a bulge; (3) the {\\emph{majority}} of the galaxy's stars form from gas gradually accreted after the merger and the creation of the thick disk; and (4) no significant merger events follow.  This model will only be consistent with cosmological scenarios where the merging rate peaks early on (i.e.\\ low $\\Omega_m$ models).  The evidence suggests that this basic sequence of events is a generic feature of the history of the majority of galaxies which appear as thin disks today.  If so, then it places a number of strong constraints on galaxy formation models:  \\begin{itemize}  \\item Many successful analytic models of disk formation treat the formation of the disk as a monolithic collapse (i.e.\\ Fall \\& Efstathiou 1980, Dalcanton et al.\\ 1997, van den Bosch 1998, 2002). These models tend to produce realistic disks, even though their entire theoretical basis seems to be in conflict with a hierarchical model for the assembly of galaxies.  However, our data suggests that, for galaxies which are disk dominated today, major hierarchical mass accretion probably ends early and involves only a small fraction of the galaxies' mass.  This suggests that it is probably legitimate to treat the formation of the disk as a monolithic dissipative collapse.  \\item The observation that significant mergers are unlikely to have occured in the last $6-8\\Gyr$ for very late-type disk galaxies can place strong constraints on the input parameters for semi-analytic models of galaxy formation, particularly once the mass-accretion threshold for thick disk formation is better constrained by realistic merging simulations of primordial thin disks. Cosmological models which have merging rates that increase to the present day would be less favored by these observations.  \\item In most semi-analytic models of galaxy formation, it is assumed that bulges form through merging of two galaxies with comparable masses, and that any disk is accreted subsequent to the merger. However, the observations of pervasive thick disks suggest that some of the significant mergers early in a galaxy's history lead to the formation of a thick disk, and do not neccessarily produce a bulge (although they might, possibly depending on the gas mass fraction of the merging progenitors).  Thus, it may be necessary to revise the criteria for how bulges are produced in semi-analytic models.  \\item If merging leads to the formation of both a thick disk and a bulge (e.g.\\ Kauffmann et al.\\ 1993), then the age constraints on thick disks place indirect age limits on bulges that form via mergers. Our data therefore suggests that bulges must form early, lest the epoch of merging also create thick disks that are younger than is observed.  However, bulge formation which takes place via secular processes such as bar instabilities (e.g.\\ Pfenniger 1993) is still permitted at any epoch.  \\item Assuming that the thick disks in our sample are produced by merging, and that they persist down to the mass scale of transitional dwarf spheroidals as we have argued above, then there must have been merging sub-units on even smaller mass scales.  This decreasing mass scale sets an upper limit on a possible smoothing scale for the primordial power spectrum, and limits the masses of possible Warm Dark Matter candidates.  \\item Because the thick disk stars are older than those in the thin disk, the thick disk isolates baryonic material from an earlier epoch. Thus, it may be possible to use the relative dynamics and radial distributions of the thick and thin disks to constrain how the specific angular momentum distribution changes as a function of time. This could potentially resolve subtle discrepancies between the angular momentum distribution of thin disks and theoretical models (van den Bosch 2001).  \\item The apparent universality of thick disks down to very low mass scales suggests that it may be difficult to measure truly ``primordial'' Helium abundances for constraining Big Bang nucleosynthesis.  The existence of the red envelope suggests that even the lowest mass (40-60$\\kms$) galaxies with the youngest, bluest star forming disks experienced an even earlier generation of star formation.  This early star formation would be likely to pollute the gas which is currently in the disk, and thus any metallicity measurements made from HII regions would have been enriched not just by the current generation of stars, but a previous one as well.  \\end{itemize}"
    },
    "0207/astro-ph0207198_arXiv.txt": {
        "abstract": "{Absolute and differential chemical abundances are presented for the largest group of massive stars in M31 studied to date. These results were derived from intermediate resolution spectra of seven B-type supergiants, lying within four OB associations covering a galactocentric distance of 5$-$12\\, kpc. The results are mainly based on an LTE analysis, and we additionally present a full non-LTE, unified model atmosphere analysis of one star (OB78-277) to demonstrate the reliability of the differential LTE technique. A comparison of the stellar oxygen abundance with that of previous nebular results shows that there is an offset of between $\\sim 0.15 - 0.4$\\,dex between the two methods which is critically dependent on the empirical calibration adopted for the $R_{23}$ parameter with [O/H]. However within the typical errors of the stellar and nebular analyses (and given the strength of dependence of the nebular results on the calibration used) the oxygen abundances determined in each method are fairly consistent. We determine the radial oxygen abundance gradient from these stars, and do not detect any systematic gradient across this galactocentric range. We find that the inner regions of M31 are not, as previously thought, very 'metal rich'. Our abundances of C, N, O, Mg, Si, Al, S and Fe in the M31 supergiants are very similar to those of massive stars in the solar neighbourhood.} ",
        "introduction": "\\label{intro}  \\begin{table*} \\caption[]{Observational details for the M31 targets. The OB-association numbers are from van den Bergh (\\cite{vnB64}), while the stellar identifications are from Massey et al. (\\cite{Mas86}). For example OB8-17 is star number 17 in association 8. The alternative identifications are from Berkhuijsen et al. (\\cite{Ber88}). Visual magnitudes and spectral types are taken from Massey et al. (\\cite{Mas95}) -- the latter are identified by the authors initials, MAPPW. Also listed are the spectral types and heliocentric radial velocities, $v_r$ (in km s$^{-1}$), estimated from our spectra as discussed in Sect. \\ref{obsdata} and \\ref{sptype}. The data for OB10-64 have been previously discussed by Smartt et al. (\\cite{Sma01b}).} \\begin{flushleft} \\centering \\begin{tabular}{llcllccc} \\hline \\hline  Star      &  Alternative  &   V    &   \\multicolumn{2}{c}{Spectral Type} & Time    & s/n & $v_r$ \\\\ Name      &      ID       &        &  MAPPW     & Here         \t& (hours) & ratio & \\\\ \\hline \\\\ \\\\ OB8-17    &  41-2178      & 18.01  &  O9-B1I    & B1Ia     \t& 4.0\t& 50 \t& $-102\\pm10$ \\\\ OB8-76    &               & 18.52  &  B0III     & B0.5Ia\t& 4.0  \t& 30\t& $-34\\pm7$ \\\\ OB10-64   &  41-2265      & 18.10  &  B1I       & B0Ia          & 4.5   & 50    & $-113\\pm11$ \\\\ OB48-234  &               & 18.50  &  B1I       & B1.5Ia \t& 2.5   & 30 \t& $-125\\pm8$ \\\\ OB48-358  &               & 18.70  &  B0-1I     & B0Ia \t        & 2.5   & 30 \t& $-107\\pm22$ \\\\ OB78-159  &  40-1876      & 17.97  &  B0I\t& B0Ia  \t& 7.0   & 40\t& $-467\\pm20$ \\\\ OB78-277  &  40-1939      & 17.35  &  B1I\t& B1.5Ia   \t& 7.0   & 65\t& $-567\\pm9$ \\\\ OB78-478  &  40-2028      & 17.50  &  B0-1I     & B1.5Ia      \t& 5.0   & 55\t& $-561\\pm12$ \\\\ \\hline \\end{tabular} \\end{flushleft} \\label{obstable} \\end{table*}  Until recently studies of the spatial distribution of chemical species in M31 and many other external galaxies mainly involved H {\\sc ii} regions and supernovae remnants (SNR's) (see for example: Dennefeld \\& Kunth \\cite{Den81}; Blair et al. \\cite{Bla82}; Galarza et al. \\cite{Gal99}). In M31, the H {\\sc ii} regions have low excitation, thus the [O {\\sc III}] lines used in estimating the electron temperature are often too weak to be detected. Without a direct method of determining the electron temperature, empirical calibrations must be implemented and in some cases extrapolated to estimate abundances of these H {\\sc ii} regions. Depending on the empirical calibrations adopted significantly different abundance estimates are derived (Pagel et al. \\cite{Pag80}; Mc Gaugh \\cite{McG91}; Zaritsky et al. \\cite{Zar94}; Kobulnicky et al. \\cite{Kob99}, Pilyugin \\cite{Pil00}, \\cite{Pil01a}, \\cite{Pil01b}). This indicates that an independent method of evaluating the abundance gradients of external galaxies is needed. Blue supergiants are amongst the optically brightest stellar objects in spiral and irregular galaxies, and can provide us with such a method.  Previous studies have shown that quantitative spectroscopy of B-type supergiants in Local Group galaxies can be carried out successfully using 4 m telescopes (see Monteverde et al. \\cite{Mon00}, Smartt et al. \\cite{Sma01b}). The added advantage of using blue supergiants  over H {\\sc ii} regions is that their rich metal line optical spectra provide us with a means of studying elements not observed in the emission line spectra of H\\,{\\sc ii} regions. Other than H, He and CNO, some $\\alpha$-processed and Fe-peak elements are observed in the spectrum of blue supergiants. These elements are important as they help to put constraints on nucleosynthesis theories and models of the chemical evolution of galaxies.  Studies of blue supergiants in external galaxies can also help to enhance the use of the Period-Luminosity (PL) relationship of Cepheid variables in two ways. First one can determine {\\em stellar} abundances in the fields where Cepheids are found to constrain the effect of metallicity on the PL relation. Secondly, although it is difficult to determine extinction to Cepheids themselves model atmosphere fits can accurately determine the interstellar extinction towards blue supergiants, and hence determine reddening in stellar fields where Cepheids are found. In addition, it now appears possible that blue supergiants could be used to determine extragalactic distances by a properly calibrated independent technique, using the wind momentum-luminosity relationship (WLR, see Puls et al. \\cite{Pul96}; Kudritzki et al.  \\cite{Kud99}; Kudritzki \\& Puls \\cite{Kud00}). The WLR allows distances to be determined via detailed studies of the radiatively driven winds in O, B and A-type supergiants.  Several of the stars that will be presented here have been observed by HST with both WFPC2 and STIS, in a broader project to use them  as calibrators of the WLR within the Local Group. First steps in this has occurred by  measuring the terminal velocities from the UV STIS spectra in a related paper (Bresolin et al. \\cite{bres2002}). However to carry out the WLR study in full the atmospheric parameters and chemical composition of the early-type supergiants must be reliably estimated,  and this is the focus of this paper.  Previous to this work, the largest homogeneous abundance analysis on a set of supergiants in M31 is that by Venn et al. (\\cite{Ven00} - hereafter VMLPKL). In this study, two A-type supergiants and one F-type supergiant were analysed using high resolution spectra from the Keck 10 m telescope and the HIRES spectrograph. The stellar oxygen abundances in this analysis suggested a shallow or negligible abundance gradient in M31. The nebular H\\,{\\sc ii} regions give  abundance gradients in the range of -0.013 - -0.027 dex kpc$^{-1}$ , depending on the empirical calibration used, however  the VMLPKL results were based around only three stars and the individual abundances are reasonably consistent to H\\,{\\sc ii} region results at similar galactocentric distances. Definite conclusions from three points are not warrented, and our experience should  caution us against the use of restricted datasets in determining stellar abundance gradients e.g. see the discussion of  Smartt \\& Rolleston (\\cite{smr97}) on the abundance gradient in the  Milky Way. Additionally Smartt et al. (\\cite{Sma01b}) have analysed a B0Ia supergiant, OB10-64, in the inner regions of M31 and compared it to the abundances derived in the H\\,{\\sc ii} region surrounding the OB10 association. The absolute nebular abundance is critically dependent on which calibration of the $R_{23}$ parameter is used, however there is good agreement between the stellar and at least one parameterization of the nebular result.  Here we present a spectroscopic analysis for the largest sample of supergiants in M31 considered to date and covering a range in galactocentric distance of 5 - 12 kpc. The elemental abundances of these B-type supergiants are presented and their reliability discussed including the uncertainties incurred by adopting an LTE analysis. We also determine the oxygen abundance gradient in M31 and compare it with that determined from H\\,{\\sc ii} regions and A \\& F-type supergiants (VMLPKL). ",
        "conclusions": "\\label{conc}  We have presented the results of detailed LTE absolute and differential abundance analyses of the largest group of B-type supergiants in M31 studied to date.  Non-LTE calculations have also been carried out to investigate the effect of departures from LTE on our results. It was shown that although analysed using LTE model atmospheres and line formation codes the analysis is dominated by uncertainties stemming from the quality of the data rather than non-LTE effects.  The seven B-type supergiants lie in distinct clusters which cover a galactocentric distance of 5$-$12 kpc and from the derived abundances we estimated the oxygen abundance gradient of M31. Across this fairly restricted range of the disk we do not find any evidence of a significant abundance gradient, which is similar to the result found by VMLPKL for four A-F -type supergiants.  Radial abundance gradients for the $\\alpha$-processed elements, Si \\& Mg, were also determined indicating a negligible silicon abundance gradient and a possible shallow gradient for magnesium.  However we emphasise that we still have a very restricted number of data points with which to probe the abundance gradient in M31. The work of Smartt \\& Rolleston (\\cite{smr97}), in determining abundance gradients in the Milky Way, cautions against using small numbers of stars with restricted Galactocentric radii to draw firm conclusions on existance of abundance gradients. This result is reasonably consistent with the shallow  negative oxygen abundance gradients determined from H\\,{\\sc ii} regions and supernovae remnants (Dennefeld \\& Kunth \\cite{Den81}; Blair et al. \\cite{Bla82}; Galarza et al. \\cite{Gal99}).  It has been shown that dependent on which empirical calibration of the R$_{23}$ parameter with [O/H] that is adopted, different magnitudes of the radial abundance gradient in M31 are obtained. It has also been shown that there is an offset between the abundances obtained from the H\\,{\\sc ii} regions and that of the B-type supergiants ($\\sim$ 0.15-0.4 dex), at the same galactocentric distance, which again depends on the empirical calibration implemented. These results indicate that the calibration of the R$_{23}$ parameter with [O/H] for H\\,{\\sc ii} regions with high metallicity/low excitation are clearly not accurately constrained. The empirical calibration which fits our stellar results best is that of Pilyugin (\\cite{Pil00}, \\cite{Pil01a}, \\cite{Pil01b}). Moreover this calibration gives the shallowest oxygen abundance gradient ($\\sim$ 0.013 dex kpc$^{-1}$). The main difference between this and other calibrations is that it considers the hardness of the ionising radiation and therefore the physical parameters of the H\\,{\\sc ii} regions. As the B-supergiants only probe out to 12\\,kpc from  the centre of M31, it would be highly desirable to  sample the outer regions of the M31 disk with similar stars.  Our results for the innermost supergiants of the sample, OB8-17 \\& OB10-64, indicate that M31 is not metal-rich, as previously thought from the results of H\\,{\\sc ii} regions, but actually suggests that it is of solar metallicity. This is consistent with the results obtained by Smartt et al. (\\cite{Sma01b}) and from the nebular oxygen abundances when implementing the empirical calibration of Pilyugin (2000, 2001a, 2001b). Although the star OB8-76 appears to have quite high abundances, the mean abundance of the OB8 cluster is not significantly above the solar neighbourhood.  Smartt et al. (\\cite{Sma01b}) have shown that detailed wind-analyses can be accurately carried out on B-type supergiants in M31. They also found that the wind momentum-luminosity relation of Kudritzki et al. (\\cite{Kud99}) can be applied to these B-type supergiants. In the future to calibrate the WLR for the metallicity of M31 and further the calibration of WLR in the Local Group, the atmospheric parameters and abundances will be used in conjunction with the terminal velocity (see Bresolin et al. \\cite{bres2002}) and mass-loss rates of the wind. The final aim of this work will be to provide an accurate and independent extragalactic distance scale."
    },
    "0207/astro-ph0207367_arXiv.txt": {
        "abstract": "In order to distinguish between regular and chaotic planetary orbits we apply a new technique called MEGNO in a wide neighbourhood of orbital parameters determined using standard two-body Keplerian fits for HD$\\thinspace$12661, HD$\\thinspace$38529, HD$\\thinspace$37124, HD$\\thinspace$160691 planetary systems. We show that the currently announced orbital parameters place these systems in very different situations from the point of view of dynamical stability. While HD$\\thinspace$38529 and HD$\\thinspace$37124 are located within large stability zones in the phase space around their determined orbits, the preliminary orbits in  HD$\\thinspace$160691 are highly unstable. The orbital parameters of the HD$\\thinspace$12661 planets are located in a border region between stable and unstable dynamical regimes, so while its currently determined orbital parameters produce stable regular orbits, a minor change within the margin of error of just one parameter may result in a chaotic dynamical system. ",
        "introduction": "The recent explosion in the number of newly detected extrasolar planets has brought the total number of such planets to about one hundred; it appears that the exact number changes too rapidly to be quoted. Some of the planets form multiple planetary systems around their parental stars. At the moment the total number of systems with two or more planetary-mass companions around main sequence stars has reached 11. Four planetary systems - HD$\\thinspace$12661, HD$\\thinspace$38529,  HD$\\thinspace$37124 and HD$\\thinspace$160691 - have recently been elevated to the status of multiple, following the discovery of a second planetary companion in each of them. In most of these systems the strong dynamical interaction between planets makes planetary orbital parameters (see Table 1), found using standard two-body Keplerian fits, unreliable (Laughlin \\& Chambers 2001, 2002). There is also a great uncertainty in the determination of planetary masses. All those leave us a substantial available parameter space to be explored in order to exclude the initial conditions which lead to dynamically unstable configurations.  A classical method that allows one to distinguish between regular and chaotic dynamical states is the method of Lyapunov Characteristic Numbers (LCN). Let us note that {\\it chaotic} in the Poincar{\\'e} sense means that the dynamical behavior is not quasi-periodic, and does not necessarily mean that the system will disintegrate during any limited period of time. The estimation of LCN usually requires computations over long evolutionary time, sometimes much longer than the lifetime of the system studied.  In our previous papers (see for example, Go{\\'z}dziewski et al. 2001a, 2001b) we showed that a new method developed by Cincotta \\& Sim\\'o (2000) and called {\\bf MEGNO} (the acronym of {\\it Mean Exponential Growth of Nearby Orbits}), can be successfully applied to the studies of dynamical stability of extrasolar planetary systems. This method is based on the same ideas as LCN but converges about 100 times faster, and is more sensitive. For example its application to the Gliese 876 system (Go{\\'z}dziewski et al. 2001b) clearly identified the exact location of the 2:1 mean motion resonance and its width. MEGNO helped to show (Go{\\'z}dziewski \\& Maciejewski 2001) that while the orbital parameters of the HD 82943 system derived from the Keplerian fit lead to an unstable self-destructing system, there are some small changes in this fit which lead to stable configurations.  In this work we apply the  technique to all four new planetary systems. We are especially interested in a comparative study of the global dynamics of these systems, as they are likely to represent  different types of dynamical behavior. Taking into account the rather preliminary nature of all orbital fits, we hope that this paper will provide a useful guide for available stable orbital parameters (such as { $a$},{  $e$} and {$\\omega$}) for more sophisticated self-consistent fits (Laughlin \\& Chambers 2001, 2002; Marcy et al. 2002). It may also provide tighter constraints on the parameter space available to some planets due to unsufficient number of observations  as in the case of the HD$\\thinspace$37124c planet where any eccentricity $e_c$ between 0.3 and 0.8 fits the observational data within the velocity errors (Butler et al. 2002).  \\begin{deluxetable}{llllllll} \\tablewidth{0pt} \\tablecaption{Orbital parameters of new planetary systems.}  \\tablehead{ \\colhead{Name}           & $M_p\\sin i$($M_J$)      & $M_*$($\\Msun$)  & $a$(AU) & $P$(days)          & $e$ & $\\omega$(deg) & $T_{peri}^*$(JD-2450000)}  \\startdata {\\bf HD$\\thinspace$12661b\\tablenotemark{a}} & 2.30& 1.07& 0.82& 263.3& 0.35 & 292.6 & 9943.7 \\\\ {\\bf HD$\\thinspace$12661c\\tablenotemark{a}} & 1.56& 1.07& 2.56& 1444.5& 0.20 & 147.0 & 9673.9\\\\ \\hline {\\bf HD$\\thinspace$37124b\\tablenotemark{a}} & 0.86&0.91  &0.54& 153.3& 0.10 & 97.0 & 1227\\\\ {\\bf HD$\\thinspace$37124c\\tablenotemark{a}} & 1.01&0.91  &2.95&1942.0& 0.40 & 265.0 & 1928\\\\ \\hline {\\bf HD$\\thinspace$38529b\\tablenotemark{a}} & 0.78& 1.39& 0.13& 14.3& 0.28& 90.0 & 10005.8\\\\ {\\bf HD$\\thinspace$38529c\\tablenotemark{a}} & 12.78& 1.39& 3.71& 2207.4& 0.33& 13.0 & 10043.7\\\\ \\hline {\\bf HD$\\thinspace$160691b\\tablenotemark{b} } & 1.7$\\pm$0.2& 1.08& 1.5$\\pm$0.1 & 638$\\pm$10&0.31$\\pm$0.08& 320$\\pm$30 & 50698$\\pm$30\\tablenotemark{c}\\\\ {\\bf HD$\\thinspace$160691c\\tablenotemark{b}} & 1.0 &1.08& 2.3& 1300& 0.8& 99 & 51613\\\\  \\enddata \\tablenotetext{*}{In our calculations we use as an initial orbital parameter the Mean Anomaly which is a function of $T_{peri}$} \\tablenotetext{a}{Data from  http://www.exoplanets.org as on August 10, 2002} \\tablenotetext{b}{Data from Jones et al. 2002} \\tablenotetext{c}{HJD}  \\end{deluxetable} ",
        "conclusions": "\\begin{figure*} % \\centering \\includegraphics[totalheight=19cm, width=15cm]{f3.eps} \\caption{Stability maps in the $e_b - e_c$ parameter space for all four new planetary systems. The symbols  are the same as in Fig. 2 (filled circles indicate stable regular orbits).} \\end{figure*}  Fig. 3 presents a comparative visualisation of our MEGNO stability analysis for all four new planetary systems. For this visual presentation we choose to consider orbital stability as a function of both orbital eccentricities simply because the values of $e_b$ and $e_c$ can only be changed between 0 and 1, and therefore do not require any scaling (as for example would semi-major axes) for different systems. All other orbital parameters are the nominal ones from Table 1. It is easy to see the differences in dynamical status of the four systems and to identify the ranges of eccentricities which allow stable planetary orbits in each of them. We also produced simular maps for other pairs of orbital parameters. Such ranges of stable parameters can be very useful for the improvement of preliminary orbits in the new planetary  systems discussed in this paper. In our future papers we plan to present detailed dynamical analyses of each system taking into account angular orbital parameters not constrained by observational data ($i_r$x, $\\Omega$), as well as $\\sin i$ and the resulting different planetary masses."
    },
    "0207/astro-ph0207041_arXiv.txt": {
        "abstract": "We describe the properties of a remarkable group of actively star-forming dwarf galaxies and \\hii\\ galaxies in the Abell~1367 cluster, which were discovered in a large-scale \\halpha\\ imaging survey of the cluster. Approximately 30 \\halpha-emitting knots were identified in a region approximately 150 kpc across, in the vicinity of the spiral galaxies NGC 3860, CGCG 97-125 and CGCG 97-114.  Follow-up imaging and spectroscopy reveals that some of the knots are associated with previously uncataloged dwarf galaxies (M$_B$ = $-15.8$ to $-16.5$), while others appear to be isolated \\hii\\ galaxies or intergalactic \\hii\\ regions. Radial velocities obtained for several of the knots show that they are physically associated with a small group or subcluster including CGCG~97-114 and CGCG~97-125. No comparable concentration of emission-line objects has been found elsewhere in any of the eight northern Abell clusters surveyed to date. The strong \\halpha\\ emission in the objects and their high spatial density argue against this being a group of normal, unperturbed dwarf galaxies.  Emission-line spectra of several of the knots also show some to be anomalously metal-rich relative to their luminosities. The results suggest that many of these objects were formed or triggered by tidal interactions or mergers involving CGCG 97-125 and other members of the group.  A Westerbork Synthesis Radio Telescope HI map of the region shows direct evidence for tidal interactions in the group. These objects may be related to the tidal dwarf galaxies found in some interacting galaxy pairs, merger remnants, and compact groups. They could also represent evolutionary precursors to the class of isolated ultracompact dwarf galaxies that have been identified in the Fornax cluster. ",
        "introduction": "Until recently most surveys for star-forming galaxies in the nearby universe have been restricted to imaging of previously cataloged objects, or wider-field prism surveys that are mainly sensitive to strong emission-line galaxies.  These have provided a relatively complete inventory of massive galaxies and starburst galaxies, but they provide much less complete information on the population of star-forming dwarf galaxies.  Large numbers of very nearby star-forming dwarfs have been studied (e.g., Terlevich et al. 1991, Hunter, Hawley, \\& Gallagher 1993, van Zee 2000, 2001), but complete star formation inventories, extending across the full range of galaxy types {\\it and} masses, are lacking.  As part of a larger effort to obtain complete inventories of star formation rates (SFRs) in nearby galaxy samples, we have carried out a deep, wide-field H$\\alpha$ survey of nearby clusters of galaxies using the MOSAIC CCD camera on the 0.9 m telescope at Kitt Peak National Observatory (Sakai et al. 2001a). Each field covers one square degree, and we have obtained high-quality data for 25 fields in 8 nearby northern Abell clusters, in the radial velocity range 3000 $-$ 8000 km~s$^{-1}$. A total of six fields were observed in Abell~1367, and $\\sim$250 H$\\alpha$-emitting galaxies were detected (Sakai et al. 2002, in preparation). During the course of this analysis we discovered an unusual concentration of \\halpha-emitting dwarf galaxies and \\hii\\ galaxies in the central field of the Abell 1367 cluster. A brief report was given in Sakai et al. (2001b). Recently Iglesias-Paramo et al. (2002) presented the results of an independent H$\\alpha$ survey of the center of A1367, and they comment specifically on the unusual properties of this region of the cluster (see the appendix of their paper). Our follow-up spectroscopy (\\S 4) shows that these objects are part of a low velocity dispersion group or subcluster within or behind the main cluster, which also contains two Zwicky galaxies: CGCG~97-114 and CGCG~97-125.  Although the discovery of emission-line dwarf galaxies in Abell~1367 is not extraordinary in itself, the concentration of such objects in this region is very unusual --- to date no other such concentrations have been found in any of the eight clusters we surveyed.  Moreover, emission-line spectra obtained for several of the knots reveal chemical properties that are inconsistent with the scenario that this region is simply a grouping of normal star-forming dwarf galaxies.  Instead, the observational evidence suggests that at least some of the objects are the products of galaxy interactions or other environmental processes within the group, or in conjunction with the larger A1367 cluster. Consequently this serendipitously discovered group may offer valuable clues to the physical processes that influence the evolution and formation of dwarf galaxies in groups and clusters.  The remainder of the paper is organized as follows. In \\S2 we discuss the data collected on this region, including the \\halpha\\  and follow-up broadband imaging, emission-line spectroscopy, and HI aperture synthesis observations.  In \\S3 we use these data to characterize the nature of the galaxy group and measure the SFRs and basic physical properties of the star-forming dwarfs.  Finally in \\S4 we consider possible physical explanations for the nature and formation of these objects, and tentatively conclude that they are a combination of pre-existing dwarf galaxies, intergalactic HII regions, and possibly newly formed dwarf galaxies, all triggered by tidal interactions between the larger members of the group.  We also place these results in the context of other discussions of tidally-formed dwarf galaxies (e.g., Mirabel, Dottori, \\& Lutz 1992), and \\hii\\ regions (Iglesias-Paramo \\& V\\'ilchez 2001), and the recent discovery of compact blue galaxies in the Fornax cluster (Drinkwater et al. 2001, Phillips et al. 2001). ",
        "conclusions": "At the outset of this analysis we considered three physical explanations for this unusual group of galaxies:  1)  The CGCG~97-114/125 group is a normal compact group of field galaxies observed in projection behind A1367; it is nothing more than a usual collection of spiral and irregular galaxies which happen to be actively forming stars at approximately the same time.  2)  The unusual star formation properties of the group are the result of a strong interaction with the intergalactic medium in A1367, caused by shocking of the ISM as these galaxies move through the IGM with an encounter velocity of 1600 km~s$^{-1}$.  3)  The unusual star formation properties of the group are caused by current and/or past tidal encounters between two or more of the galaxies in the group, largely independent of the cluster environment outside of the group.  Our follow-up deep imaging, spectroscopic observations, and HI data appear to strongly favor the last of these interpretations, but we first summarize the evidence against the other scenarios.  While it is certainly true that many normal galaxy groups contain as many as several strongly star-forming dwarf galaxies, the concentration of so many starbursting dwarf galaxies and \\hii\\ galaxies in such a small region ($\\sim$100 kpc) is very unusual, apart from groups containing strongly interacting galaxies.  We would expect to find far more quiescent dwarf galaxies in the region if the starbursts we observe were not triggered by a common physical mechanism.  Moreover, the high metal abundances of intergalactic knots and the HI tails/bridges provide direct evidence for the importance of tidal processes.  Induced star formation from IGM interactions is not quite as easily ruled out, especially because evidence for such processes is found elsewhere in A1367 (Gavazzi et al. 1995).  However this interpretation appears to be unlikely on a number of grounds.  None of the galaxies in the CGCG~97-114/125 group show evidence of the bow-shock structure in the \\halpha\\ or broadband images, or asymmetric HI distributions that is characteristic of the other objects of this type in A1367 and elsewhere.  And perhaps more significantly, we would expect a 1600 km~s$^{-1}$ encounter between the galaxy ISMs and the A1367 IGM to produce copious soft X-ray emission, but Chandra maps of this region do not show evidence of this type of extended emission (Sun et al. 2001, Sun \\& Murray 2002).  Instead, several lines of evidence point to the liklihood that the unusual star formation properties of this group are triggered by one or more major tidal interactions within the CGCG~97-114/125 group. To summarize they include: 1) morphological evidence for disturbed dynamical structure of CGCG~97-125 and 97-114; 2) presence of massive tidal structures connecting many of the emission regions in the HI maps;  3) the low dispersion in radial velocities of the star-forming galaxies and emission knots, including those located outside of the HI features; and 4) near-solar metal abundances in some of the apparently isolated intergalactic \\hii\\ regions along the HI arm, and suspiciously high abundances in some of the starbursting dwarf galaxies.  The most straightforward interpretation of these observations is that the largest galaxy in the group, CGCG~97-125, has undergone at least one major tidal encounter with other members of the group, including a recent merger event that has produced its shell-like outer structure. These interactions have pulled metal-rich gas out of the galaxy, into an extended tidal tail or arm, and some of the gas has collapsed to form \\hii\\ regions or tidal dwarf galaxies in the HI arms.  Similar structures and star-forming regions are observed in some nearby examples of interacting galaxy pairs and merger remnants. Perhaps the closest analog is the Antennae system NGC 4038/9, which exhibits extended HI arms (Hibbard et al. 2001) with similar HI masses and a series of massive star-forming knots that have been proposed to be newly-formed tidal dwarf galaxies (Mirabel, Dottori, \\& Lutz 1992, Braine et al. 2001).  Tidal dwarf galaxy formation also has been purported to be occurring in the gaseous arms of other nearby interacting galaxies (e.g., Duc \\& Mirabel 1998, Weilbacher et al. 2000, Braine et al. 2001).  The morphology of the faint knots observed in the CGCG~97-114/125 group bear some resemblance to these systems, particularly in HI, but the main difference is the absence of continuous stellar counterparts to the HI arms; here the \\hii\\ regions are fainter and more isolated.  This might be explained if the tidal features in this group are older, and gravitationally unbound from parent galaxies already, or if the efficiency of star formation in these interactions were much lower for some reason.   The origin of the larger dwarf galaxies Dw~1, Dw~2, and Dw~3 is less clear.  The unusually blue colors and faint, diffuse underlying stellar components in these galaxies tempts us to speculate that these objects too may have been formed relatively recently (e.g., last 1--2 Gyr) in tidal interactions.  However the observed metal abundances of Dw~1 and Dw~3 are plausibly consistent with their being old irregular galaxies that have evolved independently. Some qualities that characterize tidal dwarf galaxies include the lack of dark matter and a small fraction of old stellar population (Hunter, Hunsberger \\& Roye 2000). Deeper imaging (at visible and near-infrared wavelengths) and measurements of the stellar kinematics, or HI rotational velocity would be able to discern the presence of an older stellar population, if any, and test whether these objects contain the dark matter halos expected for normal dwarf irregular galaxies. We may have found an example of a very recent interaction in which the disk of one of the galaxies (CGCG~97-114 with its low M$_{\\rm H \\rm I}$/L$_{\\rm B}$) was severely disrupted by a much more massive object (CGCG~97-125), leaving the shreds of the outer disk behind, which we are now witnessing as small star-forming regions.  It is intriguing to speculate on the eventual fate of the metal-rich \\hii\\ regions (e.g., K2).  These objects almost certainly are newly formed from the tidal debris of the galaxy interactions in this group, but it is unclear whether the associated star clusters will remain gravitationally bound to the more massive galaxies or will form new tidal dwarf galaxies.  Again, more accurate kinematic observations of the knots and the other galaxies in the region should allow one to fit a dynamical model to the HI and optical observations, and constrain not only the orbits of the knots but also the mass distributions in the halo of CGCG~97-125 and 97-114.  We also draw attention to the possible connection between these types of objects and the isolated compact dwarf galaxies that have recently been discovered in the Fornax cluster by Drinkwater et al. (2001) and Phillips et al. (2001).  The latter objects appear to be either tidally stripped remnants of dwarf galaxies or massive isolated star clusters.  It is conceivable that some of the star-forming regions observed in the CGCG~97-114/125 group may evolve into isolated intergalactic dwarf galaxies or star clusters in A1367, though it appears that the precursors to the massive objects observed in Fornax probably were considerably more massive than the regions we have observed here.  Finally we remark briefly on the the apparent rarity of groups of this kind in nearby galaxy clusters.  No other comparable subgroups or subclusters of star-forming regions have been found elsewhere in our Abell cluster survey, which covers 25 square degrees and a search volume of approximately 300 Mpc$^3$ in 8 clusters. This may not be entirely surprising, because if the star formation observed in this group has been triggered by tidal encounters it requires low-velocity interactions of order a few hundred km~s$^{-1}$ or less, which is much lower than the typical encounter velocities in these rich Abell clusters.  These considerations suggest instead that compact groups may be the most prolific environment for this mode of star and galaxy formation.  A deep \\halpha\\ imaging survey of compact groups by Iglesias-Paramo \\& Vilchez (2001) has revealed tidally extended star-forming regions in 5 of 16 groups surveyed.  These \\hii\\ regions probably are the closest analogs to the objects studied in this paper, though most of the emission knots found in the survey of Iglesias-Paramo \\& V\\'ilchez (2001) lie on well-defined tidal arms of large galaxies, or on well-defined tidal bridges connecting the interacting galaxies.  Perhaps deep imaging of other compact groups will reveal closer analogs to the concentration of star-forming galaxies in the CGCG~97-114/125 group. Until then this remarkable region appears to be unique."
    },
    "0207/astro-ph0207088_arXiv.txt": {
        "abstract": "V838~Mon has undergone one of the most mysterious stellar outbursts on record. The spectrum at maximum closely resembled a cool AGB star, evolving toward cooler temperatures with time, never reaching optically thin conditions or showing increasing ionization and a nebular stage.  The latest spectral type recorded is M8-9. The amplitude peaked at $\\Delta V$=9 mag, with the outburst evolution being characterized by a fast rise, three maxima over four months, and a fast decay (possibly driven by dust condensation). BaII, LiI and $s-$element lines were prominent in the outburst spectra. Strong and wide (500 km/sec) P-Cyg profiles affected low ionization species, while Balmer lines emerged to modest emission only during the central phase of the outburst. A light-echo discovered expanding around the object constrains its distance to 790$\\pm$30 pc, providing $M_V=+4.45$ in quiescence and $M_V=-4.35$ at optical maximum (dependent on the still uncertain $E_{B-V}$=0.5 reddening). The visible progenitor resembles a somewhat under-luminous F0 main sequence star, that did not show detectable variability over the last half century.  V838~Mon together with M31-RedVar and V4332~Sgr seems to define a new class of astronomical objects, {\\sl Stars that Erupt into Cool Supergiants (SECS)}.  They do not develop optically thin or nebular phases, and deep P-Cyg profiles denounce large mass loss at least in the early outburst phases. Their progenitors are photometrically located close to the Main Sequence, away from the post-AGB region. After the outburst, the remnants still closely resemble the precursors (same brightness, same spectral type). Many more similar objects could be buried among poorly studied variable stars that have been classified as Miras or SemiRegulars on the base of a single spectrum at maximum brightness. ",
        "introduction": "A detailed description of the outburst of V838~Mon is given by \\cite{Munari}, to which the reader is referred. In this note only the main features are summarized with some updates on the late photometric and spectroscopic evolution of the outburst.  An updated lightcurve of the eruption of V838~Mon is presented in Figure~1. A first maximum was reached by $+10^d$ (see abscissae scale on Figure~1) when the continuum energy distribution was characterized by a temperature of 4150~K.  A second maximum at $+37^d$ peaked around 5200~K and a third one at $+68^d$ reached 4600~K. Each decline from maxima was accompanied by a monotonic cooling, with the last one taking V838~Mon to 3500~K by $+90^d$. From $+90^d$ to $+120^d$ the color temperature in the region of the $V,R_c, I_c$ bands decreased to that of an M8-9 supergiant, or 2600~K. The retracing of the $U-B$ and $B-V$ color indexes when the spectrum developed the coolest temperatures is a real effect (cf. absolute spectrophotometry in Figure~2 at $+119^d$), and it is normally seen in M-type stars of the corresponding spectral types due to progressive disappearance of TiO absorptions at the shortest wavelengths.  \\begin{figure}[t] \\centering \\includegraphics[width=15.6cm]{Munari_Fig_1.ps} \\caption{$V$, $B-V$ and $V-I_C$ lightcurves of the outburst of V838~Mon.  Dots mark NOFS data, open circles Tsukuba data. Crosses and open triangles are values from various IAUC and VSNET circulars (mainly from SAAO, D.West, P.Sobotka, L.Smelcer, F.Lomoz and J.Bedient). The solid line indicates the quiescence brightness.} \\label{lightcurve} \\end{figure}  \\begin{figure}[h] \\centering \\includegraphics[width=14.8cm]{Munari_Fig_2.ps} \\caption{The spectrum of V838~Mon for April 29, 2002 obtained with WHT 4.2m in La Palma.} \\label{Decline_spectrum} \\end{figure}  \\begin{figure}[h] \\centering \\includegraphics[angle=270,width=14.8cm]{Munari_Fig_3.ps} \\caption{Small sections of sample of Asiago Echelle spectra to document the evolution around the far-red Calcium triplet and H$\\alpha$ of the V838~Mon outburst.} \\label{CaII} \\end{figure}  \\begin{figure}[h] \\centering \\includegraphics[width=14.8cm]{Munari_Fig_4.ps} \\caption{Expansion of the light-echo around V838~Mon, revealing a previously invisible ring of circumstellar material. $U$ band 67$\\times$67 arcsec images obtained with the USNO 1m telescope. Dates from left to right and top to bottom: January 13, February 27, March 10, March 27, March 31, April 4, April 20 and April 30.} \\label{ring} \\end{figure}  The spectral evolution well followed the $V-I$ color temperature evolution. During the first three months the spectrum closely resembled a K giant, slowing progressing toward later spectral types and reaching K5 by $+90^d$. In the following month the spectrum rapidly entered the M-type realm and reached M8-9 by $+119^d$. The spectrum of V838 Mon for this date is shown in Figure~2. The spectral evolution of V838~Mon has been exciting also on a much finer scale, as Figure~3 indicates, where Echelle spectra around the CaII far-red triplet and H$\\alpha$ are presented for $+25^d$, $+56^d$, $+86^d$ and $+112^d$. P~Cyg line profiles for low-excitation species have a terminal velocity which monotonically decreased with time from the initial value of $-$500 km/sec, with Balmer lines appearing in emission with their own P-Cyg profiles only after the second maximum. BaII, LiI and $s-$elements are present in the V838~Mon spectra.  In mid-February, \\cite{Henden} discovered the formation of a light-echo around V838~Mon, when the light from the second maximum began illuminating pre-existing circumstellar material responsible for the IRAS detection of the precursor. This light-echo was followed as it expanded to a maximum diameter of 35 arcsec, a size that has remained essentially constant during the following months, as Figure~4 shows. The light-echo expansion rate of 0.44$\\pm$0.017 arcsec day$^{-1}$ sets the distance of V838~Mon to 790$\\pm$30 pc for a spherical distribution of the scattering material. The outburst light sweeping through the circumstellar material allows us to read the recent mass loss history of the progenitor: assuming 15 km~sec$^{-1}$ velocity for its wind (typical for an AGB), the light-echo has reached by April~1 material lost $\\sim$4900 years ago. High resolution imaging with HST by \\cite{Bond} confirms the spherical symmetric dust distribution around V838~Mon and reveals multiple circularly-symmetric rings, along with a central void. This void was also visible on ground-based images after V838~Mon faded.  The void is the likely reason why the light-echo was not seen until some time after the rise to second maximum. The angular separation of the concentric rings in the HST images indicate a $\\sim$500 year recurrence time in the enhanced mass loss events.  Using the 790pc distance estimate and the peak brightness, we can derive $M_V=+4.45$ for V838~Mon in quiescence and $M_V=-4.35$ at peak outburst. The precise values depend on the exact amount of reddening, here estimated to be $E_{B-V}$=0.5.  At galactic coordinates $l=217.80$ $b=+1.05$, the height over the Galactic plane is just $z=$13 pc. It is relevant to note that the progenitor of V838~Mon was not detected by H$\\alpha$ emission-line surveys in the region (these surveys discovered several faint emission line stars close to V838~Mon), and inspection of Palomar and SERC plates as well as results from many archival plates presented at this conference by \\cite{Barsukova} reveal absence of photometric variability in quiescence. Both H$\\alpha$ emission and variability would have supported an interactive binary nature of the precursor. ",
        "conclusions": ""
    },
    "0207/astro-ph0207277_arXiv.txt": {
        "abstract": "\\object{4U\\,0614+09} is a low-mass X-ray binary with a weakly magnetized neutron star primary. It shows variability on time scales that range from years down to $\\sim 0.8$ milliseconds. Before the Chandra and XMM-Newton era, emission features around 0.7 keV have been reported from this source, but recent Chandra observations failed to detect them. Instead, these observations suggest an overabundance of Ne in the absorbing material, which may be common to ultracompact ($P_{orb} \\simless 1$ hour) systems with a neon-rich degenerate dwarf secondary.  We observed \\object{4U\\,0614+09} with XMM-Newton in March 2001. Here we present the energy spectra, both from the RGS and EPIC cameras, and the Fourier power spectra from EPIC high-time resolution light curves, which we use to characterize the spectral state of the source. ",
        "introduction": "\\object{4U\\,0614+09} is a low-luminosity X-ray binary. Thermonuclear (type I) X-ray bursts from \\object{4U\\,0614+09} were observed by OSO 8 (\\cite{mmendez-c1:swa78}) and WATCH (\\cite{mmendez-c1:bra92}), identifying the central source as a neutron star (as opposed to systems with black-hole candidate primary).  In \\object{4U\\,0614+09} the X-ray flux can vary by a factor of $\\sim 2 - 4$ (e.g., \\cite{mmendez-c1:stra00}) on timescales of days to months. Using EXOSAT data, \\cite*{mmendez-c1:bar95} found an anticorrelation between the high- and low-energy X-ray emissions in \\object{4U\\,0614+09}; this anticorrelation has been observed to extend up to 100 keV (\\cite{mmendez-c1:for96}).  Observation with RXTE have revealed strong quasi-periodic oscillations (\\cite{mmendez-c1:for97}; \\cite{mmendez-c1:men97}) that extend up to $\\sim 1300$ Hz kilohertz (\\cite{mmendez-c1:stra00}). These oscillations are thought to originate from matter in Keplerian orbit close to the central object. If this so, these quasi-periodic oscillations carry information about the strong gravitational field in the vicinity of the compact object.  The energy spectrum of \\object{4U\\,0614+09} can be approximated by a combination of a power law (sometimes with an exponentially cut-off at high energies), and a soft component, both affected by interstellar absorption. The soft component fits a blackbody, and is interpreted as the combined effect of emission from the surface of the neutron star and the accretion disc. The power law component is assumed to originate via comptonization of soft photons by hot electrons in a corona around the neutron star.  In \\object{4U\\,0614+09}, observations with EINSTEIN's Solid State Spectrometer have revealed emission features at $E \\sim 0.7$ keV (\\cite{mmendez-c1:chr94}; \\cite{{mmendez-c1:chr97}}); these features are thought to originate in a corona around the neutron star or above the disc (\\cite{mmendez-c1:chr94}).  Here we present a preliminary analysis of two observations of \\object{4U\\,0614+09} carried out in March 2001 with the instruments onboard XMM-Newton. We discuss spectral (both continuum and line features) and timing properties of the source. ",
        "conclusions": "\\label{mmendez-C1_sec:dis}  Our XMM-Newton observations of the low-mass X-ray binary \\object{4U\\,0614+09} found the source in the so-called island state, during which the energy spectrum fits a relatively flat power law (photon index $\\sim 2$), and the power spectrum shows a broad-band component that extends up to $\\sim 1$ Hz, with high rms variability ($\\sim 31$\\,\\% in this case).  Striking from the spectral fits is the excess emission (above the power law emission) at $\\sim 0.65$ keV that is apparent both in the MOS and RGS spectra. A similar excess has been reported by \\cite*{mmendez-c1:chr94} using the solid state spectrometer aboard EINSTEIN, and \\cite{mmendez-c1:whi97} using the solid state imaging spectrometers aboard ASCA. The low-energy excess reported by \\cite{mmendez-c1:chr94} is centered at around 0.77 keV, and has an equivalent width of $\\sim 40$ eV. In our case the excess is centered at a slightly lower energy ($\\sim 0.65$ keV), and we measure a larger equivalent width (200--300 eV). Christian et al. propose that this excess could be due to emission by Ly$\\alpha$ \\ion{O}{VII} and He-like \\ion{O}{VIII}, and \\ion{Fe}{XVII}--\\ion{Fe}{XIX} in a corona around the central object. Similar emission has been detected recently from other X-ray binaries, e.g. EXO\\,0748--67 (\\cite{mmendez-c1:cot01a}) and \\object{4U\\,1822--37} (\\cite{mmendez-c1:cot01b}). In those cases, however, the emission lines are narrow. It is possible that the excess that we measure is due to Oxygen radiative recombination continuum produced by transitions of continuum electrons to the ground state (e.g., \\cite{mmendez-c1:lie96}).  Alternatively, it is possible that this line-like ``feature'' is a consequence of assuming that the abundance of the absorbing material along the line of sight is solar. In fact, the feature disappears when an overabundance of \\ion{Ne}{}/\\ion{O}{} in the absorbing material with respect to the solar abundance is considered. If, as suggested by \\cite*{{mmendez-c1:jue01}}, this overabundance occurs in the vicinity of the binary system, these results are relevant within the evolutionary scenario of this, and similar X-ray binaries."
    },
    "0207/astro-ph0207107_arXiv.txt": {
        "abstract": "We argue that any violent galactic winds following early epoch of star bursts would significantly weaken the potentials of galaxies, and leave lasting signatures such as a lowered dark halo density and preferentially radial/escaping orbits for halo tracers such as globular clusters.  A galaxy is disintegrated if more than half of its dynamical mass is blown off.  The presence of dense halos in galaxies and the absence of intergalactic/escaping globulars should imply an upper limit on the amount of baryons lost in galactic winds around 4\\% of the total mass of the galaxy.  This translates to limits on the baryons participating the early star bursts and baryons locked in stellar remnents, such as white dwarfs. The amount of halo white dwarfs claimed in recent proper motion searches and microlensing observations in the Galactic halo are too high to be consistent with our dynamical upper limits. Similar arguments also imply upper limits for the amount of neutron stars and stellar black holes, in galaxy halos. Nevertheless, a milder outflow is desirable, especially in dwarf galaxies, both for lowering their cold dark matter central density and for injecting metals to the intergalactic medium. ",
        "introduction": "How strong is the energetic feedback from star formation in galaxies? What are the effects of massloss from a galaxy? These are important questions in understanding galaxy formation. Baryonic gas cools and forms stars in potentials of the dark matter halos.  Some massive stars explode as supernovae and the stellar envelope is ejected with high speed.  This feedback not only injects and mixs metals in galaxies, but also drives galactic winds to pollute intergalactic medium.  A sudden loss of baryonic gas by winds after an extremely powerful burst of star formation also means a sudden weakening of the gravitational potential of the galaxy.  The weakening of the potential has several interesting consequences. For example, the galaxy halo will relax to a looser distribution after a mild loss of its mass through winds.  This reduction of halo density is interesting for overcoming the dense cusp in Cold Dark Matter halos (e.g. Gnedin \\& Zhao 2001).  Stars could also escape a lowered potential well of a galaxy, and appear as intergalactic stars.  There are also limits on the amount of massloss.  A galaxy becomes unbound if more than half of its mass is lost (Hills 1980). Real galaxies appear to be tightly bound.  The boundness of present day galaxies should set some limits on the violentness of massloss and star formation in the past.  It is now well-established that galactic outflows are ubiquitous in actively star forming galaxies, both in the local universe and at high redshift (see a review by Heckman 2001).  The amount of gas lost in the energetic outflow is largely comparable to the amount of stars formed.  Such a wind can be efficient in polluting the hot gas in the intergalactic medium (IGM). Some amount of energetic wind or outflow seems mandatory to account for the presence of metals in the QSO absorption systems at high redshift, and metals in the hot gas in galaxy clusters (e.g., Bookbinder et al. 1980).  In the Milky Way there are very few metal-poor stars with $[Fe/H] \\le -4$ in the halo and $[Fe/H] \\le -0.6$ in the disk.  There are also metals in some distant high-velocity clouds in the Local Group.  These again suggest an early phase of metal production in a burst of mainly massive stars.  Several recent claims of detection of a substantial amount of white dwarfs (WDs) in the Galactic halo and distant galaxies by direct proper motion searches or indirect microlensing observations suggest a first burst of high-mass stars narrowly peaked around $(2-3)\\msun$ in galaxy halos.  To leave behind every half-solar-mass WD, a few times more of the remnant mass is returned to the interstellar medium via the planetary phase.  Those in binaries die after $\\sim 1$ Gyr via the highly energetic Type Ia SN, which can easily power a hot and fast wind and drive all stellar ejecta out of the potential well of a galaxy.  Motivated by these considerations, we model the dynamical effect of the post-star-formation galactic wind on the density distribution of the dark particle halo.  We examine the effect of a fast wind at a redshift $z \\sim 1$ when the Milky-Way-sized halos are already largely assembled, and energetic \\snia are observed.  In general, a galaxy will expand, stellar orbits become highly radial in response to a rapid weakening of the potential well due to the mass gone with the wind.  In particular we set limits on the amount of high-mass star formation in galaxies in the past, and the amount of stellar remnents, e.g., white dwarfs, in present day galaxies.  Simple models for the wind and expansion are developed in \\S2 and the Appendix, including the effects on globular clusters. Applications to the halo white dwarf problem are made in \\S3.  We discuss our results in \\S4, in particular, the strength and starting time of the wind, and conclude in \\S5. ",
        "conclusions": "In short, the absence of obvious signs of severe expansion of the Milky Way and the absence of escaping globulars in galaxies in general suggest that star formation and feedback in galaxies are mild; galactic winds carry less than 4\\% of the total gravitational mass, or 25\\% of the baryonic mass in galaxies.  This restricts the amount of remnents in galaxy halos from early star formation, and cool halo WDs can make up 2\\% of the total mass in present day galaxies.  This is consistent with only the lower end of the halo WD fraction of $2\\% \\le f_{WD} \\le 50\\%$ from direct and indirect detections of WDs; it is consistent with Oppenheimer et al. (2001) detection, but inconsistent with the higher values from Alcock et al. (1997, 2000) and Ibata et al. (2000).  Our dynamical limit is also tighter for very cool WDs than for hot WDs.  There is very limitted room to fit in a large extra population of very cool WD population $\\le 4000$ K, something that future deeper surveys should take into account.  Our model implies stringent limits on the amount of black hole (BH) or neutron star (NS) formation as well. Assuming these are remnents of massive stars formed in galaxy halos, then typically neutron stars or stellar black holes are formed from very massive ($8\\msun-1000\\msun$) progenitors with $m_{PG}/m_{BH} \\ge m_{PG}/m_{NS} \\sim 6-30$.  So the massloss is even more severe than for forming WDs.  The progenitors evolve very quickly off the main sequence, and some explode as Type II SN, which immediately power a very fast wind. To keep the wind from severly damage a galaxy (cf. eq.~\\ref{eps} and eq.~\\ref{mild}) would imply that these remnents make up a fraction \\beq f_{BH} \\le f_{NS} \\le (0.1-1)\\% \\eeq of the total mass of a galaxy. In comparison, the upper limit on stellar BHs from microlensing is $f_{BH} \\le (30-100)\\%$ (Lasserre et al. 2000, Alcock et al. 2001). So our dynamical limit is much tighter. Some models of galaxy formations suggest a significant amount of massive black holes in galaxy halos, with black hole mass $m_{BH} \\sim 10^2-10^6\\msun$ (Lacey \\& Ostriker 1985, Madau \\& Rees 2001).  These black holes, however, must be pregalactic, in which case our dynamical limits do not apply.  In summary, we show that intergalactic stars and globular clusters and the present day density profiles of galaxies offer a new diagnosis of early star bursts of massive stars in halos of galaxies, and the amount of remnents, such as halo WDs and BHs. A mild wind-induced expansion of the dark halos might have played some role in lowering the dense cold dark matter central density in dwarf galaxies (Gnedin \\& Zhao 2001) and high surface brightness galaxies as well.   HSZ thanks Bernard Carr, Rodrigo Ibata, Mike Irwin, Gerry Gilmore, Donald Lynden-Bell, Jerry Ostriker, Tom Theuns and the anonymous referee for helpful comments.  \\appendix"
    },
    "0207/astro-ph0207331_arXiv.txt": {
        "abstract": "{The linear stability of MHD Taylor-Couette flow of infinite vertical extension  is considered for various magnetic Prandtl numbers Pm. The calculations are  performed for a wide gap container with $\\hat\\eta=0.5$ with an axial uniform magnetic field excluding counterrotating cylinders. For both hydrodynamically stable and unstable flows the magnetorotational instability produces characteristic minima of the Reynolds number for certain (low) magnetic field amplitudes and Pm $>$ 0.01. For Pm $\\lsim$ 1 there is a characteristic magnetic field amplitude beyond which  the instability sets in in form of nonaxisymmetric spirals with the azimuthal number $m=1$. Obviously, the magnetic field is able to excite nonaxisymmetric configurations despite of the tendency of differential rotation to favor axisymmetric magnetic fields which is known from the dynamo theory. If Pm is too big or too small, however,  the axisymmetric mode with $m$=0 appears to be the most  unstable one possessing  the lowest  Reynolds numbers --  as it is also true for hydrodynamic Taylor-Couette flow or for very weak fields. That the most unstable mode for modest Pm proves to be nonaxisymmetric must  be  considered as a strong indication for the possibility of dynamo processes in connection with the magnetorotational instability. ",
        "introduction": "In order to discuss possible experimental realizations of the magnetorotational instability as the main transporter of angular momentum in all kinds of accretion disks there are several recent studies of Taylor-Couette flow for electro-conducting fluids between rotating cylinders under the influence of an uniform axial magnetic field (Ji et al. 2001; R\\\"udiger \\& Zhang 2001; Willis \\& Barenghi 2002). The numbers describing the geometry of the container and the magnetic Prandtl number of the fluid have been considered as the free parameters. For a given magnetic field amplitude (the Hartmann number) the critical angular velocity of the inner cylinder (the critical Reynolds number) is computed for the onset of an instability of the rotation law between the cylinders.  In R\\\"udiger \\& Shalybkov (2002) the instability pattern is considered as axisymmetric. The main result for resting outer cylinder is that for high magnetic Prandtl number  for weak magnetic field the excitation of the instability is easier than without magnetic field but for strong magnetic field the excitation of the instability is more complicated. The effect, however, disappears for small magnetic Prandtl number, i.e. for lower electric conductivity of the fluid as it may be realized in protoplanetary disks.  On the other hand for rotating outer cylinder, when no instability without magnetic field exists, the magnetic field always produces critical Reynolds numbers which, however, are running with 1/Pm. For Pm of order $10^{-5}$ the critical Reynolds number is of order $10^6$ which is just the experimental limit. \\begin{figure} \\psfig{figure=cylinder_new.ps,width=6cm,height=11cm} \\caption{\\label{geometry} Cylinder geometry of the  Taylor-Couette flow} \\end{figure}  In the present paper the nonaxisymmetric perturbations are included into the consideration. This is of particular relevance for the question whether the Cowling theorem for dynamo action can be fulfilled, after which a dynamo can only work with nonaxisymmetric fields. We shall find that indeed for certain parameters  -- despite the smoothing action of the differential rotation -- nonaxisymmetric modes can be excited easier than axisymmetric modes. This is in great contrast to earlier results of Taylor-Couette flow without magnetic fields where always the axisymmetric modes possess the lowest Reynolds numbers (Roberts 1965; DiPrima 1961)\\footnote{For counterrotating cylinders, however, the preference of nonaxisymmetric modes is already known, see Kr\\\"uger et al. 1966; Chen \\& Chang   1998 }.  Here,  the dependence of a real Taylor-Couette flow on the magnetic Prandtl number and on the azimuthal `quantum number $m$' is investigated.  The simple model of uniform  density fluid contained between two vertically-infinite rotating cylinders is used with a constant magnetic field parallel to the rotation axis. The unperturbed state is a stationary circular flow with $\\Om$ \\begin{equation} \\Om(r) = a+b/{R}^2, \\label{Om} \\end{equation} where $a$ and $b$ are two constants related to the angular velocities $\\Om_{\\rm in}$ and $\\Om_{\\rm out}$ with which the inner and the outer cylinders are rotating. If $R_{\\rm in}$ and $R_{\\rm out}$ ($R_{\\rm out}>R_{\\rm in}$) are the radii of the two cylinders then \\begin{equation} a={\\hat \\mu-{\\hat\\eta}^2\\over1-{\\hat\\eta}^2} \\ \\Om_{\\rm in}, \\quad\\quad\\quad b=  {1-\\hat\\mu\\over1-{\\hat\\eta}^2}\\ \\Om_{\\rm in} \\ R_{\\rm in}^2, \\label{ab} \\end{equation} with \\begin{equation} \\hat\\mu=\\Om_{\\rm out}/\\Om_{\\rm in} \\q {\\rm and} \\q \\hat\\eta=R_{\\rm in}/R_{\\rm out}. \\label{mueta} \\end{equation} After the Rayleigh stability criterion, $d(R^2 \\Om)^2/dR>0$, rotation laws are hydrodynamically stable for $\\hat\\mu>\\hat\\eta^2$. Taylor-Couette flows with resting outer cylinders ($\\hat\\mu=0$) are thus never stable. Here in order to isolate the MRI, we are also interested in flows with rotating outer cylinders  so that the hydrodynamical stability criterion,   $\\hat\\mu>\\hat\\eta^2$,  is fulfilled. Our standard examples are formed with $\\hat\\eta=0.5$, $\\hat\\mu=0$ and $\\hat\\mu=0.33$, resp. The first example ($\\hat\\mu=0$) is hydrodynamically unstable and the second one  ($\\hat\\mu=0.33$) is hydrodynamically stable. We are here only interested  in flow patterns  in containers with positive $\\hat \\mu$. ",
        "conclusions": "We have shown that  a Taylor-Couette flow which is stable in the hydrodynamic regime ($\\hat \\mu  \\geq \\hat \\eta^2$) is destabilized by a weak axial magnetic field. Below a critical Hartmann number of order 10... 100  the instability sets in in form of axisymmetric rolls while above this value the instability  forms nonaxisymmetric field and flow modes. This phenomenon exists despite of the observation (e.g. in dynamo theory) that differential rotation is known as suppressing the formation of nonaxisymmetric magnetic fields.  On the other hand, after the Cowling theorem of dynamo theory a magnetic field  can only be maintained if it is nonaxisymmetric. Considering a  number of typical magnetic Prandtl numbers we find  that for our container with conducting cylinders the dominance of the nonaxisymmetric modes only occurs  for not too high and not too low magnetic Prandtl number. Obviously, the dissipation processes are more important for nonaxisymmetric modes rather than axisymmetric modes. Hence the dissipation  allows  nonaxisymmetric modes only to be preferred if both the dissipation values have nearly the same order of magnitude."
    },
    "0207/astro-ph0207020.txt": {
        "abstract": "{In this paper we study long slit spectra in the region of H$\\alpha$ emission line of a sample of 111 spiral galaxies with recognizable and well defined spiral morphology and with a well determined environmental status, ranging from isolation to non-disruptive interaction with satellites or companions.  The form and properties of the rotation curves are considered as a function of the isolation degree, morphological type and luminosity. The line ratios are used to estimate the metallicity of all the detected HII regions, thus producing a composite metallicity profile for different types of spirals. We have found that isolated galaxies tend to be of later types and lower luminosity than the interacting galaxies. The outer parts of the rotation curves of isolated galaxies tend to be flatter than in interacting galaxies, but they show similar relations between global parameters. The scatter of the Tully-Fisher relation defined by isolated galaxies is significantly lower than that of interacting galaxies. The [NII]/H$\\alpha$ ratios, used as metallicity indicator, show a clear trend between Z and morphological type, t, with earlier spirals showing larger ratios; this trend is tighter when instead of t the gradient of the inner rotation curve, G, is used; no trend is found with the interaction status. The Z-gradient of the disks depends on the type, being almost flat for early spirals, and increasing for later types. The [NII]/H$\\alpha$ ratios measured for disk HII regions of interacting galaxies are higher than for normal/isolated objects, even if all the galaxy families present similar distributions of H$\\alpha$ Equivalent Width. \\keywords {Galaxies: spiral -- kinematics and dynamics -- structure -- interaction} } ",
        "introduction": "The analysis of the rotation curves of disk galaxies is the most direct way to obtain information on the mass distribution of spiral galaxies. The ionized gas has been used for long as a tracer of their kinematics. During the 80's, Rubin and collaborators started a systematic effort to obtain accurate rotation curves of spiral galaxies of all morphological types and luminosity (Rubin et al. 1991, and references therein). The accumulation of data from different sources helped to get an overall picture of the form of the rotation curve of spirals, and its relation with other galactic properties. It is now recognized that the maximum rotational velocity, V$_{m}$, is related with the total mass (and luminosity) of the galaxy, with the optical scale radius of the disk and with the morphological type (see Persic et al. 1996). The flatness of the outer rotation curve in most cases also led to accept the existence of massive dark halos in spiral galaxies (see Rubin et al. 1991; Sofue 1998).  Most of those analysis were based on data sets assembled with no completeness criteria. In particular, the galaxies were considered as field or cluster objects, and no further attention was payed to their environmental status, in spite of the expected influence of even small companions on the mass distribution, and star formation history of a given galaxy. Spiral galaxies in very close isolated pairs were studied by Keel (1993, 1996). Trying to %study understand the effects of the interaction on the dynamics of disk galaxies, M\\'arquez \\& Moles (1996; hereafter Paper I) studied the properties of isolated (see below for the definition of isolation) spiral galaxies, to set a zero-point for the effects of the interaction; see also Mathewson et al. 1992, and Courteau 1997, for an analysis of field spirals). M\\'arquez \\& Moles (1999; hereafter Paper II), studied also the properties of spirals in isolated pairs, and compared them to those of the isolated galaxies in Paper I. It was found that the main differences %were - is the presence of type II disk profiles in interacting systems (but not in isolated galaxies), and a flatter outer rotation curve in isolated galaxies. No distorted curves were found among isolated disk galaxies.  More recently, 2D Fabry-Perot rotation curves have been obtained for a number of cluster spirals in order to determine the environmental effects in such large aggregates. The results show a complex pattern (Amram et al. 1996). Barton et al.  (2000, 2001) have also analyzed the rotation curves of spiral galaxies in close pairs and in the general field in order to put some constraints on the cluster effects on the kinematical properties of galaxies, and the consequences in their use for distance estimation by means of the Tully-Fisher relation. Their results do confirm the earlier results in Paper II, in the sense of a more scattered T-F relation for non isolated objects. Similarly, galaxies with interacting companions in the recent analysis by Kannappan et al. (2002) fall on the high luminosity/low velocity width side of the TF and show more scatter.  We emphasize that the so called field galaxies should be carefully investigated since some of them could still be perturbed by small companions or satellites, that could produce significant effects (Athanassoula 1984; Conselice \\& Gallagher 1999; Conselice et al. 2000). In Paper I and II a quantitative criterion of isolation was given, trying to identify truly isolated objects to build up a reference for the analysis of the effects of gravitational interaction. We will use a similar approach here.  The same long slit spectroscopic data used for the study of the gas kinematics can also be used, through the flux ratios of the observed emission lines, to trace the metallicity, Z, along the disk, and to determine the existence of Z-gradients. The existing analysis point out that the global metallicity is related to the mass (hence, to V$_{m}$), and that Z decreases gently outwards (see the review by Henry \\& Worthey 1999). Ferguson et al. (1998) have extended the analysis towards the extreme outer regions of disks, finding that Z drops there abruptly, but keeping values far from pristine.  In the present paper, we will study a sample of 111 spiral galaxies with a well studied environmental status, ranging from isolation to mild interaction with satellites or companions. In all cases however the interaction is non disruptive (they have been selected to have recognizable and well defined spiral morphology). The data comprise new long slit spectroscopy for 85 spiral galaxies. The remaining data are from Paper I. The form and properties of the rotation curves are considered as a function of the isolation degree, morphological type and luminosity.  The line ratios are used to estimate the metallicity of all the detected HII regions, thus producing a composite metallicity profile for different types of spirals.  Section 2 is devoted to the description of the sample and the determination of the interaction status. In Section 3, the observations and data reduction procedures are presented. Section 4 and 5 deal with the rotation curves and the Tully-Fisher relation, respectively. In Section 6 the properties of nuclear and extranuclear HII regions are described. The summary and conclusions are given in Section 7. ",
        "conclusions": ""
    },
    "0207/cond-mat0207289_arXiv.txt": {
        "abstract": "We analyze the phenomena of condensate collapse, as described by Donley et al \\cite{JILA01b,Claussen03}, by focusing on the behavior of excitations or fluctuations above the condensate, as driven by the dynamics of the condensate, rather than the dynamics of the condensate alone or the kinetics of the atoms. The dynamics of the condensate squeezes and amplifies the quantum excitations, mixing the positive and negative frequency components of their wave functions thereby creating particles which appear as bursts and jets. By analyzing the changing amplitude and particle content of these excitations, our simple physical picture explains well the overall features of the collapse phenomena and provide excellent quantitative fits with experimental data on several aspects, such as the scaling behavior of the collapse time and the amount of particles in the jet. The predictions of the bursts at this level of approximation is less than satisfactory but may be improved on by including the backreaction of the excitations on the condensate. The mechanism behind the dominant effect -- parametric amplification of vacuum fluctuations and freezing of modes outside of horizon -- is similar to that of cosmological particle creation and structure formation in a rapid quench (which is fundamentally different from Hawking radiation in black holes). This shows that BEC dynamics is a promising venue for doing `laboratory cosmology'. ",
        "introduction": "We introduce a new perspective in the analysis of the phenomena of condensate collapse, described by Donley et al \\cite{JILA01b,Claussen03}, by focusing on the behavior of fluctuations above the condensate, rather than the condensate itself. We show that the condensate dynamics squeezes, amplifies, and mixes positive and negative frequency components of the wave functions of the condensate excitations. In addition to providing a good qualitative understanding of the general picture our theory also produces precise predictions, specifically, on the critical number of particles at the first instance when the instability sets in, the scaling of the waiting time $t_{collapse}$ and the number of particles in a jet. In this rendition we point out the analogy between the evolution of quantum excitations of the collapsing condensate and the vacuum fluctuations parametrically amplified by the background spacetime in the Early Universe, suggesting a new venue for ``laboratory cosmology''.  A condensate formed from a gas of cold ($3$nK) Rubidium atoms is rendered unstable by a sudden inversion of the sign of the interaction between atoms. After a waiting time $t_{collapse},$ the condensate implodes, and a fraction of the condensate atoms are seen to oscillate within the magnetic trap which contains the gas (see below and \\cite{JILA01b}). These atoms are said to belong to a ``burst''. In the experiments described by Donley et al. -- to single out this controlled BEC collapse experiment from the others, we shall adhere to the namesake Bose Novae -- the interaction is again suddenly turned off after a time $\\tau_{evolve}.$ For a certain range of values of $% \\tau _{evolve},$ new emissions of atoms from the condensate are observed, the so-called ``jets''. Jets are distinct from bursts: they are colder, weaker, and have a distinctive disk-like shape.  The Donley et al. experiment takes full advantage of the tunability of the effective atomic interaction due to a Feshbach resonance characteristic of $% ^{85}$Rubidium \\cite{JILA98,JILA00}. The resonance is caused by the presence of a bound state whose binding energy may be tuned by means of an external magnetic field. In later experiments \\cite{JILA02a,JILA02b}, observed fluctuations in the number of particles in the condensate have been well-explained as arising from oscillations between the usual atomic condensate and a molecular state \\cite {KGB02,KH02,MSJ02a,MSJ02b,Mackie02,Yin03}.  These oscillations were observed for magnetic fields in the order of $160$G, where the effective scattering length is of the order of $500a_{0}$ (and positive) ($a_{0}=0.529$\\ $10^{-10}$m\\ \\ \\ is the Bohr radius) and the frequency of oscillations is of hundreds of KHz \\cite{JILA02a,JILA02b}. By contrast, in the Donley et al. experiment \\cite{JILA01b}\\ typical fields were around $167$G, the scattering length was only tens of Bohr radii (and negative) and the frequency of atom - molecule oscillations may be estimated as well over ten MHz \\cite{JILA03}. Under these conditions it is unlikely that the molecular condensate plays any important dynamical role, and indeed no oscillations are reported in the original paper (for the opposite view, see \\cite{MMH03}). For these reasons and to highlight the mechanism particular to this experiment, we shall not include explicitly a molecular condensate in our model but discuss in detail the one - field model. However, if need to, this may be done in a very simple way, by including a second field to describe molecular destruction and creation operators \\cite {MSJ02a,MSJ02b,MMH03}. We will elaborate on this point in a later subsection  There is a vast literature attempting to provide theoretical explanations of collapsing condensates \\cite{SHU96,SSH98,KMS98,BR00,HUA01,AD02,SU03}. In addition to speculations that Bose Novae is due to molecular oscillations as alluded above (which we view as of secondary importance) the most serious theoretical attempt is based on the Gross-Pitaevskii equation with explicitly introduced nonlinear terms to account for multiparticle interactions \\cite{SU03,SRH02}. We will show that the primary mechanism responsible for the main features of the Bose Novae experiment originates from the dynamics of quantum fluctuations around the background condensate field(s). We start with the Heisenberg operator for the many body wave function and split it into a c-number part describing the condensate amplitude and a q-number part describing collective excitations (not individual atoms) above the condensate. We then derive an evolution equation for the wave function operator of the quantum (non-condensate) excitations under an improved Hartree approach, the so-called Popov approximation \\cite {POP87,GRI93,GRI96,PS02}.  In this paper, we use a ``test field'' approximation, by adopting (rather than deriving) the specific evolution of the condensate extracted from the experiment as given and study the dynamics of the excitations riding on this dynamics. Note that the experimentally given condensate dynamics is different from the mean field dynamics obtained from a solution of the GP equation, because the former includes the dynamical effects of the fluctuations. Finding a self-consistent solution of the evolution equations for both the condensate and its fluctuations is called the `backreaction problem'. It has been studied in detail in problems of similar contexts such as cosmological particle production (see below). Theoretical investigations for BEC fluctuations dynamics can be found in Refs. \\cite{KGB02,KB02,DS02}. The squeezing of quantum unstable modes and its back reactions on the condensate has been considered before, e.g., as a damping mechanism for coherent condensate oscillations \\cite{KM01}, and also applied to the description of condensate collapse \\cite{VM02,Y02,VYA01,YBJ02,R76}. Field theory methods have recently been applied to the problem of formation and stability of Bose condensates \\cite{field,BWLYA01}. Fluctuations have also been considered by G\\'{o}ral et. al. \\cite{GGR00} and Graham et al. \\cite {GWCFW96}. Our work differs from them in the emphasis we place on the behavior of the quantum excitations as a consequence of condensate dynamics.  Particularly relevant to the present work is Ref. \\cite{Y02}, where condensate collapse is analyzed from the point of view that the physics is mainly due to the dynamics of quantum fluctuations, the same view as we hold here. There, the trapping potential is replaced by a normalizing box, whose volume is eventually taken to infinity. Our analysis in Sections 2 and 3 is for a more realistic geometry, which enables us to compare quantitatively to experiments. The analysis of bursts versus jets given in Section 4 however originates from a new concept inspired by cosmological processes.  To the extent that many phenomena observed in connection to the collapse of this nature (Bose Novae) are essentially the result of a quantum fluctuation field (the non-condensate) interacting with a time dependent background (the condensate), as we believe it is, there is a close analog with similar processes in the early universe, specifically, vacuum particle creation from an time dependent external field \\cite{Schwinger} or in a curved background spacetime \\cite{Parker,Zeldovich}. (For a squeezed state depiction of this process, see, e.g., \\cite{HKM} and references therein.) One could view condensate collapse as a laboratory realization of cosmological particle creation during quenching. (Note this is not the physical process behind black hole particle creation, as in the Hawking effect, much attention drawn to its detection in BEC notwithstanding \\cite{HawEffBEC}.) In this process there is a competition between two (inverse) time scales, the physical frequency of the mode under consideration, and the inverse collapse (expansion) rate of the condensate. In cosmology the inverse expansion rate is the Hubble constant for the background spacetime. While a mode whose physical frequency is higher than the Hubble constant, we refer to it as ``inside the horizon'', and its behavior is oscillatory. When the converse obtains, the mode is ``outside the horizon''. They are depicted as `frozen' because they do not oscillate (see below), but are amplified \\cite{STA86}. This amplification is largely responsible for the observed primordial density contrast in the Universe \\cite{BM}.  In the Bose-Novae collapse problem, the role of the ``Hubble'' constant is played by the inverse growth (exponential) rate of the most unstable mode of the condensate, which is determined by the instantaneous number of particles in the condensate. Modes whose natural frequency is greater than the corresponding scale are relatively impervious to the dynamical condensate, but when the converse obtains, consequences are drastic. When the exponential growth is the dominant factor, the mode is frozen; instead of oscillating, it is being amplified, a process which is analogous to the growth of fluctuations during spinodal decomposition \\cite{SD}.  In the same way that modes that left the horizon during inflation return during the radiation and matter dominated eras, giving rise to acoustic oscillations, as the unstable condensate sheds its atoms and approaches stability, the band of ``frozen'' modes narrows: we say that modes ``thaw'' as they turn from exponentially increasing to oscillatory behavior. The crux of the matter is that only oscillating modes are detected through destructive absorption imaging (see below). Whenever a mode thaws, it is perceived as if particles were being created. In the conditions of the experiment the initial number of actual particles above the condensate is negligible, and we may describe the phenomenon as particle creation from the vacuum.  To summarize, the key idea in our understanding of the phenomena associated with a condensate collapse is that of a dynamical background field of the condensate squeezing and mixing the positive and negative frequency components of its quantum excitations, thereby creating particles from the vacuum. The viewpoint of this work may be easily incorporated in a first principles approach as taken in e.g. \\cite{STO99,DS01b,GZ00,GAF01}. The remarkable analogy between condensate collapse and quantum processes in the Early Universe and spinodal decomposition in phase transitions may stimulate new related experiments in BEC to be carried out to address these problems in cosmology and condensed matter physics \\cite{other}.  This paper is organized as follows. Section II we briefly review the phenomenology of condensate collapse and set up the basic mathematical model. In Section III we give a discussion of the onset of instability and of the scaling of the waiting time $t_{collapse}$ with the scattering length. In Section IV, we turn to a discussion of bursts and jets, based on the distinction between frozen and thawed modes. By postulating a specific condensate evolution (extracted from the experiment) we obtain quantitative predictions for the number of particles in a jet as a function of the time $% \\tau _{evolve}$ (when the scattering length is brought to zero). Our results are summarized in the final Section. A few technical details are left to the Appendixes. ",
        "conclusions": "In this paper, we have applied insights from the quantum field theory of particle creation and structure formation in cosmological spacetimes and the theory of second order phase transitions to a specific scenario of controlled collapse of a Bose-Einstein condensate, the so-called Bose Novae phenomena. We have described these phenomena as resulting from particle creation from the vacuum, induced by the time dependent condensate. This time dependence squeezes and amplifies the field operator describing excitations above the condensate. A key concept in our analysis borrowed from theories of cosmological structure formation is the drastic difference in the physical effects of frozen versus oscillatory modes: those whose physical frequencies are higher than the collapse rate oscillate and are rather impervious to the condensate, while those below (frozen modes) grow in time and get amplified, in a way similar to the growth of fluctuations during spinodal decomposition. As the condensate stabilizes and the collapse rate decreases the frozen modes begin to thaw. The appearance of oscillatory modes (in second quantized language) is described as particle creation appearing in jets and bursts, as described in detail above.  In order to focus on the key ideas we have adopted a number of simplifying assumptions. We take the condensate evolution as a given input from the experiments, rather than deriving it from fully self-consistent equations. We have treated excitations within the Popov approximation, which improves on the Hartree approach but is known to break down as the number of particles above the condensate increases. We have neglected the coupling between different excitation modes, considering only the coupling of each to the condensate.  These simplifications render certain aspects of the problem more amenable to others because they are rather insensitive to the assumptions. The scaling of $t_{collapse}$ is shown to depend on the behavior of a few modes setting the characteristic time scale of the problem - therefore the prediction is not affected by the underestimation of the coupling to other modes.  Even within these simplications, we have obtained good quantitative predictions for the onset of instability, the scaling of the waiting time $% t_{collapse}$ (when the condensate implosion really begins after the inversion of the scattering length) with the scattering length, and also for the number of particles in a jet as a function of $\\tau _{evolve}$, when the interaction between atoms is switched off.  Another success of the model is to provide a simple explanation for the widely different appearance of bursts and jets. As remarked earlier, jets may only appear if the turn - off time $\\tau _{evolve}$ is earlier than the formation of the remnant, because once the condensate is stable again, there are no more frozen modes to thaw, but, on the other hand, jets will appear for $\\tau _{evolve}<t_{collapse}$, when the condensate has not yet shed any particles. Also jets must be less energetic than bursts, since they are composed of lower modes.  Considering the success these simple ideas and light calculations brought about we believe our approach might have captured the essence of the physics behind these phenomena. The physical paradigms we used in bringing forth these ideas also suggest that understanding the basic mechanism of important processes in cosmology, critical dynamics, and Bose-Einstein condensation may share more than a superficial ground.\\newline  \\textbf{Acknowledgement} We first learned about this neat experiment from Bill Phillips who also introduced us to Keith Burnett whose `industrial strength' colloquium at Maryland and discussions afterwards provided us with the background information about Bose Novae. Our colleagues Ted Jacobson and Stefano Liberati have been looking for Hawking radiation and analog models of gravity in BEC, thus providing our group with a supercool ambience. We thank them all for enticing us to join in the BEC fun. The audiences in EC's seminar at UMD and BLH's seminar at NIST reporting on these results are to be thanked for their questions. We are obliged to E. Donley and S. Kokkelmans for their prompt and informative responses to our questions on BECs and for communicating key unpublished data. EC also acknowledges discussions with Eric Bolda. This research and EC's visits to UMD are supported in part by a NSF grant PHY98-00967, a NIST grant and by CONICET, UBA, Fundacion Antorchas and ANPCyT under grant PICT99 03-05229."
    },
    "0207/astro-ph0207457_arXiv.txt": {
        "abstract": "We report the first results from a study of the internal kinematics, based on spatially resolved H$_{\\alpha}$ velocity profiles, of three galaxies at redshift z$\\sim$0.6 and one at redshift z$\\sim$0.8, detected by ISOCAM in the Hubble Deep Field South. The kinematics are derived from high resolution near-infrared VLT spectroscopy. One of the galaxies is a massive spiral which possesses a very large rotational velocity of 460 km s$^{-1}$ and contains a mass of 10$^{12}$ M$_{\\odot}$ (within 20 kpc), significantly higher than the dynamical masses measured in most other local and high redshift spirals. Two of the galaxies comprise a counter-rotating interacting system, while the fourth is also a large spiral. The observed galaxies are representative examples of the morphologies encountered among ISOCAM galaxies. The mass-to-light (M$/$L$_{bol}$) ratios of ISOCAM galaxies lie between those of local luminous IR galaxies and massive spirals. We measure an offset of  1.6$\\pm$0.3 mag in the rest frame B-band and of 0.7$\\pm$0.3 mag in the rest frame I-band when we compare the four ISOCAM galaxies to the local Tully-Fisher B and I-band relations. We conclude that the large IR luminosity of the ISOCAM population results from a combination of large mass and efficient triggering of star formation. Since ISOCAM galaxies contribute significantly to the Cosmic Infrared Background our results imply that a relatively small number of very massive and IR luminous objects contribute significantly to the IR background and star formation activity near z$\\sim$0.7. ",
        "introduction": "Most studies of high redshift galaxies have targeted their statistical properties such as galaxy counts, colors and redshift distributions (e.g. Koo \\& Kron~1992; Lilly et al.~1995). The information obtained this way has been used to construct the luminosity function which when compared with various models can give us clues about the evolution of galaxies at earlier epochs (e.g. Gronwall \\& Koo~1995). The advent of 10m class telescopes combined with advances in detector technology offers the possibility of a more detailed investigation of the high redshift population aimed at better understanding their physical properties and internal kinematics, the latter being ultimately related to the question of their evolution. Such a detailed study of the internal kinematics of high-redshift systems has the advantage of studying the evolution of galaxies on a galaxy by galaxy basis, in contrast to redshift surveys that study a large sample of similar but not exactly identical galaxies and derive ``mean'' properties that are then applied to the entire sample.  However, measuring the evolution of mass$/$luminosity of high-redshift galaxies directly is a challenging task. The Faber-Jackson relation for spheroid-dominated systems (Faber \\& Jackson~1976) and the Tully-Fisher relation for spiral galaxies  (Tully \\& Fisher~1977, TF) are important tools for investigating galaxy kinematics. By comparing high-redshift galaxies with morphologically similar local galaxies we can study fundamental properties of the high-z systems, investigate the processes by which disk-formation occurs (e.g. Navarro \\& Steinmetz 2000; Mo \\& Mao 2000) and eventually understand galaxy formation. Internal kinematic information is crucial since it provides a direct estimate of the masses of galaxies and allows galaxy evolution to be traced directly by mass as distinct from light. To date, datasets on kinematics of high-redshift galaxies contain few galaxies and the results have been somewhat discrepant. For instance, Rix et al.~(1997), Simard \\& Pritchet (1998) find significant luminosity evolution for their low to intermediate (0.2$<$z$<$0.4) redshift galaxies while Vogt et al.~(1996; 1997) find evidence for only modest brightening for their (0.1$<$z$<$1) sample. Apart from measuring offsets from the local TF, examining differences in the slope of the TF is also crucial (e.g. Barton et al.~2001). For instance Ziegler et al.~(2002) report that the TF slope changes with redshift. An increase in the TF slope with redder passbands going from L$\\propto$v$^{3}$ in the optical to L$\\propto$v$^{4}$ in the near infrared has also been found (e.g. Verheijen~2001).   The discovery by COBE of an extragalactic infrared background with an intensity at least as large as that of the optical background suggests that a significant fraction of the star formation in the Universe takes place in dust enshrouded regions and is thus missed by optical surveys. With the advent of the Infrared Space Observatory (ISO,  Kessler et al.~1996) deep mid--IR surveys for distant galaxies have been successfully carried out for the first time. Operating in the 5 -- 18 $\\mu$m band sensitive to warm dust and PAH$/$UIB emission features, ISOCAM on board ISO was more than 1000 times more sensitive than IRAS and  thus had the potential to study infrared bright galaxies at redshifts beyond 0.5.  A number of cosmological surveys have been performed with ISOCAM especially in the LW3 filter (12 -- 18$\\mu$m) resulting in the detection of galaxies down to flux density limits of 100 $\\mu$Jy. We are engaged in a followup effort to investigate the nature of galaxies discovered through deep ISOCAM surveys. This new population of galaxies is at z$\\sim$1 and shows strong cosmological evolution (Elbaz et al.~1999, Franceschini et al.~2001). In Rigopoulou et al.~(2000), \\citetext{hereafter Paper I} we addressed the issue of the energetics by studying the properties of a sample of galaxies detected by ISOCAM observations of the Hubble Deep Field South (HDFS, hereafter ISOHDFS galaxies). Based on our low resolution near infrared (rest-frame R-band) spectra we concluded that ISOHDFS galaxies are powerful dusty starbursts, consistent with their large mid-IR luminosities. These observations indicate that this population traces an important phase of galaxy evolution. The next and most challenging step is to investigate their mass content and to compare it with the ongoing rate of star formation quantified in Paper I.  In this paper we present our results on spatially resolved rotation curves of four ISOHDFS galaxies selected from the sample of ISOCAM detections in the Hubble Deep Field South. We compare their properties to those of local galaxies and investigate the offset from the local TF. We find that one of the four galaxies studied is a {\\em \\bf{ massive spiral with an extraordinarily high mass of 10$^{12}$\\ M$_{\\odot}$}}. ",
        "conclusions": "We have presented first results from internal kinematic studies of four intermediate redshift (z = 0.6 -- 0.8) galaxies detected by ISOCAM in the HDFS. The kinematics are derived from high resolution spatially resolved H$_{\\alpha}$ velocity profiles. The galaxies studied, two large Sbc spirals and a double pair$/$interacting system, are representative of the morphologies of the entire ISOHDFS sample. Our results are summarized as follows:  ISOHDFS galaxies tend to be massive galaxies with masses of a few x 10$^{11}$M$_{\\odot}$. One of the galaxies studied, ISOHDFS 27, contains a mass of 10$^{12}$ M$_{\\odot}$ (within 20 kpc) substantially higher than the masses measured in other intermediate redshift galaxies and, comparable to the mass of giant ellipticals and hosts of luminous QSOs and radio galaxies. The presence of a central AGN in this galaxy is ruled out by three independent pieces of evidence: the extend$/$ morphology of the H$_{\\alpha}$ emission, the H$_{\\alpha}/$[NII] ratio of 3 (typical of starbursts) and the mid-infrared LW3(15$\\mu$)m$/$LW2(7$\\mu$m) ratio which again is typical of values found in local luminous IR galaxies. Based on the MIR luminosities and assuming L$_{FIR}$=10 L$_{MIR}$ we have calculated the mass-to-light ratio M$/$L$_{bol}$ of ISOHDFS galaxies. The M$/$L$_{bol}$ values are between the values found for local IR luminous star-forming galaxies and massive spirals.  We have compared the four ISOCAM galaxies to the local B and I-band Tully-Fisher relations. We have measured offsets of 1.6$\\pm$0.3 and 0.7$\\pm$0.3 mag from the B and I-band TF relations, respectively. Although our sample is very small to draw statistical conclusions it is evident that a clear offset from the local TF relations exist. We discuss the possible existence of a mass-dependent luminosity evolution where small mass (size) galaxies are more likely to undergo merging$/$interactions and thus show substantial luminosity evolution. In addition, recent episodes of star-forming activity, thus strong emission-line objects, are bound to show higher luminosity evolution than the more quiescent ones. Overall, we find that selection effects are very crucial to studies of the TF relation in intermediate$/$high redshift systems. We stress the uniqueness of our sample (an infrared selected one), when comparing it to other studies.  We have examined the global morphological properties of the ISOHDFS galaxies: 40\\% of ISOHDFS galaxies are double pair$/$interacting systems (akin to ISOHDFS 25), with the remaining being field disk-dominated spirals (similar to ISOHDFS 55). A further $\\sim$10\\% of the spirals are large early--type systems like ISOHDFS 27. Based on the current dynamical mass estimates and estimates of the baryonic masses for the majority of the ISOHDFS galaxies we find that ISOCAM galaxies are indeed massive systems. Our results imply that the IR background is dominated by a small number of massive and IR luminous star-forming galaxies near $<$z$> \\sim$0.7."
    },
    "0207/astro-ph0207382_arXiv.txt": {
        "abstract": "{ We present K-band echelle spectra of the cataclysmic variable SS~Cyg and the pre-cataclysmic variable V471 Tau in order to measure the strengths of the $^{12}$CO and $^{13}$CO bands at 2.3525 and 2.3448 $\\mu$m, respectively, and so perform the observational test of the common-envelope model of close binary star evolution proposed by \\citet{sarna95}. Although we find evidence of an absorption feature coincident with the expected wavelength of $^{13}$CO in both objects, we attribute it instead to a cluster of neutral atomic absorption features (primarily due to Ti\\,{\\small I}) possibly arising from star-spots on the surfaces of the rapidly rotating secondary stars in these systems, thereby rendering the test inconclusive. We present a modified observational test of common-envelope evolution, based on the observation of the $^{13}$CO bands at 2.3739 and 2.4037 $\\mu$m, which is insensitive to spectral contamination by star-spots. ",
        "introduction": "Well over half of all stars are believed to be binary or multiple systems, with about half of these, in turn, consisting of close binary systems where the two component stars are unable to complete their normal evolution without being influenced by the presence of the other (see \\citet{duquennoy91} and references therein).  The orbital separations of close binary systems containing at least one compact object -- such as the cataclysmic variables (CVs) and low-mass X-ray binaries (LMXBs) -- are significantly smaller than the radii of the stars which were the progenitors of the compact objects in these systems. This means that significant orbital shrinkage must have occurred, probably in a process known as common-envelope (CE) evolution. According to the CE model of close binary star evolution (in this case, as applied to CVs; \\citealt{paczynski76}), the more massive (primary) star fills its Roche lobe when it reaches its giant or asymptotic giant branch phase, while its lower mass (secondary) companion remains on the main sequence. Under these conditions, mass transfer to the secondary is dynamically unstable and occurs at such a high rate that the transferred material cannot be accreted by the secondary and so forms a CE, in which the binary is immersed. Through the action of drag forces, the main-sequence star spirals towards the core of the giant, generating luminosity which drives off the CE.  What remains is often called a post-common envelope binary (PCEB), with typical orbital periods of a few days.  These systems are thought to become CVs or LMXBs when magnetic braking or gravitational radiation extracts sufficient orbital angular momentum for the main-sequence secondary star to fill its Roche lobe (\\citealt{spruit83}; \\citealt{rappaport83}). The theory of CE evolution is therefore of fundamental importance in astrophysics, and is probably a key step in the production of some of the most exotic inhabitants of our Galaxy, including the binary radio pulsars, black-hole X-ray binaries and Type Ia supernovae. For a recent review of CE evolution, see \\citet{iben93}.  Although there is general agreement that most close binary systems have evolved through a CE phase, there is little direct evidence to support the CE model.  The best evidence to date for the reality of CE evolution comes from the observation of planetary nebulae with close binary nuclei (\\citealt{iben93}; \\citealt{livio96}; \\citealt{bond92}). There is no direct evidence, however, that whole classes of important objects such as LMXBs, CVs and their immediate precursors, the so-called pre-CVs (e.g. \\citealt{catalan94}), have in fact evolved through a CE phase.  As a result, \\citet{sarna95} proposed a direct observational test of CE evolution. The idea is based on the fact that the ratio of $^{12}$C/$^{13}$C decreases from a value of 84 in main-sequence stars like the Sun \\citep{harris87} to a value of about 17 in giants \\citep{harris88}, due to the different nuclear burning and mixing processes which occur in these stars. During the CE phase, the main-sequence secondary effectively exists within the atmosphere of the giant primary and will accrete material from it, thereby altering the $^{12}$C/$^{13}$C ratio from solar-like values towards giant-like values. By measuring the $^{12}$C/$^{13}$C ratio it is therefore possible to determine whether a binary has passed through a CE stage.  This test has already been performed by \\citet{dhillon95a}, who made a tentative detection of $^{13}$CO in the K-band spectrum of the pre-CV V471~Tau. To confirm this detection, we observed V471 Tau again, along with the CV SS Cyg. The results of these new, much higher quality observations are presented in this paper, together with a discussion of the recent results of \\citealt{catalan01} (who independently performed similar observations to the ones we describe here). ",
        "conclusions": "Our K-band echelle spectra of the CV SS Cyg and the pre-CV V471 Tau show evidence of a spectral feature at 2.3448 $\\mu$m. This is coincident with where we would expect to observe the $^{13}$CO (2--0) molecular band, which would imply that these objects have passed through a CE phase during their evolution. However, our spectra also show enhanced Na\\,{\\small I} absorption, which can be attributed to star-spots. These star-spots also contribute to a cluster of strong, neutral atomic absorption features around 2.3448 $\\mu$m, rendering our test of CE evolution inconclusive. We therefore propose a new test based on the measurement of the $^{13}$CO bands at 2.3739 and 2.4037 $\\mu$m, which we show would be uncontaminated by star-spots. We intend to perform this revised test in the near future."
    },
    "0207/astro-ph0207219_arXiv.txt": {
        "abstract": "s{ This paper surveys our current knowledge of the hard X-ray emission properties of old accreting neutron stars in low mass X-ray binaries. Hard X-ray components extending up to energies of a few hundred keV have been clearly detected in sources of both the Atoll and Z classes. The presence and characteristics of these hard components are discussed in relation to source properties and state. An overall anticorrelation between the fraction luminosity in hard X-rays and mass accretion rate is apparent over different sources spanning a large range of luminosities as well as individual source undergoing state changes. Evidence for a second, yet unknown, parameter controlling the hard X-ray emission is emerging. We draw a parallel with the spectral properties of  X-ray binaries hosting a stellar mass accreting black hole, and conclude that, at a merely phenomenological level, there appears to be a close analogy between the spectral properties of black hole candidates in their high and intermediate states and Z-sources. We briefly mention models that have been proposed for the hard X-ray emission of neutron star low mass X-ray binaries and comment on perspectives in the INTEGRAL era.} ",
        "introduction": "A variable hard component dominating the spectrum of Sco X-1 above $\\sim 40$~keV was detected as early as 1966 (Peterson \\& Jacobsen 1966; see also Riegler et al. 1970; Agrawal et al. 1971; Haymes et al. 1971). In other occasions the hard tail in Sco X-1 was not found, perhaps owing to pronounced variations (e.g., Miyamoto \\& Matsuoka 1977, and references therein; Soong \\& Rothschild 1983; Jain et al. 1984; Ubertini et al. 1992). Evidence for a hard component was also found in Cyg~X-2 (Peterson 1973) and GX~349+2 (Greenhill et al. 1979). These results received relatively little attention, probably because the nature of Sco~X-1-like and bright galactic bulge X-ray sources remained not understood or, at least, controversial until late seventies. On the contrary, that Cyg~X-1 hosts an accreting black hole candidate, BHC, had become clear as early as 1972 (Bolton 1972; Webster \\& Murdin 1972). The conspicuous hard X-ray emission of this source (see e.g.\\ Tanaka \\& Lewin 1995, and references therein) was therefore considered the prototypical \"hard spectrum\" of an accreting black hole and provided much of the observational basis for model development. Optically thick, geometrically thin accretion disk models that were developed in those years proved inadequate to explain the hard power-law like spectrum of Cyg~X-1, that extended without a break up to $\\sim 80-100$~keV in the hard state and was (likely) detected up to energies of $\\sim 1$~MeV in soft/intermediate states (e.g.\\ Liang \\& Nolan 1984). Standard accretion disk models were thus modified to include the presence of a hot inner disk region or corona, where unsaturated thermal Comptonisation of soft photons from the optically thick disk up to energies of many tens or hundreds keV takes place (Eardley, Lightman, \\& Shapiro 1975; Eardley \\& Lightman 1976; Galeev, Rosner \\& Vaiana 1979).  Renewed interest in the hard X-ray emission properties of neutron star low mass X-ray binaries was motivated by the SIGMA/GRANAT discovery of a spectral component extending up to energies of $\\sim 100-200$~keV in Terzan 2 (Barret et al. 1991), KS 1731-260 (Barret et al. 1992), SLX 1735-269 (Goldwurm et al. 1996) and Terzan 1 (Borrel et al. 1996). Unlike Sco~X-1 (and similar sources), the sources above have relatively low luminosity ($\\sim 10^{36}-10^{37}$ ergs/s) and emit type I X-ray bursts; moreover some of them are transients. ",
        "conclusions": ""
    },
    "0207/astro-ph0207505_arXiv.txt": {
        "abstract": "We consider the most general parametrization of flat topologically compact universes, complementing the work of Scannapieco, Levin and Silk to include non-trivial shapes. We find that modifications in shape of the fundamental domain will lead to distinct signatures in the anisotropy of the cosmic microwave radiation. We make a preliminary assessment of the effect on three statistics: the angular power spectrum, the distribution of identified ``circles'' on the surface of last scattering and the correlation function of antipodal points. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207396_arXiv.txt": {
        "abstract": "The archetypal model for the recently discovered dark energy component of the universe is based on the existence of a scalar field whose dynamical evolution comes down today to a non--vanishing cosmological constant. In the past -- before big--bang nucleosynthesis \\textcolor{red}{for that matter} -- that scalar field could have gone through a period of kination during which the universe has expanded at a much higher pace than what is currently postulated in the standard radiation dominated cosmology. I~examine here the consequences of such a period of kination on the relic abundance of neutralinos and find that the latter could be much higher -- by three orders of magnitude -- than what is estimated in the canonical derivation. I~shortly discuss the implications of this scenario for the dark matter candidates and their astrophysical signatures. ",
        "introduction": "\\label{sec:introduction}  \\vskip 0.1cm The \\textcolor{red}{recent WMAP} observations of the Cosmic Microwave Background (CMB) anisotropies \\cite{WMAP}, combined either with the determination of the relation between the distance of luminosity and the redshift of supernovae SNeIa \\cite{supernovae_omegaL}, or with the large scale structure (LSS) information from galaxy and cluster surveys \\cite{2dF}, give independent evidence for a cosmological average \\textcolor{red}{ matter density of $\\omegaM = 0.27 \\pm 0.04$} \\cite{WMAP}. This value may be compared to a baryon density of \\textcolor{red}{ $\\omegaB = 0.044 \\pm 0.004$} as indicated by nucleosynthesis \\cite{nucleosynthesis} and the relative heights of the first acoustic peaks in the CMB data. A significant fraction of the matter in the universe is dark and non--baryonic. The cosmological observations also point towards a flat universe \\textcolor{red}{with $\\Omega_{\\rm tot} = 1.02 \\pm 0.02$} and strongly favour the existence of a cosmological constant which contributes a fraction \\textcolor{red}{$\\omegaL = 0.73 \\pm 0.04$} to the closure density. The pressure--to--density ratio $w$ of that fluid is negative with a value of $w = - 1$ in the case of an exact cosmological constant. That $\\omegaL$ component is called dark energy as opposed to the $\\omegaM$ dark matter contribution.  \\vskip 0.1cm The nature of the astronomical dark matter is still unresolved insofar. The favorite candidate for the non--baryonic component is a weakly--interacting massive particle (WIMP). The so--called neutralino naturally arises in the framework of supersymmetric theories. Large efforts have been devoted in the past decade to pin down these evading species. New experimental techniques have been devised to look for the direct and indirect astrophysical signatures of the presence of neutralinos in our Milky--Way \\cite{gamma_neutralino_MW} as well as in extra--galactic systems \\cite{gamma_neutralino_M87}. The uncertainty on the theoretical estimates of the various signals has been considerably reduced. As an illustration, the energy spectrum of secondary spallation antiprotons -- the natural background to a putative neutralino--induced antiproton extra radiation -- is now well under control \\cite{secondary_antiprotons}. Another example of the level of sophistication which the theoretical investigations have reached is provided by the calculations of the neutralino relic abundance $\\omegachi$. The observation that this relic density -- depending on the numerous parameters of the model -- falls in the ballpark of the measured value for $\\omegaM$ has been a crucial argument in favor of supersymmetric particles as a viable option to non--baryonic dark matter. A large number of processes -- typically $\\sim$ 2000 -- are now taken into account and the corresponding diagrams are automatically generated and calculated with the help of numerical codes such as micrOMEGAs \\cite{micromegas}. Co--annihilations are taken into account and the thermal averaging $\\left< \\sigma v \\right>$ of the product of the velocity by the cross section is performed.  \\vskip 0.1cm Surprisingly enough, calculations of $\\omegachi$ are based on the assumption that the universe is dominated by radiation when neutralinos decouple from the primordial plasma and reach their relic density. This hypothesis is presumably correct as soon as primordial nucleosynthesis (BBN) sets in at a time of $\\sim$ 1 second. We have however little information on the earlier pre--BBN period, a crucial stage during which neutralinos freeze out. If the expansion of the universe is modified with respect to a pure radiation--dominated behaviour, the quenching of these species could be drastically modified. An increase in the expansion rate $H$ accelerates the  decoupling of neutralinos and translates into larger values for the relic density $\\omegachi$.  \\vskip 0.1cm Exploring the effects of a modified expansion history of the universe onto the relic abundance of neutralinos is no longer a mere academic exercise. Such an analysis has become mandatory inasmuch as a new and unexpected component -- the dark energy $\\omegaL$ -- has been discovered. The potential interplay between that component and its matter counterpart $\\omegaM$ is worth being explored and may have unexpected consequences. The archetypal model for the cosmological dark energy is the so--called quintessence \\textcolor{red}{\\cite{quintessence,steinhardt}} and relies on the existence of a neutral scalar field $\\Phi$ with Lagrangian density \\beq {\\cal L} \\; = \\; \\frac{1}{2} \\, g^{\\, \\mu \\nu} \\, \\partial_{\\mu}  \\Phi \\, \\partial_{\\nu}  \\Phi \\, - \\, V \\left( \\Phi \\right) \\;\\; . \\label{scalar_neutral} \\eeq Should the field $\\Phi$ be homogeneous and the metric be flat, the energy density may be expressed as \\beq \\rho_{\\Phi} \\equiv T^{0}_{\\; 0} \\; = \\; {\\displaystyle \\frac{\\dot{\\Phi}^{2}}{2}} \\, + \\, V \\left( \\Phi \\right) \\;\\; , \\label{energy_density} \\eeq whereas the pressure obtains from $T_{i j} \\equiv - g_{\\, i j} \\, P$ so that \\beq P_{\\Phi} \\; = \\; {\\displaystyle \\frac{\\dot{\\Phi}^{2}}{2}} \\, - \\, V \\left( \\Phi \\right) \\;\\; . \\label{pressure} \\eeq If the kinetic term ${\\dot{\\Phi}^{2}}/{2}$ is negligible with respect to the contribution of the potential $V \\left( \\Phi \\right)$, a pure cosmological constant with $w_{\\Phi} = P_{\\Phi} / \\rho_{\\Phi} = - 1$ is recovered since $\\rho_{\\Phi} = - \\, P_{\\Phi} = V \\left( \\Phi \\right)$. As indicated by cosmological observations, this is the case today. \\textcolor{red}{ But the field $\\Phi$ has been continuously rolling down. Should the kinetic term ${\\dot{\\Phi}^{2}}/{2}$ have dominated over the potential $V \\left( \\Phi \\right)$ in the early universe, a period of kination -- \\ie, domination by the kinetic energy of the field $\\Phi$ -- would have ensued with drastic effects on the expansion rate of the universe \\cite{joyce_kination}.}  \\vskip 0.1cm \\textcolor{red}{ In section~\\ref{sec:kination}, we briefly recall why a pure cosmological constant should be disregarded and replaced by a dynamical dark energy component in the form of a scalar field, the so--called quintessence whose salient features are presented. The existence of tracking solutions provides a natural solution to the problem of initial conditions. We also pay some attention to the difficulty of generating a kination--dominated expansion in the early universe together with a cosmological constant today \\cite{steinhardt}. We show that this difficulty may be circumvented depending on the potential $V \\left( \\Phi \\right)$ that drives the evolution of the scalar field and we propose examples where quintessence boosts the expansion rate in the past while it still accounts for $\\omegaL$ today.} Following a suggestion by \\cite{joyce_idea}, we investigate in section~\\ref{sec:decoupling} the effects of kination on the thermal decoupling of neutralinos and derive an approximate relation between their relic abundance $\\omegachi$ and their annihilation cross section in the presence of kination. Section~\\ref{sec:discussion} is devoted to a discussion of the consequences of this scenario on the astrophysical signatures of neutralino dark matter. ",
        "conclusions": "\\label{sec:discussion}  \\vskip 0.1cm We have shown that the neutralino relic abundance increases if a period of kination takes place during the freeze--out of the species. We derive now an estimate of the corresponding boost factor as a function of $\\etaphi$ and $m_{\\chi}$. If $\\tilde{a}$ dominates over $\\tilde{b} \\, \\xF$ in the expression of the annihilation cross section, we may even simplify further relation~(\\ref{OMH2_APRX}) in order to get \\beq \\Ochi \\; \\sim \\; \\left\\{ {\\displaystyle \\frac {2 \\times 10^{-27} \\; {\\rm cm^{3} \\; s^{-1}}}{\\tilde{a}}} \\right\\} \\;\\; \\eeq in the conventional radiation dominated cosmology. We have taken a value of $\\yF \\sim 20$ for the mass--to--decoupling temperature ratio in that case. On the contrary, if quintessence is the dominant form of energy with a large value for $\\etaphi$, the neutralino fossile abundance becomes \\beq \\Ochi \\; \\sim \\; {\\displaystyle \\frac{\\sqrt{\\etaphi} \\; m_{100}} {\\ln \\left( 2 \\, u \\right)}} \\, \\left\\{ {\\displaystyle \\frac {1.3 \\times 10^{-23} \\; {\\rm cm^{3} \\; s^{-1}}}{\\tilde{a}}} \\right\\} \\;\\; , \\eeq where $u = \\sqrt{\\alpha} \\, \\xF$ has already been defined. With a value of $\\yF \\sim 10$ -- see the previous section -- this implies \\beq u \\; \\simeq \\; 1.5 \\times 10^{4} \\; \\sqrt{\\etaphi} \\; m_{100} \\;\\; , \\eeq so that the logarithm yields a contribution $\\sim$ 10. The parameter $m_{100}$ denotes the neutralino mass in units of 100 GeV. A crude estimate of the relic abundance in this regime ensues \\beq \\Ochi \\; \\sim \\; \\sqrt{\\etaphi} \\; m_{100} \\, \\left\\{ {\\displaystyle \\frac {1-2 \\times 10^{-24} \\; {\\rm cm^{3} \\; s^{-1}}}{\\tilde{a}}} \\right\\} \\;\\; . \\eeq We derive a boost factor of $\\sim 10^{3} \\, \\sqrt{\\etaphi} \\; m_{100}$ with respect to the conventional cosmology. If now $\\tilde{b} \\, \\xF$ is the leading term as regards the annihilation cross section, we find that the relic abundance which is normally given by \\beq \\Ochi \\; \\sim \\; \\left\\{ {\\displaystyle \\frac {10^{-25} \\; {\\rm cm^{3} \\; s^{-1}}}{\\tilde{b}}} \\right\\} \\;\\; , \\eeq is increased to \\beq \\Ochi \\; \\sim \\; \\sqrt{\\etaphi} \\; m_{100} \\, \\left\\{ {\\displaystyle \\frac {1.2 \\times 10^{-22} \\; {\\rm cm^{3} \\; s^{-1}}}{\\tilde{b}}} \\right\\} \\;\\; . \\eeq in the presence of quintessence. The boost factor is still of the order of $\\sim 10^{3} \\, \\sqrt{\\etaphi} \\; m_{100}$. In our fiducial illustration, we actually obtained an increase of the neutralino relic abundance by a factor of $3,000$ with $m_{100} = 2.5$ and $\\etaphi = 1$ in good agreement with the bench mark value which has been derived here.  \\vskip 0.1cm The increase of $\\Ochi$ with $\\etaphi$ has interesting consequences and brings up new perspectives as regards neutralino dark matter. To commence, the various avatars of the minimal or non--minimal supersymmetric extensions of the standard model start to be constrained -- should R parity be conserved -- by the accelerator data on the one hand side and by the requirement that the neutralino relic abundance should not overclose the universe or even exceed the observed value of $\\omegaM \\, h^{2} = 0.135 \\pm 0.009$ \\cite{ellis}. If a period of kination takes place in the pre--BBN period, the various SUSY configurations in the $(\\Ochi , m_{\\chi})$ plane that are so far allowed are shifted upwards with the consequence of becoming forbidden. Exploring in greater detail this question is a worthwhile project.  \\vskip 0.1cm We already anticipate that configurations with a very small relic density -- for instance those for which poles dominate in the annihilation mechanism -- would become cosmologically attractive if $\\etaphi$ is large enough. The difference with the conventional cosmology lies in the significant enhancement of the annihilation cross section of neutralino dark matter candidates. At fixed $\\Ochi$, notice that $\\Sa v$ increases precisely by the same factor of $\\sim 10^{3} \\, \\sqrt{\\etaphi} \\; m_{100}$ which we have derived above. This means a general enhancement of the various indirect signatures for supersymmetric dark matter. If neutralinos dominate the mass budget of the Milky Way halo, they should still annihilate today and produce gamma--rays, antiprotons and positrons which may be detected through the corresponding distortions in the various energy spectra. As a matter of fact, the recent HEAT experiment has confirmed \\cite{heat} an excess around 8 GeV in the positron spectrum of cosmic rays. A large boost factor of $\\sim 10^{3}$ -- $10^{4}$ in most of the supersymmetric parameter space is needed to explain that excess in terms of a homogeneous distribution of annihilating galactic neutralinos. A certain degree of clumpiness is actually expected in most of the numerical simulations but even in the extreme case of \\cite{moore_clumps}, it does not exceed a few hundreds. A period of kination in the early universe could provide an alternate explanation for that boost factor.  \\vskip 0.1cm Another potential consequence of quintessence is the rehabilitation of a fourth generation heavy neutrino in the realm of the dark matter candidates. In the conventional cosmology, a 100 GeV stable neutrino provides today a contribution of $\\sim 10^{-4}$ to the closure density. Once again, kination at the time of decoupling would enhance that relic density and make that species cosmologically relevant.  \\vskip 0.1cm Notice finally that scenarios with extra--dimensions have the same effect as kination. The expansion rate is also increased and may even evolve as $T^{4}$ -- to be compared to a $T^{3}$ behaviour in our case and to a $T^{2}$ dependence in the conventional radiation dominated universe. Implications of such scenarios on neutralino dark matter should be investigated. In the case of a low reheating temperature at the end of inflation, neutralinos are not thermally produced. Depending on the details of the scenario -- decay of an inflaton field \\cite{giudice} or on the contrary decay of moduli fields \\cite{khalil} -- the relic abundance is decreased or increased.  \\vskip 0.1cm A key ingredient of our study is the contribution $\\etaphi$ of the quintessential scalar field to the overall energy density at the onset of BBN. A detailed analysis of the light element yields in the presence of kination \\cite{melchiorri} is mandatory at that stage in order to explore a promising scenario or to derive constraints on $\\etaphi$. The existence of dark energy opens up an exciting line of research and the study of its implications on the astronomical dark matter problem will certainly bring surprising results."
    },
    "0207/astro-ph0207169_arXiv.txt": {
        "abstract": "We analyse the structural and dynamical properties of a sample of 324 nearby elliptical and dwarf elliptical galaxies observed during an extensive NIR survey in H-band (1.65\\micron). The Fundamental Plane (FP) is determined and a significant tilt is assessed. The origins of such a tilt are investigated by means of a spherically symmetric, isotropic pressure supported dynamical model relying on the observed surface brightness profiles. The systematic variation of the shape coefficient converting the measured central velocity dispersion $\\sigma_0$ into the virial rms velocity $\\sigma_{rms}$ is found to be the main cause of the tilt, due to aperture effects. Moreover the ratio between the dynamical mass $M_{dyn}$ and the total H-band luminosity $L_H$ turns out to be roughly constant along the luminosity sequence of ellipticals: H-band luminosity is therefore a reliable and cheap estimator of the dynamical mass of the Es. ",
        "introduction": "The assumption that normal galaxies are in dynamical equilibrium implies that, for a given type of dynamics (i.e. rotation or pressure supported), the dynamical status of the system is strongly related to its mass distribution. Using the further assumption that the mass inside the ``observable'' radius of a galaxy is traced by the (stellar) light, the dynamical parameters should be, at least in first approximation, determined by the structural parameters describing the light distribution. In the case of elliptical galaxies, these systems have been proved to be mainly pressure supported and the fundamental dynamical parameter is the central velocity dispersion $\\sigma_0$. The simplest way of describing the light distribution of a galaxy is to measure the half light (or effective) radius \\re and the average surface brightness $I_e$ inside \\re. If the elliptical galaxies formed a homologous family, both from the structural and the dynamical point of view, and the $M/L$ ratio were constant, the virial theorem would imply a linear relation between the logarithm of the three parameters given by: \\begin{equation} \\mathrm{Log}~R_e=2~\\mathrm{Log}~ \\sigma_0 - \\mathrm{Log}~ I_e + k \\end{equation} A linear relation has been actually found \\citep[the Fundamental Plane, FP,][]{ddFP,dresslerFP}, but the coefficients differ significantly from those predicted by the virial theorem (this is known as the ``tilt'' of the FP). This implies the non-constancy of $M/L$ and/or the breaking of the homology. However the existence of the FP as a tight relation requires that the variations of $M/L$ or the breaking of the homology happen in a very systematic way.\\\\ Most of the early studies devoted to this problem assumed homology and concluded that a systematic increase of $M/L$ with luminosity is needed both in optical \\citep[see e.g.][]{k3} and in NIR pass-bands \\citep[see e.g.][]{pahre98b}. However, many studies, in which the surface brightness profiles are analysed both in optical \\citep[see e.g.][]{djorg_etal85,1985nagp.meet..257D, 1986ApJS...60..603S,1987ApJS...64..643S,1987MNRAS.226..747J,1988ApJS...68..173D, 1988AJ.....96..487C} and in NIR pass-bands \\citep[see e.g.][]{dEvirgo,c31marco}, have found systematic variations of the profile shape and concentration index (i.e. the structural parameter quantifying how much the light distribution is centrally peaked\\footnote{Many definitions of concentration index were proposed in literature; the definition hereafter adopted is $c_{\\mathrm 31}=r_{75}/r_{25}$, the ratio between the radii enclosing 75\\% and 25\\% of the total luminosity of the galaxy \\citep[see][]{dEvirgo, c31marco}}) among ellipticals, implying a breaking of the homology. In pioneering work \\cite{prugniel97} and \\cite{graham97} produced evidence in favour of a strong influence of structural non-homology on the tilt of the FP. \\cite{busarello97} concluded that most of the tilt could be accounted for by dynamical non-homology, although spatial non-homology and stellar population effects can give significant contributions as well. More recently, \\cite{Bertin02} showed how the departure from spatial homology contributes to the tilt of the FP in the B band for a small sample of nearby E galaxies imaged with high S/N.\\\\ The claim by \\cite{FPmarco}, that the amount of tilt is significantly lower in the NIR bands than in the optical ones, implies, however, a significant role of different stellar populations (age, metallicity) in determining the M/L and the structural parameters. NIR pass-bands are the most suitable for studying the structural parameters because they trace much better than visible pass-bands the bulk of the luminous mass of the galaxies which sits in old stellar populations. Moreover they are less sensitive to dust obscuration and line blanketing, and in turn this reduces the age/metallicity effects on the $M/L$ ratio. We have performed a systematic and extensive investigation of the structural properties in the H (K') band for a sample of nearby galaxies, covering all the morphological types and extending from high to low luminosities, i.e. to the dwarf regime \\citep[see][~and references therein]{paperV,dEvirgo}. Based on these surface brightness profiles, \\cite{c31marco} demonstrate a systematic relationship between the concentration index $c_{31}$ and the total H-band luminosity which is almost completely independent from the eye-ball morphological classification: from this relationship the structural non-homology of the elliptical family can be inferred. The structural and dynamical properties of the whole sample of surveyed galaxies are analysed by \\cite{pieriniKspace} adopting the $\\kappa$-space formalism \\citep{k3}.\\\\ In this paper we present a study of the relationships between the structural H-band (1.65\\micron) properties of 324 elliptical and dwarf elliptical galaxies. The relationships with dynamical parameters are analysed for a subsample of 135 galaxies ranging from the highest luminosities to the dwarf regime and extending the original sample of 73 galaxies studied by \\cite{FPmarco}. The sample is described in Sec. \\ref{sample}. In Sec. \\ref{kormendy_sect} we study the ``Kormendy relation'' between the effective surface brightness \\mue and the effective radius \\re. In Sec. \\ref{FP_sect} a derivation of the H-band FP is presented using different fitting methods, with an analysis of the contribution to the FP tilt due to $M/L$ and homology breaking is performed in Sec. \\ref{tilt_sect}, using a simple dynamical model relying on the measured surface brightness profiles. The method, independently developed by \\cite{Zibetti_tesi}, is very similar to that of \\cite{Bertin02}. It should be stressed, however, that we analysed a 10 times larger sample using NIR photometric data, as opposed to the B band data analysed by \\cite{Bertin02}. A brief discussion and the conclusions of this work are given in Sec. \\ref{discussion} and \\ref{conclusions}. ",
        "conclusions": "We have determined the Fundamental Plane of elliptical galaxies in nearby rich clusters (mainly Virgo and Coma) in NIR H pass-band. Our result may be written: $\\mathrm{Log}~R_e=(1.38\\pm0.1)\\mathrm{Log}~\\sigma_0 -(0.88\\pm0.07)\\mathrm{Log}~I_e + 5.47$\\\\ The relation is tight, showing a dispersion of $0.14~dex$ in \\re, while typical errors are $0.09~dex$.\\\\ The origins of the tilt of the fitted plane with respect to the virial predictions for a homologous family of galaxies have been investigated by means of a simple dynamical model. Spherical symmetry, hydrodynamical equilibrium and isotropic velocity dispersion are assumed. Constant $M/L$ within each galaxy has been assumed as a free parameter and used to calculate density and velocity dispersion profiles. $M/L$ was then adjusted in order to match predicted values of $\\sigma_0$ with the measured ones.\\\\ The obtained values of $M/L$ do not show any dependence on the total luminosity. Systematic variations of $M/L$ are then ruled out as the main cause of the tilt of the Fundamental Plane. We showed the ratio between the rms velocity and the central velocity dispersion is systematically varying as a function of \\re, mainly due to the different slit aperture relative to \\re ~in the spectroscopic measurement of $\\sigma_0$. This variation is responsible for most of the amount of the tilt.\\\\ The constancy of $M/L$ makes the H-band total luminosity a reliable and cheap first-order estimator for the dynamical mass of elliptical galaxies. This matches the analogous claim by \\cite{gavrel} for disk galaxies.\\\\ Applying the dynamical model to datasets extended to low-luminosity and dwarf elliptical galaxies, that will become available in the future, will provide crucial informations in order to determine whether the hypotheses of constant $M/L$ and pure isotropic pressure support can describe (at least in first approximation) the whole family of ellipticals. The structural continuity between the giant- and the dwarf- regime is shown by the distribution of galaxies in the \\mue-\\re ~plane. The existence of an upper-limit to the effective surface brightness has been found which follows the classical Kormendy relation for the giants.\\\\"
    },
    "0207/astro-ph0207443_arXiv.txt": {
        "abstract": "The RXTE satellite observed the Coma cluster for $\\sim$177\\,ks during November and December 2000, a second observation motivated by the intriguing results from the first $\\sim$87\\,ks observation in 1996. Analysis of the new dataset confirms that thermal emission from isothermal gas does not provide a good fit to the spectral distribution of the emission from the inner 1$^o$ radial region. While the observed spectrum may be fit by emission from gas with a substantial temperature gradient, it is more likely that the emission includes also a secondary non-thermal component. If so, non-thermal emission comprises $\\sim 8\\%$ of the total 4--20 keV flux. Interpreting this emission as due to Compton scattering of relativistic electrons (which produce the known extended radio emission) by the cosmic microwave background radiation, we determine that the {\\it mean, volume-averaged} magnetic field in the central region of Coma is $B \\sim 0.1-0.3$ $\\mu G$. ",
        "introduction": "X-ray spectra of clusters of galaxies have long been expected to show structure beyond that of a single temperature thermal model, mainly due to non-isothermality of intracluster (IC) gas in the outer cluster region. In addition, non-thermal (NT) X-ray emission in clusters was predicted (\\eg, Rephaeli 1977) from Compton scattering of relativistic electrons by the Cosmic Microwave Background (CMB) radiation. There is at least some observational evidence for radial variation of the gas temperature in a few clusters (\\eg, Markevitch 1996, Honda \\ea 1996, Donnelly \\ea 1999, Watanabe \\ea 1999). In the Coma cluster, recent XMM measurements indicate (Arnaud \\ea 2001) that the temperature is remarkably constant within the central region where temperature variation was previously deduced from ASCA measurements. After a long search (for a recent review, see Rephaeli 2001), NT X-ray emission seems to have finally been measured in Coma (Rephaeli, Gruber \\& Blanco 1999, hereafter RGB, Fusco-Femiano \\ea 1999), A2256 (Fusco-Femiano 2000), A2319 (Gruber \\& Rephaeli 2002), and perhaps also in A2199 (Kaastra \\ea 2000).  Appreciable deviation from isothermality may have significant impact on modeling the structure and evolution of IC gas, and on use of the gas as a probe to determine the total cluster mass (assuming hydrostatic equilibrium). The exact gas density and temperature profiles are also very much needed in analysis of measurements of the Sunyaev-Zeldovich (S-Z) effect and its use as a cosmological probe. There clearly is strong motivation for a more realistic characterization of cluster X-ray spectra for an improved description of IC gas, and the study of NT phenomena in clusters.  We have previously analyzed $\\sim 87$ and $\\sim 160$ ks RXTE measurements of the Coma cluster and A2319, respectively, in order to search for NT emission from these clusters which have well documented extended regions of radio emission. Analyses of these measurements yielded strong evidence for a second spectral component in both clusters. While the second component could possibly indicate a temperature variation across the cluster, we have concluded (RGB, Gruber \\& Rephaeli 2002) that the deduced spectral parameters are more consistent with power-law emission. NT emission in Coma seems to have been detected directly -- in the 25-80 keV range -- by BeppoSAX (Fusco-Femiano \\ea 1999). This provided further impetus to propose a longer observation of this cluster with RXTE. Here we briefly report the results from a joint analysis of these and the previous RXTE measurements, with a total integration time of $\\sim 264$ ks, and discuss some of their direct implications. ",
        "conclusions": "The consistent results from the analysis of the second (year 2000) and combined (years 1996 \\& 2000) datasets further substantiate the reality of the detection of a second component in the Coma spectrum. If thermal, this component could help determine the thermal structure of IC gas, which in turn would have important implications for the use of clusters as cosmological probes in general, and the most extensively researched nearby rich cluster in particular. If this component is NT emission from a population of \\rel electrons, then a new dimension for the study of NT IC phenomena will have been opened. Since the combined $\\sim 264$ ks RXTE observations have not yielded direct detection of power-law emission at energies $> 30$ keV, we have to invoke other observational and theoretical considerations in order to identify more uniquely the origin of the extra emission we have deduced.  It is unrealistic to expect that IC gas is fully isothermal outside the central region of the cluster. A more likely behavior is at least some decrease of the temperature outside a central $\\sim 2-3$ core radii region. Indeed, previous ASCA measurements of Coma (Honda \\ea 1996) and other clusters (e.g., Markevitch 1996) seem to have shown some deviations from isothermality. In Coma, the Honda \\ea (1996) analysis of measurements from 14 different pointings of an area extending $\\sim 1^{o}$ from the center indicated that the IC gas temperature varied by $\\pm 50\\%$ -- with respect to the overall mean value of $\\sim 8$ keV -- in two azimuthal regions $40'$ from the center. However, these two regions cover only a small part of the projected area of the cluster, and since the temperature is higher than the mean in one region, while it is lower in the other, the overall change of the spectral flux, as compared with that from an isothermal gas at the mean cluster temperature, is negligible. More recent results from high spectral and spatial resolution measurements of Coma with XMM it was concluded that the temperature is constant in the central $\\sim 10'$ radial region, with a best fit value of $8.2 \\pm 0.1$ keV (Arnaud \\ea 2001). While the fits to the XMM data do not seem to have included two-temperature models, it is clear that an appreciable emission at a significantly different temperature than this mean value would have been deduced from the XMM measurements through a larger variance, if not in the form of a systematic temperature gradient.  The results of our RXTE analysis yield a statistically most probable temperature combination with a prohibitively high value for the second component, $kT_2 \\simeq 37.1$. At its lowest boundary, the 90\\% contour region in the ($kT_{1}$, $kT_{2}$) plane does include the more acceptable values $kT_{1}=5.5$ keV, and $kT_{2}=9$, but this only if the respective 4--20 keV fractional fluxes of these two components are 24\\% and 76\\%. It is quite unlikely that about a quarter of the flux could come from a component with a significantly lower temperature than the mean value deduced by virtually all previous X-ray satellites. In particular, such a component would have been detected in the high spatially resolved measurements with ROSAT and XMM.  To assess the possibility that a two-temperature gas model is just a simplified representation of a more realistic continuous temperature distribution, we have repeated the following simple procedure we employed in our analysis of the first RXTE observations (RGB): Assuming a polytropic gas temperature profile of the form $T(r) \\propto n(r)^{\\gamma -1}$, with the familiar $\\beta$ density profile for the gas density, $n(r) \\propto (1+r^{2}/r_{c}^2)^{-3\\beta /2}$, where $r_c$ is the core radius, we calculated the integrated flux and the mean emissivity-weighted temperatures as functions of $\\gamma$, $\\beta$, and $r$. These quantities were then calculated in the regions $[0, r]$ and $[r, R_0]$ by convolving over the triangular response of the PCA with $R_0 \\simeq 58'$. From ROSAT observations, $r_c \\sim 10.0'$, and $\\beta \\simeq 0.70 \\pm 0.05$ (Mohr \\ea 1999). We sought the range of values of $r$, $\\beta$, and $\\gamma$ for which the two mean emissivity-weighted temperatures and respective fluxes from these regions are closest to the values deduced from our spectral analysis in Section 3. The results of these calculations indicate that for $0.5 \\leq \\beta \\leq 0.9$ and $1 \\leq \\gamma \\leq 5/3$, there is no acceptable polytropic configuration that matches the observationally deduced values of the temperatures and fractional fluxes. For low values of $\\gamma$ the temperature gradient is too shallow, while for high values the implied central temperature is unrealistically high. This simple plausibility check suggests that the two thermal components model is somewhat inconsistent with the RXTE results. However, the gas distribution may be more complicated than considered here, so that a temperature structure as implied here cannot be altogether ruled out.  Of particular interest is the somewhat more likely possibility that the secondary spectral component is NT. Since emission from an AGN in the FOV is not likely (see details in RGB), it is natural to consider that this emission is due to Compton scattering of \\rel electrons whose presence in Coma is directly inferred from many measurements of spatially extended region of radio emission (\\eg Kim \\ea 1990, Giovannini \\ea 1993). From the measured radio spectral index, $1.34\\pm 0.1$, it readily follows that the predicted power-law (photon) flux from Compton scattering of these electrons by the CMB has an index $2.34 \\pm 0.1$, a value which is quite consistent with what we have inferred, $2.1 \\pm 0.5$ (all errors are at 90\\% confidence). With the measured mean radio flux of $0.72 \\pm 0.21$ Jy at 1 GHz, the power-law X-ray flux deduced here, and the {\\it assumption} that the spatial factors in the theoretical expressions for the two fluxes are roughly equal, we can easily compute (see more details in RGB) the mean volume-averaged value of the magnetic field, $B_{rx}$. Taking into account the full 90\\% range of values of the radio flux, radio index, and the power-law X-ray flux, we get $B_{rx} \\simeq 0.1 - 0.3 \\mu$G. This range of values for $B_{rx}$ is consistent with our previous estimate (RGB), and the range (0.14 - 0.25 \\,$\\mu$G) deduced by Fusco-Femiano \\ea (1999).   Since we have assumed that the spatial factors in the theoretical expressions for the radio and NT X-ray fluxes are roughly equal, it follows that the mean value of the deduced magnetic field is independent of the source size and distance. To determine the \\rel energy density we do have to specify the radius of the emitting region. Scaling to the observed radius of the diffuse radio emission, $R \\sim 20'$, and integrating the electron energy distribution over energies in the observed radio and X-ray bands, we obtain $\\rho_{e} \\simeq (8 \\pm 3)\\times 10^{-14} (R/20')^{-3}$ erg\\,cm$^{-3}$; a distance of $139$ Mpc (with $H_0 = 50$ km\\,s$^{-1}$\\,Mpc$^{-1}$) was used. Based on the high Galactic proton to electron energy density ratio of cosmic rays, it can be conjectured that the energetic proton energy density is considerably higher than this value.  The strength of IC magnetic field can also be estimated from Faraday rotation measurements of background radio sources seen through clusters, yielding a different mean field value, $B_{fr}$. Analyses of such measurements usually yield field values that are a few $\\mu$G (see, \\eg, Clarke, Kronberg, and B\\\"ohringer 2001, and the review by Carilli \\& Taylor 2002). Clearly, the mean strength of IC fields has direct implications on the range of electron energies that are deduced from radio measurements, and therefore on the electron (synchrotron and Compton) energy loss times. Higher electron energies imply shorter energy loss times, with possibly important ramifications for \\rel electron models (\\eg, Rephaeli 1979, Sarazin 1999, Ensslin \\ea 1999, Brunetti \\ea 2001, Petrosian 2001).  Much has been written about the apparent discrepancy between deduced values of $B_{rx}$ and $B_{fr}$. Indeed, it is sometimes claimed that this discrepancy makes Compton interpretation of cluster power-law X-ray emission untenable. However, $B_{rx}$ and $B_{fr}$ are actually quite different measures of the field: Whereas the former is essentially a volume average of the \\rel electron density and (roughly) the square of the field, the latter is an average of the product of the line of sight component of the field and gas density. All these quantities vary considerably across the cluster; in addition, the field is very likely tangled, with a wide range of coherence scales which can only be roughly estimated. These make the determination of the field by both methods considerably uncertain. Thus, the unsatisfactory observational status (stemming mainly from lack of spatial information) and the intrinsic difference between $B_{rx}$ and $B_{fr}$, make it clear that these two measures of the field cannot be simply compared. Even ignoring the large observational and systematic uncertainties, the different spatial dependences of the fields, \\rel electron density, and thermal electron density, already imply that $B_{rx}$ and $B_{fr}$ will in general be quite different. This was specifically shown by Goldshmidt \\& Rephaeli (1993) in the context of reasonable assumptions for the field morphology, and the known range of IC gas density profiles. It was found that $B_{rx}$ is typically expected to be smaller than $B_{fr}$. Various statistical and physical uncertainties in the Faraday rotation measurements, and their impact on deduced values of IC fields, were investigated recently by Newman, Newman \\& Rephaeli (2002); their findings further strengthen the assessment that a simple minded comparison of values of $B_{rx}$ and $B_{fr}$ is meaningless, and that it is quite premature to draw definite conclusions from the apparent discrepancy between values deduced by these very different methods to measure IC magnetic fields.  As we have mentioned in the previous section, the excess EUV emission in Coma could also be NT, and based on the similar morphologies of the EUV emission and low frequency radio emission, Bowyer \\& Berghofer (1998) interpreted this emission as Compton scattering of the CMB by a population of low energy electrons. They adopted a value for the power-law index which is somewhat lower than the value used here, but deduced a similar value ($\\sim 0.2\\, \\mu$G) for the mean magnetic field. More recently, Tsay, Hwang \\& Bowyer (2002) have explored whether a Compton origin for the observed EUV excess can be maintained even if the field is as high (few $\\mu$G) as is currently deduced from Faraday rotation measurements. They conclude that this is possible within a limited class of lower energy ($\\sim 100$ MeV) NT electron models, and that in this case a different (second) population of \\rel electrons is required to explain the measurements of NT X-ray emission by RXTE and BeppoSAX. Note that significant IC density of sub-relativistic electrons could in principle also produce high energy X-ray emission by NT bremsstrahlung (Kaastra \\ea 1998, Sarazin \\& Kempner 2000). However, the properly normalized contribution of such electrons to the power-law emission deduced here from the RXTE measurements is too small (Shimon \\& Rephaeli 2002) to affect our estimated value of the magnetic field.  Further evidence for the NT nature of the second spectral component in Coma could possibly come from the scheduled 500 ks observation of this cluster with with IBIS imager aboard the INTEGRAL satellite. The moderate $\\sim 12'$ spatial resolution of IBIS can potentially yield crucial information about the location and size of high energy NT X-ray emission."
    },
    "0207/astro-ph0207325_arXiv.txt": {
        "abstract": "Using a consistent set of models, we parameterized the \\xray\\ spectra of all accreting pulsars in the \\RXTE\\ database which exhibit \\CRSFS\\ (\\crsfs, or cyclotron lines). These sources in our sample are Her X-1, 4U 0115+63, Cen X-3, 4U 1626-67, XTE J1946-274, Vela X-1, 4U 1907+09, 4U 1538-52, GX 301-2, and 4U 0352+309 (X Per). We searched for correlations among the spectral parameters, concentrating on how the cyclotron line energy relates to the continuum and therefore how the neutron star $B$-field influences the X-Ray emission. As expected, we found a correlation between the \\crsf\\ energy and the spectral cutoff energy. However, with our consistent set of fits we found that the relationship is more complex than what has been reported previously. Also, we found that not only does the width of the cyclotron line correlate with the energy (as suggested by theory), but that the width scaled by the energy correlates with the depth of the feature. We discuss the implications of these results, including the possibility that accretion directly affects the relative alignment of the neutron star spin and dipole axes. Lastly, we comment on the current state of fitting phenomenological models to spectra in the \\rxte/\\sax\\ era and the need for better theoretical models of the \\xray\\ continua of accreting pulsars. ",
        "introduction": "\\label{sec:intro}  Accretion powered \\xray\\ pulsars \\citep{whi83,nag89,bil97} provide a unique laboratory for the study of matter in extremes of temperature and magnetic as well as gravitational fields. After more than two decades of research, however, there is still no compelling model for the generation of the hard \\xray\\ spectrum in these objects. This reflects the difficulties and complexities of radiative transport and magnetohydrodynamics in the environment found at the neutron star magnetic polar caps.  The observed hard \\xray\\ emission from these objects originates primarily from one or two ``hotspots'' found at the neutron star magnetic poles. Due to their large magnetic fields ($B\\gtrsim10^{12}$\\,G), material accreted from a nearby companion couples to the neutron star $B$-field at several hundred neutron star radii. The material is then channeled onto the neutron star surface, forming accretion structures at the two poles. It is the combination of the beaming properties of these structures with the rotation of the star that gives rise to the pulsed emission seen by a distant observer. By analyzing the rotation-averaged spectral properties of these hotspots, we hope to improve our understanding of the physical conditions and properties of the emission regions of these accretion structures. This is also a step towards interpreting and understanding pulsar spectra as a function of neutron star rotation phase (pulse phase resolved spectroscopy).  For the analysis presented here, we focus on the effects of the magnetic field on the resulting hard \\xray\\ spectrum. By using sources with known magnetic field strengths (from the measurement of cyclotron features) we remove one uncertainty from the class analysis. Our results, however, should also extrapolate to accreting pulsars with unknown $B$-field strengths. Another motivation was to provide an observational base for theoretical investigations into the production of the pulsar hard \\xray\\ continuum, as well as to guide future calculations and simulations.  The list of our sources, along with some of their properties, is given in Table~\\ref{table:systems}.  To perform this analysis, we used sources where there was a direct measurement of the pulsar $B$-field using \\CRSFs\\ (\\crsfs), also commonly referred to as ``cyclotron lines.'' These line-like spectral features arise due to the resonant scattering of photons by electrons whose energies are quantized into Landau levels by the magnetic field \\citep{meszaros92}. The fundamental energy where these features appear is given by \\begin{equation}\\label{eq:crsfeq} \\crsfeq{}\\rm{\\,keV} \\end{equation} where $B$ is the magnetic field (in Gauss) in the scattering region, and $z$ is the gravitational redshift. The quantized energy levels of the electrons are to first order harmonically spaced, with features at 2\\ecy{}, 3\\ecy{}, etc. both predicted and, in some sources, observed \\citep[e.g.][]{hei99comp0115,san99,cus98}. At sufficiently high magnetic fields, relativistic effects can introduce a slight anharmonicity in the rest frame resonant photon energies. In these cases the cyclotron energies are given by \\citep{har91} \\begin{equation}\\label{eq:revcrsfeq} \\revcrsfeq{} \\end{equation} where $B'=B (\\hbar e)/(m_{e}^{2}c^{3})$ is the magnetic field scaled to the QED field scale, $n$ is the harmonic number, and $\\theta$ is the angle of propagation of the photon relative to the magnetic field. Since these energies depend on the angle $\\theta$, the emergent spectral features are influenced heavily by the spatial distribution of electrons in the scattering region.  The dependence of $\\theta$ in Eq.~\\ref{eq:revcrsfeq} indicates that, even at nonrelativistic energies, the \\crsf\\ energy is not the only source of information about the scattering region. From the Monte Carlo simulations of \\citet{ara99} and \\citet{ise98,ise98b}, it is found that the shape of the fundamental can be quite complex, and in general depends heavily on the details of the emission and scattering geometries, as well as the physical parameters such as the electron temperature and density in the scattering region. These features are also sometimes observed to vary as a function of rotation phase of the star \\citep[e.g.][]{bur00,biff99,cla90,soo90a,vog82}, allowing for the detailed study of a single accretion structure using pulse phase resolved spectroscopy.  In this paper we summarize spectral fits to ten accreting pulsars, and present observational evidence for the effect of the $B$-fields on the underlying hard \\xray\\ continua of these pulsars. This was part of a larger analysis of \\rxte\\ archival data that encompassed 25 accreting pulsars in total. These ten (see Table~\\ref{table:systems}) were selected due to the fact that their spectra exhibited \\crsfs. They represent a complete sample of pulsars with \\crsfs\\ in the \\rxte\\ database. The remaining sources, the ones without detectable cyclotron lines, will be discussed in a future publication.  In \\S\\ref{sec:rxte} we discuss the \\RXTE\\ (\\rxte) satellite. In \\S\\ref{sec:method} we present a summary of the methodology and spectral models we used. In \\S~\\ref{sec:obs} we discuss the 10 pulsars in our sample, along with fits to their \\rxte\\ spectra. In \\S\\ref{sec:fits} the results of the fitting are presented, along with a discussion of how correlations were found and the checks that were done using Monte Carlo simulations. In \\S\\ref{subsec:results} we discuss our findings and their physical implications. Finally, in \\S\\ref{sec:summary} we conclude with a brief summary of our primary results and discoveries. ",
        "conclusions": "\\label{sec:summary}  There are three principle results of this class analysis of the spectra and cyclotron features of accreting \\xray\\ pulsars. The first two results involve correlations among the shape parameters of the \\crsfs\\ themselves. We find that the observed \\crsf\\ widths are roughly proportional to their energy. If the widths are primarily due to Doppler broadening, then this implies a viewing angle selection bias in finding sources that exhibit cyclotron features. Since 6 of the 10 sources discussed here are in systems that are viewed nearly edge on, a selection bias on viewing angles further suggests a preferred offset angle between the dipole and spin axes of the neutron star.  The next result is that, at least for the fundamental features, deeper \\crsfs\\ are also broader, even when scaled by the centroid energy. This is difficult to understand simply in terms of the relativistic cross sections alone, which are in the \\emph{opposite} sense. This implies that other effects, such as photon spawning or the non-isotropic angular redistribution of photons, are important and should continue to be considered in theoretical efforts.  Lastly, we find a correlation between the magnetic field strength and the spectral cutoff energy. The existence of a correlation indicates that the observed spectral break is either a magnetic effect, or perhaps is tied to the magnetic field through some intermediate quantity. In the correlation itself there is a departure from the power-law observed by \\citet{mak99}, either a roll over or break in the slope, near a cyclotron resonance energy of 35\\,keV. This might be an indication that the break is indeed tied to another quantity, such as the electron temperature, that is then tied to the magnetic field up to a saturation point.  We also discussed a departure in observed pulsar spectra from the standard pulsar continuum shape. While the standard shape is still applicable over most of the hard \\xray\\ band, modern satellites have shown that it is inadequate near 10\\,keV. This highlights the need for theoretical work in how the \\xray\\ continuum in these pulsars is formed, and when fitting spectra a departure away from the purely phenomenological models currently being used."
    },
    "0207/astro-ph0207113_arXiv.txt": {
        "abstract": "The  interpretation  of the old, cool  white  dwarfs  recently  found by Oppenheimer  et  al.  (2001)  is  still   controversial.  Whereas  these authors  claim that they have  finally  found the elusive  ancient  halo white dwarf population that contributes significantly to the mass budget of the galactic halo, there have been several other  contributions  that argue that these white dwarfs are not genuine halo members but, instead, thick disk stars.  We show here that the  interpretation  of this sample is  based  on  the  adopted   distances,   which  are  obtained  from  a color--magnitude  calibration, and we demonstrate  that when the correct distances  are used a sizeable  fraction  of these  putative  halo white dwarfs belong indeed to the disk  population.  We also perform a maximum likelihood  analysis of the  remaining  set of white  dwarfs and we find that they most  likely  belong to the thick  disk  population.  However, another  possible  explanation  is that this sample of white  dwarfs has been  drawn  from a 1:1  mixture  of  the  halo  and  disk  white  dwarf populations. ",
        "introduction": "White  dwarfs are the most  common  end--points  of  stellar  evolution. Since they are long-lived and well understood  objects, they  constitute an invaluable tool to study the evolution and structure of our Galaxy in general and of the Galactic  halo in  particular  (Isern et al.  1998a). Moreover,  the  discovery of  microlenses  towards the Large  Magellanic Cloud (Alcock et al.  2000; Lasserre et al.  2001) has generated a large controversy about the possibility that white dwarfs could be responsible for these  microlensing  events and, thus, could  provide a  significant contribution  to the mass budget of our Galactic  halo.  However,  white dwarfs as viable dark matter  candidates are not free of problems, since an excess of them would  imply as well an  overproduction  of red dwarfs and Type II  supernovae.  In order to  overcome  this  problem  Adams \\& Laughlin (1996) proposed a non--standard  initial mass function in which the  formation  of both low and high mass stars was  suppresed.  Besides the lack of evidence for such biased initial mass  functions,  they also present additional problems.  The formation of an average mass ($\\sim \\, 0.6\\,  M_\\odot$)  white dwarf is accompanied  by the injection  into the interstellar medium of a sizeable amount of mass (typically $\\sim 1.5 \\, M_\\odot$) per white white dwarf.  Since in turn Type II  supernovae  are suppressed in biased initial mass functions,  there is not enough energy to eject this matter  into the  intergalactic  medium and a mass that is roughly  three  times the mass of the  resulting  white  dwarf has to be accomodated into the Galaxy (Isern et al.  1998).  Furthermore, the mass ejected in the process of  formation  of a white dwarf is  significantly enriched in metals (Abia et al.  2001; Gibson \\& Mould  1997).  Finally, an excess of white  dwarfs  may  translate  into an excess  of  binaries containing  such stars.  If there are many white dwarfs in binaries then the secondary  cannot be a red dwarf because  these would these would be easily  detected.  Therefore,  we are then  forced to assume  that these binaries are double degenerates, which are one of the currently proposed scenarios  for Type Ia  supernovae.  Hence  we are  forced  to face  the subsequent  increase of Type Ia  supernova  rates  which,  consequently, results in an increase in the  abundances  of the  elements  of the iron peak   (Canal,   Isern   \\&   Ruiz--Lapuente   1997).   However,   other explanations,  such as self--lensing  in the LMC (Wu 1994; Salati et al. 1999), or background  objects (Green \\& Jedamzik  2002) are possible and have not been yet totally ruled out.  The debate of whether or not white dwarfs  contribute  significantly  to the  Galactic  halo  dark  matter  has   motivated  a  large  number  of observational  searches  (Knox,  Hawkins  \\& Hambly  1999;  Ibata et al. 1999;  Oppenheimer et al.  2001;  Majewski \\& Siegel 2002, Nelson et al. 2002) and theoretical works (Reyl\\'e, Robin \\& Crez\\'e 2001; Koopmans \\& Blandford  2002; Flynn,  Holopainen \\& Holmberg 2002) and is still open. Among the  observational  surveys perhaps the most extensive one is that of Oppenheimer et al.  (2001) who  discovered 38 faint white dwarfs with large  proper  motions  in  digitized   photographic   plates  from  the SuperCOSMOS  Sky Survey.  Oppenheimer et al.  (2001)  claimed that these white  dwarfs are indeed  halo white  dwarfs  since they have very large tangential  velocities  (in excess of $\\sim 100$ km s$^{-1}$).  Based on this  assumption,  they  derived a space  density of 2\\% of the Galactic dark halo density, which is smaller than previous  claims (Alcock et al. 1997)  for halo  dark  matter  in the form of  $\\approx  0.5\\,  M_\\odot$ objects, but still  significant.  However,  Reid, Sahu \\& Hawley  (2001) challenged  this  claim by noting  that the  kinematics  of these  white dwarfs is  consistent  with the  high--velocity  tail of the thick disk. Hansen (2001)  provided  evidence that this sample  presents a spread in age that makes it more  likely to belong to the thick  disk  population. Reyl\\'e  et al.  (2001)  and  Flynn  et al.  (2002)  also  support  this interpretation.  Koopmans \\& Blanford (2002) find that the  contribution of these white dwarfs to the local halo dark matter  density is smaller, of the order of 0.8\\%, which is in good agreement  with the  theoretical results of Isern et al.  (1998b) and the  observational  findings of the EROS team (Goldman et al.  2002).  In this paper we reexamine this issue by making  use of a Monte  Carlo  simulator  (Garc\\'\\i  a--Berro  et al. 1999;  Torres et al.  1998).  The  paper is  organized  as  follows.  In section  \\S 2  we  present  the  main  properties  of  our  Monte  Carlo simulator.  In  \\S 3 we  discuss  the  effect  of  the  color--magnitude calibration  on the  distances  of the  white  dwarfs  in the  sample of Oppenheimer  et al.  (2001)  whereas  in \\S 4 we  analyze  which  is the probability of this sample to be drawn from a halo  population.  Finally in \\S 5 our conclusions are summarized. ",
        "conclusions": "We have presented evidence that the distances of the white dwarfs in the sample of Oppenheimer et al.  (2001) have not been correctly determined. The ultimate reason of this is that the authors used a calibration which is not  appropriate  for the  halo  white  dwarf  population.  Once  the correct  calibration  is adopted it turns out that the  distances to the most  luminous  white  dwarfs in the  sample  have been  underestimated, whereas the distances to the white dwarfs with small  luminosities  have been  overestimated.  We have also found that some  white  dwarfs in the sample cannot have hydrogen dominated  atmospheres, since their position in  the   color--magnitude   diagram  is  beyond  the   turn-off.  As  a consequence, once the corrected distances are taken into account, a good fraction of these putative halo white dwarfs have significantly  smaller tangential  velocities  and can be  safely  discarded  as  genuine  halo members.  The remaining  fraction of the sample of Oppenheimer  et al.  (2001) has been analyzed using our Monte Carlo  simulator.  We have computed  Monte Carlo models for the disk and the halo populations.  The disk simulation naturally  recovers both the thin and the thick disk  populations.  Then we  have  computed  the  probability  of the  stars  of  the  sample  of Oppenheimer  et al.  (2001) to belong to a randomly  selected  sample of both halo or disk white dwarfs.  Our results  indicate  that this subset of the sample of Oppenheimer et al.  (2001) does not belong  exclusively to either the halo or the disk population at the 95\\% confidence  level. Regarding  the disk  population  {\\sl as a whole} our  results  were not conclusive  because  of the  small  fraction  of thick  disk  stars in a typical Monte Carlo simulation.  However once the stars with small birth times ($\\lapprox$ 2 Gyr),  corresponding to the thick disk, are selected we find that the number of stars in the  sample  nicely  reproduces  the values  found by  Oppenheimer  et al.  (2001),  in  agreement  with  the results of Flynn et al.  (2002)  and  Reyl\\'e  et al.  (2001).  There is yet another possibility which has not been previously  explored.  Namely that the sample of Oppenheimer et al.  (2001) is drawn from a mixture of both the halo and the (thick)  disk  populations.  We have found that in this case the  probability  is  maximum  for a 1:1 to  ratio.  Hence, we conclude that the claim by Oppenheimer et al.  (2001) that, finally, the elusive halo white dwarf  population has been found should be taken with caution and more  observational  searches and theoretical work are still needed.  Finally we have  re-derived,  using the  distances  obtained in this work, the number  density of halo  white  dwarfs  predicted  by the sample of  Oppenheimer  et al.  (2001).  We have found that a safe upper limit to this density is $n=6.2 \\cdot 10^{-5}$ pc$^{-3}$,  assuming that {\\sl all} the white dwarfs found by Oppenheimer  et al.  (2001) are true halo white dwarfs.  If, as suggested by our  simulations, we assume that only half of these  stars  are  genuine  halo  members  we find a number density of $3.1 \\cdot  10^{-5}$  pc$^{-3}$,  which is in good  agreement with previous independent determinations.  \\vspace{1cm}  \\noindent {\\sl  Acknowledgements.}  This work has been  supported by the DGES grant  PB98--1183--C03--02, by the MCYT grant AYA2000--1785, by the MCYT/DAAD grant  HA2000--0038 and by the CIRIT grants  1995SGR-0602  and 2000ACES-00017.  We would like to acknowledge the advise of Nigel Hambly in  transforming  our cooling  sequences to the  appropiate  photometric passbands.  We also would like to acknowledge the very valuable comments of our referee, Chris Flynn, which greatly improved the original version of the manuscript."
    },
    "0207/astro-ph0207439_arXiv.txt": {
        "abstract": "{The absorption feature detected in the prompt X-ray emission of GRB990705 has important consequences for its circum-burst environment and therefore on its afterglow. Here we investigate whether the circum-burst environment constrained by the absorption feature could be consistent with the observed $H$-band afterglow, which exhibits an earlier power law decay ($F\\propto t^{-1.68}$) but a much faster decay ($\\alpha>2.6$; $F\\propto t^{-\\alpha}$) about one day after the burst. Two possible geometries of the afterglow-emitting regions are suggested: 1) afterglow emission  produced by the impact of the fireball on the surrounding torus,  which serves as the absorbing material of the X-ray feature, as would be expected in the models involving that a supernova explosion precedes the gamma-ray burst by some time; 2)afterglow {emission} produced in the dense circum-burst medium inside the torus. In case 1), the faster decay at the later time is attributed to the disappearance of the shock due to the counter-pressure in the hot torus illuminated by the burst and afterglow photons. For case 2), the circum-burst medium density is found to be very high ( $n\\ga 10^4-10^5~ {\\rm cm^{-3}}$ ) if the emitting plasma is a jet or even higher if it is spherical. Future better observations of afterglows of GRBs that have absorption features might make it possible to make a more definite choice between these two scenarios. \\keywords gamma rays: bursts---line: formation---radiation mechanism: nonthermal }  \\titlerunning{Afterglows of GRBs with X-ray features} ",
        "introduction": "There is increasing observational evidence favoring the existence of absorption and emission lines in the X-ray spectra of gamma-ray bursts (GRBs) and their afterglows. Emission or absorption features can provide a fundamental tool for studying  the close environment of GRBs (e.g. M\\'{e}sz\\'{a}ros \\& Rees 1998; Lazzati et al. 1999, 2002; B\\\"{o}ttcher \\& Fryer 2001). To date, five bursts have shown evidence for  iron or lighter element emission lines during the X-ray afterglow (GRB970508, Piro et al. 1998; GRB970828,  Yoshida et al. 1999; GRB991216, Piro et al. 2000; GRB000214,  Antonelli et al. 2000; GRB011211, Reeves et al. 2002) and one (GRB990705; Amati et al. 2000; hereafter A2000) displays a transient absorption feature at 3.8 KeV during the burst itself.  A few models for emission lines in the X-ray afterglows have been suggested (see Piro 2002 for a review ), including ``distant reprocessor scenario\" and ``nearby reprocessor scenario\". In the former, the line-emitting gas is located at $R\\ga 10^{15}~{\\rm cm}$ with the line variability time corresponding to the light travel time between GRB and the reprocessor (Lazzati et al. 1999; Piro 2000; Weth et al. 2000). This scenario needs the presence of an iron-rich dense medium with iron mass $M_{\\rm Fe}\\ga 0.01 M_\\odot$. The most straightforward picture is the one in which  a SN-like explosion occurs some time before the formation of the GRB. The GRB may be produced by the collapsing of the rotationally-supported newborn massive neutron star to a black hole (Vietri \\& Stella 1998), or the phase transition to a strange star (Wang et al. 2000a) . In the latter scenario, the line emission is attributed to the interaction of a long-lasting relativistic outflow from the central engine with the massive star progenitor stellar envelope at distances $R\\la10^{13}~{\\rm cm}$ (M\\'{e}sz\\'{a}ros \\& Rees 2000; Rees \\& M\\'{e}sz\\'{a}ros 2000).  While different scenarios have been suggested to explain the emission line, the properties of the {transient} absorption feature, as in GRB990705, strongly point to a unique {circum-burst environment (Lazzati et al. 2001; B\\\"{o}ttcher et al. 2002), i.e.}, 1) iron-rich absorbing matter of a few solar masses (such as the young supernova {remnant} shell ) lies between $10^{16}$ and $10^{18}~$cm from the burst site; 2) the absorbing matter is located in the line of sight between the observer and the burster.  GRB990705 has a duration of $\\sim42 ~{\\rm s }$ in the Gamma-Ray Burst Monitor (GRBM) and fluence $(9.3\\pm 0.2)\\times10^{-5}~{\\rm erg~ cm^{-2}}$ in the $2-700 {\\rm keV}$ band (A2000). During the prompt phase, it shows an absorption feature at 3.8 keV and an equivalent hydrogen column density, which disappears 13 s after the burst onset (A2000). This absorption feature was explained by A2000 as being due to an edge produced by neutral iron redshifted to $3.8\\pm0.3~{\\rm keV}$; the corresponding redshift is $0.86\\pm0.17$. Optical spectroscopy of the host galaxy gives a redshift $z=0.8435$ (Andersen et al. 2002), consistent with the inferred value from the X-ray feature. This straightforward interpretation was, however, questioned by Lazzati et al. (2001) as it requires a vast amount of iron {\\footnote { The required total mass of iron is $35 f M_\\odot$ ( see Eq.(5) in Lazzati et al. 2001), where $f$ is the covering factor of the absorbing material surrounding the burst.}}  in the close vicinity of the burster. Lazzati et al. (2001) further suggested an alternative scenario in which the feature is produced by resonant scattering from hydrogen-like iron broadened by a range of outflow velocities. In this scenario, the radius of the SN shell is fixed by the requirement that the heating timescale of the electrons in the absorbing matter is $\\sim 10~{\\rm s}$, i.e. $R_s\\sim (2-3)\\times10^{16}~{\\rm cm}$. {Our following work is based on this scenario.}    A fading X-ray afterglow of GRB990705 was detected by the Narrow Field Instruments of {\\it BeppoSAX} 11 hours after the trigger, but the statistics are not sufficient to draw a detailed conclusion on the decaying law (A2000). Masetti et al. (2000) report having detected  the counterpart of this burst twice in the near-infrared $H$ band and only once in the optical $V$ band, from a few hours to $\\sim1$ day after the GRB trigger. The first two $H$-band measurements define a power-law decay with index $\\alpha=1.68\\pm0.10$ ($F\\propto{t^{-\\alpha}})$, but a third attempt to detect the source gave an upper limit, implying a much faster decay. No radio afterglow was detected (Subrahmanyan et al. 1999; Hurley et al. 1999).  For the afterglows with X-ray {\\em emission }lines, the line-emitting gas  could lie outside of the line of sight of the burst and therefore has no direct relation with the afterglow radiation. However,  for afterglow with X-ray {\\em absorption } features, the absorbing matter (SN shell) should have a direct consequence on the afterglow radiation, because it  must lie in the line of sight of the burst. So, an examination of the self-consistency between the power-law afterglow and the X-ray absorption feature is quite necessary. ",
        "conclusions": "Emission or absorption features in the X-ray spectrum of GRBs and their afterglows provide a useful tool  for studying the close environment of GRBs and thus their possible progenitors. The absorption feature in the prompt X-ray emission of GRB990705 was originally interpreted by Amati et al. (2000) to be a photoionization  K edge of neutral iron. However, this straightforward explanation is shown by Lazzati et al. (2001) to require an improbably large  amount of iron in the close environment of the burster. Instead, Lazzati et al. (2001) interpret this as a resonant absorption line broadened by a large spread of velocities. In this scenario, the disappearance of the feature 13 s after the burst results from electron heating due to the illuminating photons and it severely constrains the radius of the absorbing materials ($R\\sim2-3\\times10^{16}~{\\rm cm}$, see Eq. (13) of Lazzati et al. 2001). A reasonable scenario for this requirement is the supranova-like scenarios  ( Vietri \\& Stella 1998; Wang et al. 2000a), in which a young supernova remanent  is located at the close vicinity of the burster. Based on these studies, in this paper we investigated whether the circum-burst environment constrained by the absorption feature could be consistent with the observed afterglows of GRB990705.  We discussed  two  possible locations of the afterglow-emitting region: one is in the torus where the afterglows are produced by the impact of the fireball jet on this torus and the other is in the dense circum-burst medium inside the torus. In the former scenario, the impact of the fireball on the torus will generate a forward shock propagating into the torus. This forward shock will be decelerated by the dense matter in the torus into a sub-relativistic phase in quite a short time and to a lower and lower velocity as  time elapses. The heating/cooling processes of the torus by the burst and afterglow photons may bring its temperature to $T_s\\sim10^7~{\\rm K}$. Once  the ram pressure ($\\sim \\rho_b v^2$ ) of the fireball  falls low enough to be equal to the thermal counter-pressure ($n_skT_s$) of the hot torus,  the forward shock is damped down very rapidly (Vietri et al. 1999) and the afterglow emission will cut off accordingly. We found that the $H$-band afterglow of GRB990705 can be fitted in terms of this model.  In the latter scenario, as in many other afterglows, the steeping of light curve decay of GRB990705 one day after the burst is attributed to the jet evolution in a uniform density medium or a spherical fireball undergoing a transition to non-relativistic expansion. The broken power-law decay behavior of the $H$-band afterglow requires the shock radius at the light curve break time or at the Sedov phase, respectively, to be smaller than the torus location. This in turn requires that the circum-burst medium density must be $n\\ga10^4-10^5~{\\rm cm^{-3}}$ or $n\\ga10^6~{\\rm cm^{-3}}$, respectively. In this scenario, the fireball will also hit the surrounding torus finally. The abrupt density jump  might cause a rise and a successive decline in the afterglows (see Dai \\& Lu 2002 for a relativistic case).  A noticeable point relevant to the high density circum-burst medium is that the true energy reservoir  of GRB990705 may be much greater than what was estimated by Frail et al. (2001), $E_\\gamma=3.9\\times10^{50}~{\\rm erg}$, derived from the jet model by assuming an interstellar medium of density $n=0.1~{\\rm cm^{-3}}$, since the calculated fireball true energy depends on $\\theta_j^2$ which in turn depends on $n^{1/4}$.  In summary, the geometry requirement of  the X-ray absorption feature of GRB990705 is shown to be also consistent with its afterglows, although the sparse data of the afterglow makes it impossible to reach a definite conclusion on the two scenarios. Future better broad-band observations of the afterglow spectra and light curves for GRBs that have absorption features could tell which one is true and thereby provides a more valuable insight into the environment and the central engine."
    },
    "0207/astro-ph0207263_arXiv.txt": {
        "abstract": "During an \\xmm\\ observation, the eclipsing polar UZ For was found in a peculiar state with an extremely low X-ray luminosity and occasional X-ray and UV flaring. For most of the observation, UZ For was only barely detected in X-rays and $\\sim\\!800$ times fainter than during a high state previously observed with \\rosat. A transient event, which lasted $\\sim\\!900\\:$s, was detected simultaneously by the X-ray instruments and, in the UV, by the Optical Monitor. The transient was likely caused by the impact of $10^{17}$--$10^{18}\\:$g of gas on the main accretion region of the white dwarf. The X-ray spectrum of the transient is consistent with $\\sim\\!7\\:$keV thermal bremsstrahlung from the shock-heated gas in the accretion column. A soft blackbody component due to reprocessing of X-rays in the white dwarf atmosphere is not seen. The increase in the UV flux during the transient was likely caused by cyclotron radiation from the shock-heated gas. Two more flaring events were detected by the Optical Monitor while the X-ray instruments were not operating. We conclude from our analysis that the unusual flaring behavior during the low state of UZ For was caused by intermittent increases of the mass transfer rate due to stellar activity on the secondary. In addition to the transient events, the Optical Monitor detected a roughly constant UV flux consistent with 11000--K blackbody radiation from the photosphere of the white dwarf. We find a small orbital modulation of the UV flux caused by a large, heated pole cap around the main accretion region. ",
        "introduction": "UZ For is an eclipsing member of the subclass of cataclysmic variables called AM Her binaries or polars. In these binaries, the strong magnetic field of the white dwarf primary causes it to rotate synchronously with the orbital motion. The magnetic field also prevents the formation of an accretion disk around the white dwarf. The accretion stream from the Roche-lobe filling secondary is funneled along the magnetic field lines and impacts the white dwarf near a magnetic pole. Slightly above the surface, the accretion stream forms a stand-off shock that heats the gas to temperatures in excess of $10^8\\:$K. The shock-heated plasma then cools and settles on to the white dwarf while strongly emitting cyclotron radiation (IR to UV) and thermal bremsstrahlung (mostly X-rays). The photosphere below the shock is heated by reprocessing of X-rays and emits blackbody radiation visible at soft X-ray and UV energies. Many polars have been observed to decline in brightness by several magnitudes and remain in a faint state for days to years. The causes of these low states are not known, but, in the absence of an accretion disk, the large brightness variations must be due to changes in the mass transfer rate from the companion star. A comprehensive review of polars is given in \\citet{1995cvs..book.....W}.   UZ For (EXO 033319-2554.2) was first detected as a serendipitous X-ray source with \\exosat\\ \\citep{1987IAUC.4486....1G,1988ApJ...328L..45O}. Subsequent optical spectroscopy and polarimetry established UZ For as an eclipsing polar \\citep{1988A&A...195L..15B,1988ApJ...329L..97B}. The 126.5--min orbital period is close to the lower edge of the 2--3 hr \"period gap\", a sparsely populated region in the orbital period distribution of cataclysmic variables. UZ For is a high inclination system ($i\\approx80^\\circ$), and both accretion regions are eclipsed by the white dwarf for at least half of an orbital cycle. The optical spectrum of UZ For shows strong cyclotron emission lines which indicate magnetic fields of $\\sim\\!53\\:$MG and $\\sim\\!48\\:$MG for the two accretion poles \\citep{1996A&A...310..526R}. The best available estimates for the distance and the white dwarf mass are $d=208\\pm40\\:$pc and $M_{WD}\\approx0.6$--$0.8\\:M_\\odot$ \\citep{1989ApJ...337..832F,1991MNRAS.253...27B}.   In this paper we present X-ray and UV data obtained with \\xmm\\ while UZ For was in a state of unusually low and irregular accretion. We study the properties of an X-ray/UV flare and show that it was caused by accretion on to the white dwarf. We argue that the intermittent accretion rate increase was due to stellar activity on the companion star. Part of the \\xmm\\ data has recently been published in \\citet{2001ApJ...562L..71S}. The goal of our paper is to utilize all available \\xmm\\ data and present additional results not previously published. In particular, we find two more UV flares and clearly identify eclipses in the UV light curves. We show that the beginning of the X-ray/UV flare coincides with the eclipse egress of the main accretion region. This provides strong evidence that the flare was an accretion event on the white dwarf. We estimate the total accreted mass and show that it is consistent with the mass ejected by stellar flares on M dwarfs. We also detect a weak orbital modulation of the X-ray and UV fluxes and demonstrate that it is most likely due to emission from the main accretion region. ",
        "conclusions": "During the \\xmm\\ observation, UZ For was found in an extremely low accretion state with an X-ray luminosity $\\sim\\!800$ times fainter than during a high state previously observed with \\rosat. Occasional X-ray and UV flaring was detected by the X-ray instruments and the Optical Monitor. The largest flare lasted $\\sim\\!900\\:$s and increased the X-ray flux by a factor of $\\sim\\!30$. We found that the beginning of this flare coincided to within a few seconds with the eclipse egress of the main accretion region. This provides strong evidence that the flaring was caused by accretion on to the white dwarf. The X-ray spectrum of the flare is consistent with $\\sim\\!7\\:$keV thermal bremsstrahlung from the accretion column. A blackbody component, as seen with \\rosat\\ during the high state, was not found. It is plausible that, because of a larger accretion region or the absence of blobs in the accretion stream, the blackbody temperature was too low for a detection by the X-ray instruments. The increase in the UV flux seen during the flare was probably caused by cyclotron radiation from the accretion column. A significant contribution of blackbody radiation to the UV flare emission is unlikely as this would require a very large soft excess. Under the assumption that all accretion energy is emitted as bremsstrahlung, we estimate an accretion rate of $2\\times10^{13}\\:\\mathrm{g\\:s^{-1}}$ during the flare. However, since at this low rate most of the energy is emitted as cyclotron radiation, the actual accretion rate was probably $10^1$--$10^2$ times higher. We therefore estimate that during the flare a total of $10^{17}$--$10^{18}\\:$g of gas was accreted on to the white dwarf. The likely cause of the flaring observed in UZ For is stellar activity on the companion star that intermittently increased the mass transfer rate near the L1-point. The mass that was accreted on to the white dwarf during the large transient is consistent with the mass ejected by a stellar flare.    Before and after the X-ray transient, extremely weak X-ray emission, possibly due to the regular low-state accretion on to the main region, was detected. The observed X-ray luminosity corresponds to an accretion rate of $6\\times10^{11}\\:\\mathrm{g\\:s^{-1}}$. Since cyclotron radiation dominated the energy output, the actual accretion rate was probably $10^{13}$--$10^{14}\\:\\mathrm{g\\:s^{-1}}$, which is similar to the rates estimated for previous low states. In addition to the flare emission, we detect a roughly constant UV flux consistent with blackbody radiation from a 11000--K white dwarf. A small orbital modulation of the UV flux indicates the presence of a large, heated pole cap around the main accretion region.   Flaring during a low state has only been observed for a few polars. This may be due to insufficient monitoring of polars in low states. Yet low-state observations are essential since, during high and intermediate states, flaring caused by stellar activity on the companion star is likely overlooked and mistakenly attributed to accretion stream instabilities. Future low-state observations will reveal how common flaring due to stellar activity is among polars. An interesting question that might also be answered is, whether flaring occurs preferentially at the beginning or end of low states. Monitoring of irregular accretion in low-state polars may provide a new way to study stellar flares or other types of mass ejections that are too faint to be observed directly."
    },
    "0207/gr-qc0207012_arXiv.txt": {
        "abstract": "Assuming a Friedmann universe which evolves with a power-law scale factor, $a=t^{n}$, we analyse the phase space of the system of equations that describes a time-varying fine structure 'constant', $\\alpha$, in the Bekenstein-Sandvik-Barrow-Magueijo generalisation of general relativity. We have classified all the possible behaviours of $\\alpha (t)$ in ever-expanding universes with different $n$ and find new exact solutions for $\\alpha (t)$. We find the attractors points in the phase space for all $n$. In general, $\\alpha $ will be a non-decreasing function of time that increases logarithmically in time during a period when the expansion is dust dominated ($n=2/3$), but becomes constant when $n>2/3$. This includes the case of negative-curvature domination ($n=1$). $\\alpha $ also tends rapidly to a constant when the expansion scale factor increases exponentially. A general set of conditions is established for $\\alpha $ to become asymptotically constant at late times in an expanding universe. ",
        "introduction": "Stimulated by the observations for small variations in atomic structure controlled by the fine structure constant in quasar absorption lines at redshifts $z=1-3$, \\cite{murphy,webb2,webb}, there has been much recent interest in the theoretical predictions of gravity theories which extend general relativity to incorporate space-time variations of the fine structure 'constant'. These have been primary formulated as Lagrangian theories with explicit variation of the velocity of light, $c$, \\cite {moffatal,am,ba}, or of the charge on the electron, $e,$ \\cite {bsbm,bsm1,bsm2,bsm3}. Theories of the latter sort offer the possibility of matching the magnitude and trend of the quasar observations and have been studied numerically and by means of matched asymptotic approximations. They are also of particular interest because they predict that violations of the weak equivalence principle should be observed at a level that is within about an order of magnitude of existing experimental bounds \\cite {bsm4,zal,mof}. They are consistent with all other astrophysical and experimental limits of time variation of the fine structure constant and predict effects of the microwave background radiation, primordial nucleosynthesis, and the Oklo natural reactor that are too small to conflict with current observational bounds. A range of variant theories have been investigated with attention to the possible particle physics motivations and consequences for systems of grand and partial unification in references \\cite {banks,guts,olive}. In this paper we will give a full qualitative analysis of the properties of Friedmann cosmological models in a sub-class of these theories developed initially by Bekenstein \\cite{bek2} to generalise Maxwell's equations to include varying $e$ and then generalised by Sandvik, Barrow and Magueijo \\cite{bsbm} to include gravitation. We refer to these as BSBM theories. We provide a phase-space analysis of the non-linear propagation equation for the scalar field which carries the variations of the fine structure constant. Some new exact solutions are also given and all the asymptotic behaviours classified. ",
        "conclusions": "Using a phase plane analysis we have studied the cosmological evolution of a time-varying fine-structure 'constant' $\\alpha (t)=\\exp [2\\psi ]$, in the BSBM theory. We have considered the cases created by power-law evolution of the expansion scale factor of the universe. We have shown that in general $% \\alpha $ increases with time or asymptotes to a constant value at late times. We have found a new exact solution for the case of a universe dominated by a stiff fluid or massless scalar field. We have also found general asymptotic solutions for all the different and possible behaviours via the analysis of the critical points of the system that determines the evolution of $\\alpha $. In particular, we have found asymptotic solutions for the dust, radiation and curvature-dominated FRW universe which also generalise the asymptotes found in \\cite{bsm1}.These solutions correspond to late-time attractors that describes the $\\psi $ and $\\alpha $ evolution in time and will enable a more detailed analysis to be made of the fit between theoretical expectations of varying $\\alpha $ theories and observations of relativistic fine structure in atoms at high redshift.   \\vspace{2cm} \\noindent \\textbf{Acknowledgements }We would like to thank% \\textbf{\\ }H\\aa vard Sandvik and Jo\\~{a}o Magueijo for discussions. DFM is supported by Funda\\c{c}\\~ao para a Ci\\^encia e a Tecnologia, Portugal, through the research grant BD/15981/98.  \\vspace{2cm} \\noindent \\appendix\\textbf{\\appendixname :} \\vspace{.5cm} \\noindent \\textbf{\\ General analysis of the phase plane bifurcations}  In previous sections we have analysed the $\\psi $ evolution equation (\\ref {psi}) for a range of variables which are physically realistic and correspond to expanding universes. We will now analyse the whole range for variables of the system (\\ref{llog}). As before we see there are two critical points in the $(w,v)$ plane, at $(0,-1+{\\frac{3n}{2}})$ and $% ((1-3n)(\\frac{3n}{2}-1),0)$. Linearising (\\ref{llog}) about $% (w_{c_{1}},v_{c_{1}})=(0,-1+\\frac{3n}{2})$ and $% (w_{c_{2}},v_{c_{2}})=((1-3n)(\\frac{3n}{2}-1),0)$ we obtain the following characteristics matrices:  $M_1=% \\bordermatrix{ & &\\cr & 2-3n  & 0   \\cr & 1 & 1-3n \\cr} $ \\qquad $M_2=% \\bordermatrix{& & \\cr & 0 & (3n-1)(2-3n)    \\cr & 1 & 1-3n    \\cr} $  The characteristic matrices are non singular except when $n=\\frac{1}{3}$ or $% n=\\frac{2}{3}$. In the non-singular cases the critical points will be simple and the system defined by these differential equations is structurally stable \\cite{andronov}, and there will be no 'strange' chaotic behaviour outside the neighbourhood of the critical points. Hence, the linearised system will have the same phase portrait as non-linearised one in the neighbourhood of the critical points.  The evolution, with respect to changing $n$, of the signs of the determinant and the trace of these two matrices is given in the table. This show us that there are always two critical points in our system, an unstable saddle and an attractor (which changes from a spiral to a node).  \\vspace{.5cm}  \\begin{tabular}[hb]{|c|c|c|} \\hline \\textbf{$n$} & \\multicolumn{2}{|c|}{\\textbf{Critical Points $(w_{c},v_{c})$}} \\\\ \\cline{2-3} & $\\left( \\left( 3n-1\\right) \\left( 1-\\frac{3n}{2}\\right) ,0\\right) $ & $% \\left( 0,1-\\frac{3n}{2}\\right) $ \\\\ \\hline $(-\\infty ;\\frac{1}{3})$ & Saddle Point (non-physical) & Unstable Node \\\\ & det $M_{1}<0$, Tr $M_{1}>0$ & det $M_{2}>0$, Tr $M_{2}>0$ \\\\ \\hline\\hline $\\frac{1}{3}$ & Origin & Axis \\\\ & det $M_{1}=0$, Tr $M_{1}=0$ & det $M_{2}=0$, Tr $M_{2}>0$ \\\\ \\hline\\hline $(\\frac{1}{3};\\frac{1}{2})$ & Stable Spiral & Saddle Point \\\\ & det $M_{1}>0$, Tr $M_{1}<0$ & det $M_{2}<0$, Tr $M_{2}>0$ \\\\ \\hline\\hline $\\frac{1}{2}$ & Stable Spiral & Saddle Point \\\\ & det $M_{1}>0$, Tr $M_{1}<0$ & det $M_{2}<0$, Tr $M_{2}=0$ \\\\ \\hline\\hline $(\\frac{1}{2};\\frac{3}{5})$ & Stable Spiral & Saddle Point \\\\ & det $M_{1}>0$, Tr $M_{1}<0$ & det $M_{2}<0$, Tr $M_{2}<0$ \\\\ \\hline\\hline $\\frac{3}{5}$ & Stable Spiral (node) & Saddle Point \\\\ & det $M_{1}>0$, Tr $M_{1}<0$ & det $M_{2}<0$, Tr $M_{2}<0$ \\\\ \\hline\\hline $(\\frac{3}{5};\\frac{2}{3})$ & Stable Node & Saddle Point \\\\ & det $M_{1}>0$, Tr $M_{1}<0$ & det $M_{2}<0$, Tr $M_{2}<0$ \\\\ \\hline\\hline $\\frac{2}{3}$ & Stable Axis & Stable Axis \\\\ & det $M_{1}=0$, Tr $M_{1}<0$ & det $M_{2}=0$, Tr $M_{2}<0$ \\\\ \\hline\\hline $(\\frac{2}{3};\\infty )$ & Saddle Point (non-physical) & Stable Node \\\\ & det $M_{1}<0$, Tr $M_{1}<0$ & det $M_{2}>0$, Tr $M_{2}<0$ \\\\ \\hline \\end{tabular}  \\vspace{.5cm}  The cases $n=\\frac{1}{3}$ or $n=\\frac{2}{3}$, where the determinant of the characteristic matrixes vanishes, lead to a bifurcation of codimension $1$, in particular, of Saddle-Node type \\cite{wiggins}, since they correspond to points where the determinants of the characteristic matrices change sign, det% $M_{1}$ $=$ det$M_{2}=0$. At these values of $n$ the nature of the system will change. Cosmologically, these points represent a change in the behaviour of the time evolution of the fine structure 'constant' as can be seen from the figures: \\ref{n025}, \\ref{n13}, \\ref{n12}, \\ref{n35}, \\ref{n23}% , \\ref{n1}, which display the time evolution of $\\psi$. When $n$\\ starts to grows from $-\\infty $\\ to $\\frac{1}{3}$\\ the two critical points slowly\\emph{% \\ }converge at $n=\\frac{1}{3}$\\emph{. }For example, in the $n=0$ case where we may without loss of generality set $a=1$, equation (\\ref{psi}) has the exact solution  \\[ \\alpha =\\exp [2\\psi ]=A^{-2}\\cosh ^{2}[AN^{1/2}(t+t_{0})] \\] where $A,t_{0}$ are constants. This is an unrealistically rapid growth asymptotically , $\\psi \\varpropto t,$caused by the absence of the inhibiting effect of the cosmological expansion. The case of $n=1/4$ is shown in Figure 6, which shows the phase space trajectories and the evolution of $\\psi $ vs. $\\ln t.$  \\begin{figure}[hb] \\centering \\epsfig{file=n025.ps,height=6cm} % \\epsfig{file=n025psi.ps,height=6cm} \\caption{{\\protect\\small \\textit{Numerical plots of the phase space in the $% (w,v)$ coordinates, and the $\\protect\\psi$ evolution with $x=log(t)$ for $n=\\frac{1}{4}$. The '$+$' sign is a saddle point and the square is an unstable source. }}} \\label{n025} \\end{figure}  At $n=\\frac{1}{3}$ we are in the situation where the two critical points collapse into a unique one at the origin, creating a saddle-node bifurcation and a concomitant change in the behaviour and evolution of $\\psi $. As $n$ keeps growing the single critical point splits into two critical points again. They move apart until the radiation value is reached, $n=\\frac{1}{2}$ (Tr $M_{2}=0$). In this case, $\\psi $ is a asymptotically monotonic growing function of time, with some small oscillations near the Planck epoch. However, note that in our universe the asymptote giving an increase of $\\psi $ behaviour with time is never reached before the dust-dominated evolution takes over \\cite{bsbm}, \\cite{bsm1}, \\cite{bsm2}.  As the universe evolves to the dust-dominated epoch, and $n$ approaches the intermediate behaviour $n=\\frac{3}{5}$, the two critical points start to coalesce again into a single point. When $n=\\frac{3}{5}$ is reached, $\\psi $ becomes a strictly monotonically growing function of time. When $n$ reaches the value corresponding to a dust-dominated universe, $n=\\frac{2}{3}$, another saddle-node bifurcation occurs. The two critical points collapse into a single one. Again there will be change in the behaviour of $\\psi $ for larger values of $n$. Accordingly, when $n>2/3,$ the two critical points reappear once again and $\\psi $ becomes asymptotically constant in value.  Notice, that although a bifurcation is something that 'spoils' the smooth behaviour of a system, in our case, that won't happen, due to the physical constraints of our variables. In reality due to those constraints, the physical system will never 'feel' the abrupt change at $n=\\frac{1}{3}$ and $% n=\\frac{2}{3}$. This is also due to the fact that the attracting critical point always lies in the physical range of the variables, while the unstable one disappears form the physical system when the bifurcations occur, as can be seen from the phase plane plots."
    },
    "0207/astro-ph0207524_arXiv.txt": {
        "abstract": "The acoustic spectra in sunspots are known to be richer in higher frequency power. We have attempted a generalized study of the effect of magnetic fields on the shape of the acoustic spectrum using GONG+ bread-board data (spatial scale of $\\sim$ 2 arc-sec per pixel) of 11 May, 2000 and 12 June, 2000. The mean power spectra of the velocity oscillations were obtained by averaging over several spectra for different values of the magnetic field. With increasing magnetic field, the acoustic power increases at higher frequencies and decreases at lower frequencies with a transition at $\\simeq$ 5 mHz. This behaviour is slightly different from earlier results obtained from SOHO/MDI data. ",
        "introduction": "The typical distribution of acoustic power of photospheric oscillations shows a maximum at around 5 min (Leighton, Noyes, and Simon, 1962) with a decrease to negligible power at higher frequencies. This behaviour of the acoustic spectrum has been understood in terms of trapped oscillations in a cavity. The eigen-functions of the natural modes of oscillations peak at different depths in the solar interior. The band of oscillations in the region of 5 minutes peak in the convection zone, which seems to be the dominant source of excitation of the solar oscillations (Goldreich, Murray, and Kumar, 1994). There have been some investigations on the behaviour of the acoustic spectrum in sunspots (Kumar {\\em et al.}, 2000, and references therein). It has been generally accepted that the oscillatory power decreases in sunspots (Thomas, 1984; Braun, and Duvall, 1990; Lites {\\em et al.}, 1998) and has been attributed to a variety of mechanisms, e.g. reduction in the efficiency of {\\it p}-mode excitation by turbulent convection (Goldreich, and Keeley, 1977; Goldreich, and Kumar, 1988, 1990), absorption of acoustic power (Braun, Duvall, and LaBonte, 1988; Braun, 1995; Cally, 1995; Rosenthal and Julien, 2000), and modification of {\\it p}-mode eigenfunctions by the magnetic field (Brown, 1994; Hindman, Jain, and Zweibel, 1997). The earlier study of Kumar {\\em et al.}, (2000) showed a displacement in the peak of the acoustic sepctrum in sunspots to lower frequencies. Venkatakrishnan, Kumar, and Tripathy (2001) showed that the peak of the spectrum for small portions of the solar surface varies randomly with a spread of 200 $\\mu$Hz resulting from the stochastic nature of the excitation. When several spectra were averaged, there was no discernible change in the peak within magnetic regions, and only a marginal N-S difference was detected between the hemispheres.  In this paper, we investigate the behaviour of the acoustic spectrum as a function of the underlying magnetic field using GONG+ data, which has a spatial scale of $\\sim$ 2 arc-sec per pixel. We find that the high frequency power increases with magnetic field while the low frequency power decreases, with a transition at $\\simeq$ 5 mHz. This corresponds to the predicted acoustic cut-off frequency of about 5.3 mHz at the temperature minimum for the quiet Sun in theoretical approaches to model the solar atmosphere (Gurman, and Leibacher, 1984; Balmforth and Gough, 1990). Observationally, some estimates for the quiet Sun have already been obtained (Fossat {\\em et al.}, 1992, and references therein), yielding results which lie between 5.3 and 5.7 mHz. Hindman and Brown (1998) have found a similar increase in high frequency powers using SOHO/MDI data. In section~2, we describe our analysis and present our results, while these results are compared with the SOHO/MDI results in section~3.  \\input epsf \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=4.5in\\epsfbox{fig1.ps} \\caption{Histograms representing the peak frequency of average power spectra over pixels having different values of magnetic inductions for 11 May, 2000 ({\\it dashed line}) and 12 June, 2000 ({\\it solid line}). The error-bar at the right top corner indicates $1\\sigma$ value.} \\end{center} \\end{figure}    \\input epsf \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=4.5in\\epsfbox{fig2.ps} \\caption{Plots showing the average power spectra over pixels having different values of mean magnetic inductions, which are obtained after averaging on both the days. The average magnetic induction over each set of pixels is indicated in the figure.} \\end{center} \\end{figure}   \\input epsf \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=4.5in\\epsfbox{fig3.ps} \\caption{Plots showing the ratio of powers in magnetized regions to that in quiet regions, which are obtained after averaging on both the days. The average magnetic induction over each set of pixels is indicated in the figure} \\end{center} \\end{figure} ",
        "conclusions": "It has been already noted that high frequency oscillations like 3 minute oscillations are enhanced in sunspots (Beckers and Schultz, 1972; Kneer and Uexkuell, 1983; Thomas {\\em et al.}, 1987; Horn, Staude, and Landgraf, 1997). Our results show that there is a general increase in the high frequency power. What is more significant is the fact that this enhancement is a function of the magnetic field. The most remarkable fact is the transition of behaviour from low frequency to high frequency around 5 mHz corresponding to the acoustic cut-off frequency of the quiet photosphere at temperature minimum. Our results are similar to those of Hindman and Brown (1998) at different frequencies and magnetic fields. The major difference seems to be a deficiency in power at 6 mHz as compared to Hindman and Brown (1998). For other frequency bands, even the numerical values of the power ratios are equal to the SOHO/MDI results. This augurs well for the fully deployed GONG+ system and establishes the potential of the new GONG instrument for the work in local helioseismology. A large number of active regions have to be examined before one can conclude on the deficiency in power at 6 mHz relative to the Hindman and Brown (1998) results.  In summary, it can be noted that both GONG+ and SOHO/MDI data show that the ratio of acoustic power in magnetic regions to that in the quiet Sun is a function of the magnetic field over a large range of frequencies."
    },
    "0207/astro-ph0207238_arXiv.txt": {
        "abstract": "Previous modelling has demonstrated that it is difficult to reproduce the SCUBA source counts within the framework of standard hierarchical structure formation models if the sources are assumed to be the high-redshift counterparts of local ultra-luminous infrared galaxies with dust temperatures in the range 40 -- 60 K. Here, we show that the counts are more easily reproduced in a model in which the bulk of the sub-millimetre emission comes from extended, cool (20 -- 25 K) dust in objects with star formation rates of 50 -- 100 $\\textrm{M}_{\\odot}\\textrm{yr}^{-1}$. The low temperatures imply typical sizes of $\\sim 1 (S_{850}/1\\textrm{mJy})^{1/2}$ arcsec, a factor two to three larger than those predicted using starburst-like spectral energy distributions. Low dust temperatures also imply a ratio of optical/UV to 850-$\\mu$m flux which is 30 -- 100 times smaller, for the same optical depth, than expected for objects with a hot, starburst-like SED. This may help explain the small overlap between SCUBA sources and Lyman--break galaxies. ",
        "introduction": "Studies of galaxy properties in the infrared (IR) have added much to our understanding of how and when galaxies form their stars (e.g. \\citealt{1997MNRAS.289..490R}; \\citealt{1997AJ....114...54M}; \\citealt{1998hdf..symp..219D}; \\citealt{1998ApJ...505L.103L}; \\citealt{1998ApJ...508..539P}; \\citealt{1999ApJ...517..148F}; \\citealt*{1999ApJ...521...64M}; \\citealt{1999ApJ...519....1S}). At low redshifts, the bright end of the infrared galaxy luminosity function is dominated by strongly starbursting systems, which show a rapid increase in space density towards larger redshifts (see review by \\citealt{1996ARA&A..34..749S}). The large number of bright submillimetre (sub-mm) sources that were detected by the SCUBA instrument on the JCMT nevertheless came as a surprise (\\citealt*{1997ApJ...490L...5S}; \\citealt{1998Natur.394..241H}; \\citealt*{1999ApJ...518L...5B}; \\citealt{1999ApJ...512L..87B}; \\citealt{2000AJ....120.2244E}; \\citealt{2002MNRAS.331..817S}). If the the sub-mm sources are the high-redshift equivalent of the strong starbursts observed at low redshifts, the high number counts imply that strong evolution in the space density of bright IR sources would have to continue beyond $z\\sim$ 3 -- 4.  Many authors have devised phenomenological models to reproduce the sub-mm source counts. These models assume a local luminosity function of the sources and a parametrized law for the evolution of the sources in luminosity and/or density that is constrained to fit the counts at various wavelengths (see \\citealt{2001ApJ...549..745R} and references therein).  It is nevertheless important to understand the observed sub-mm counts within the standard cosmological paradigm. Previous modelling based on semi-analytic (\\citealt{1998MNRAS.295..877G}; \\citealt{2000A&A...363..851D}) or gas dynamical \\citep{astro-ph/0107290} simulations, has encountered considerable difficulties in reproducing the brightest sub-mm sources. If the bright SCUBA sources are the high-redshift counterparts of the local ultra-luminous infrared galaxies (ULIRGs), their dust temperature should be high (40 -- 60 K). They are then inferred to be forming stars at rates of a few hundred to a thousand solar masses per year, the exact value depending on the detailed assumptions regarding extinction, initial mass function and the duration of the burst. Such extreme star formation rates are very difficult to achieve in standard hierarchical galaxy formation models. \\citet{2000A&A...363..851D} show that under these assumptions, they are only able to reproduce the sub-mm counts with a model in which essentially {\\em all} baryons present in massive dark matter haloes at redshifts $\\sim 3$ turn into stars over timescales of order $10^8$ years. They propose that high merging rates in the early Universe may be responsible for this extremely efficient conversion of gas into stars, but do not attempt to model this in detail. On the other hand, \\citet{1999MNRAS.309..715B} have modelled the merging rates of dark matter haloes in hierarchical cosmologies and find that, in order to fit the counts, the halo mass-to-infrared light ratio of a typical galaxy merger must be 200 times smaller at redshift 3 than at the present day. This is clearly an extreme requirement.  The form of the spectral energy distribution (SED) of the sub-mm galaxies is, however, a source of major uncertainty in the predictions. Recently \\citet{astro-ph/0107290} modelled the sub-mm counts using N-body plus smoothed particle hydrodynamics simulations and showed that better agreement with the data was obtained if the dust temperatures were at the lower end of the range usually assumed for starbursting galaxies. This follows suggestions by \\citet{2001ApJ...549..745R} and \\citet{cirrus2.ps} that the sub-mm emission from high-redshift galaxies could arise from cirrus-like emission of cold extended dust.  An extended distribution of cold dust may, for example, be caused by wind-driven dust outflows (\\citealt*{1990ApJS...74..833H}; \\citealt*{1999A&A...343...51A}).  In this paper, we explore the degeneracy between the inferred star formation rate and dust temperature explicitly and treat temperature as a free parameter when fitting the sub-mm counts. We demonstrate that hierarchical galaxy formation models can reproduce the sub-mm counts with moderate star formation rates if the sub-mm emission arises from cool, extended dust and discuss the implications for the size of the emitting region. We use the galaxy formation model described in \\citet*{1993MNRAS.264..201K} and \\citet{1999MNRAS.303..188K} and the diffuse cirrus emission SEDs of ERR.  In \\S\\ref{semianalytic}, we describe the galaxy formation model used to predict the star formation rates and outline the difficulties of hierarchical models in accommodating a high space density of strong starbursts out to large redshifts. \\S\\ref{SED} discusses our modelling of the SED and explores the degeneracies in the dust models. In \\S\\ref{fit}, we present our results and a comparison with observations and in \\S\\ref{discuss} we discuss their implications for the nature of the sub-mm sources. \\S\\ref{conculsions} presents our conclusions. ",
        "conclusions": "\\label{conculsions}  The sub-mm source counts predicted by hierarchical models of galaxies formation have been studied for a wide range of SEDs.  We confirm the results of previous studies that, with plausible assumptions about star formation timescales, star formation efficiencies and IMF, hierarchical models for galaxy formation under-predict the number of bright sub-mm sources by a large factor if hot (40 -- 60 K) starburst-like SEDs are assumed.  The sub-mm source counts alone do not constrain the dust temperatures and SEDs of the sources. Dust temperatures and inferred star formation rates are highly degenerate. Lower dust temperatures require substantially lower star formation rates to produce the same 850-$\\mu$m flux.  For dust temperatures in the range 20 -- 25 K, the observed sub-mm counts agree very well with those predicted by hierarchical galaxy formation models. These low temperatures imply typical radii that correspond to angular sizes of $1 (S_{850}/1\\textrm{mJy})^{1/2}$ arcsec if we assume an extended distribution of dust and stars  with $A_V=10$. This is a factor two to three larger than predicted for starburst-like SEDs. Observational limits on the size of sub-mm sources are scarce. Marginally resolved observations of the SCUBA galaxy Lockman850.1 with the Plateau de Bure interferometer, and follow up near infrared imaging, may indicate that the sub-mm sources are as large as predicted by our cold extended dust models.  The typical redshift of the sources should then lie in the range $2\\lesssim z\\lesssim 4$, somewhat, but not significantly, smaller than in models which assume a starburst-like SED. These low temperatures substantially reduce the predicted mid-IR emission of the sub-mm sources. If our models are correct, SCUBA-selected sources may contribute less than expected to the counts in upcoming mid-IR surveys, which will be dominated by sources with hotter dust.  The low temperatures required to reconcile the sub-mm counts with the predictions of hierarchical galaxy formation models are not plausible for compact star-bursting galaxies, but they are expected if the emission comes from a more extended distribution of stars and dust. If most sub-mm galaxies have indeed extended distributions of cold dust with large sub-mm to UV/optical ratios this would naturally explain the small overlap with the population of Lyman--break galaxies, which is biased towards high surface brightness regions in actively star-forming galaxies."
    },
    "0207/astro-ph0207462_arXiv.txt": {
        "abstract": "The $\\lambda$~Bootis stars, a group of late B to early F-type population\\,{\\sc I} stars, have surface abundances that resemble the general metal depletion pattern found in the interstellar medium. Inspired by the recent result that the fundamental parameters of these peculiar stars differ in no respect from a comparison sample of normal stars, the hypothesis of an interaction between a star and a diffuse interstellar cloud is considered as a possible explanation of the peculiar abundance pattern. It is found that such a scenario is able to explain the selective accretion of interstellar gas depleted in condensable elements as well as the spectral range of the $\\lambda$~Bootis phenomenon. ",
        "introduction": "The $\\lambda$~Bootis stars are late B to early F-type population\\,{\\sc I} stars, which show a peculiar surface abundance pattern: while the light elements (C, N, O and S) are roughly solar, the Fe-peak elements show underabundances of up to a factor of 100. \\citet{Venn} are the first who noticed the similarity between this abundance pattern and the depletion pattern of the interstellar medium (ISM) and suggested the accretion of interstellar or circumstellar gas to explain the $\\lambda$~Bootis stars. In this respect they differ from the rest of the peculiar A-type stars where the abundance pattern is caused by separation processes in the stellar atmosphere itself. The abundances are ascribed to diffusion in the presence of slow rotation (Am stars) or strong magnetic fields (Ap stars).  \\citet{Paunzen02} scrutinized the available observational data to put constraints on any model trying to explain the $\\lambda$~Bootis phenomenon. A comparsion between the $\\lambda$~Bootis stars and a reference sample of normal stars showed that both groups share the same fundamental parameters, like effective temperature, gravity, mass, rotational velocity and age. But most surprisingly, the Na abundance of the $\\lambda$~Bootis stars revealed a correlation with nearby local interstellar column densities of Na\\,{\\sc i}.  This discovery, although so far tentative, because of the inhomogeneity of the stellar Na abundances, motivated a detailed analysis of the interaction between a star and a diffuse ISM cloud as an explanation for the $\\lambda$~Bootis phenomenon. ",
        "conclusions": ""
    },
    "0207/astro-ph0207181_arXiv.txt": {
        "abstract": "A pulse timing orbit has been obtained for the X-ray binary \\src\\ using observations made with the Proportional Counter Array on board the Rossi X-ray Timing Explorer. The mass function obtained of \\sqig 16M\\sun\\ together with the detection of an extended near-total eclipse confirm that the primary star is a supergiant as predicted. The orbital eccentricity is found to be very low with a best fit value of 0.04 $\\pm$\\ 0.02. The orbital period is also refined to be 6.0724 $\\pm$ 0.0009 days using an improved and extended light curve obtained with RXTE's All Sky Monitor. Observations with the ASCA satellite provide an improved source location of R.A. = 18$^{h}$ 55$^{m}$ 31.3$^{s}$, decl. = -02\\degrees\\ 36\\arcmin\\ 24.0\\arcsec\\ (2000) with an estimated systematic uncertainty of less than 12\\arcsec. A serendipitous new source, AX J1855.4-0232, was also discovered during the ASCA observations. ",
        "introduction": "The X-ray source, \\src, was discovered during Rossi X-ray Timing Explorer (RXTE) scans along the galactic plane  (Corbet et al. 1999; hereafter Paper I) . The source showed pulsations at a period of 361 s and a light curve obtained with RXTE's All Sky Monitor (ASM) showed modulation at a period of 6.067$\\pm$ 0.004 days which was interpreted as the orbital period of the system. The X-ray spectrum above \\sqig3 keV could be fitted with an absorbed power law model with a high-energy cut-off, and an iron emission line at approximately 6.4 keV. These results, in particular the location of the source in the orbital period/spin period diagram (Corbet 1986), were interpreted as indicating that \\src\\ is likely to consist of a neutron star accreting from the wind of an O or B supergiant primary. A less likely interpretation was that \\src\\ is instead a Be/neutron star binary, in which case it would have an unusually short orbital period for such a system.  Here we present the results of observations made with the RXTE Proportional Counter Array (PCA) that were performed over the course of one complete orbital cycle. The observations were performed with the aims of (i) measuring the orbital parameters to determine the X-ray mass function and thus the nature of the primary star, and (ii) determining whether a eclipse is present in the light curve. A system containing a supergiant primary rather than a main-sequence Be star would be much more likely to exhibit an eclipse due to the much greater size of the primary star. We also report on observations made with the imaging detectors onboard the ASCA satellite which enable the source position to be refined. RXTE ASM observations utilizing this improved position and extending over six years allow further refinement of the orbital period. Spectroscopic results from both satellites are not discussed here and will be presented elsewhere. ",
        "conclusions": "The light curve and pulse timing orbit clearly show \\src\\ to be a supergiant X-ray binary as predicted in Paper I. With the detection of the eclipse and timing measurements over an entire orbit the system parameters can now be determined. Future pulse timing observations would enable a search for orbital period changes as seen in some other high mass X-ray binaries (e.g. Clark 2000, Levine, Rappaport, \\& Zojcheski 2000, and references therein).  If an optical or IR counterpart could be found and its radial velocity orbit measured this would be valuable as the system would then be a ``double-lined\" eclipsing binary and the neutron star mass could be directly determined. While the optical reddening to this object implied by the measured X-ray absorption is high (N$_H$ = 15$\\times$10$^{22}$ cm$^{-2}$ $\\Rightarrow$ E(B-V) = 24, Paper I) at least some of this absorption may be local to the X-ray source rather than genuinely interstellar. Tighter constraints on the eclipse duration will also be valuable in obtaining precise measurements of the system parameters."
    },
    "0207/astro-ph0207148_arXiv.txt": {
        "abstract": "The amount and nature of the evolution of the X-ray properties of clusters of galaxies provides information on the formation of structure in the universe and on the properties of the universe itself.  The cluster luminosity - temperature relation does not evolve strongly, suggesting that the hot X-ray gas had a more complicated thermodynamic history than simply collapsing into the cluster potential well. Cluster X-ray luminosities do evolve. The dependence of this evolution on redshift and luminosity is characterized using two large high redshift samples. Cluster X-ray temperatures also evolve. This evolution constrains the dark matter and dark energy content of the universe as well as other parameters of cosmological interest. ",
        "introduction": "The evolution of cosmic structure is strongly dependent on the cosmology of the universe. Jenkins et al. (1998) is one of the many papers describing this well known result. Clusters of galaxies, as the most massive bound objects known, are the ultimate manifestations of cosmic structure building. The evolution of clusters is simple, being driven by the gravity of the underlying mass field of the universe and of a collisionless collapse of cluster dark matter. It should be possible to calculate this evolution reliably compared to that of other objects visible at cosmological distances such as galaxies or AGN. Clusters are luminous X-ray sources. The X-ray emission mechanism is optically thin thermal radiation from a medium nearly in collisional equilibrium, about the simpliest situation imaginable. Thus observations of the X-ray evolution of clusters provide a robust measure of the evolution of cosmic structure and thereby constrain the cosmology of the universe.  The cosmological model is described by a set of cosmological parameters. The present value of the Hubble parameter is $\\mathrm{H_0 \\equiv 100\\:h\\:km\\:s^{-1}\\:Mpc^{-1}}$. When needed h = 0.5 will be used, but almost nothing about evolution depends on the precise value of h since data at two epochs are always compared. The present matter and dark energy densities in terms of the critical density are $\\Omega_{m0}$ and $ \\Omega_{\\Lambda0}$ respectively. The amount of structure in the universe is described by $\\sigma_8$, the present rms matter fluctuations in spheres of $\\mathrm{8\\:h^{-1}\\:Mpc}$. This parameter is a complicated way to describe the present normalization of the spatial power spectrum of matter density fluctuations, P(k), on a scale of k $\\sim\\:0.2 \\mathrm{\\:h\\: Mpc^{-1}}$: $\\mathrm{\\sigma_8 \\approx [P(0.172\\:h\\:Mpc^{-1})/3879\\:h^{-3}\\:Mpc^3]^{1/2}}$ (Peacock, 1999 equations 16.13 and 16.132). The dark energy equation of state is P = w $\\rho$ c$^2$. If w = -1, then the dark energy is the cosmological constant, if $\\mathrm{-1 < w < 0}$ it is termed Quintessence. Recall that w = 0 is cold dark matter and w = 1/3 is radiation. Cluster temperature evolution provides constraints on all of the above cosmological parameters except h.  When needed, two specific cases will be considered. The X-ray astronomer's universe where $\\Omega_{m0} = 1.0, \\Omega_{\\Lambda0} = 0.0$. This combination was known to be correct ten years ago. The other case may be called the bandwagon universe where $\\Omega_{m0} = 0.3, \\Omega_{\\Lambda0} = 0.7$. This combination is known to be correct today. Both universes are spatially flat, $\\Omega_{m0} + \\Omega_{\\Lambda0} = 1$. ",
        "conclusions": "The L - T relation shows that the hot X-ray cluster gas experienced a more complicated thermodynamic history than that resulting from the cluster collapse. There was possibly some heating prior to collapse and/or cooling with lesser heating after collapse.  There is no longer a question whether cluster X-ray luminosities evolve, they do. Seven nearly independent surveys show some evolution and systematics do not appear to be a big effect. Large high redshift samples, i.e. those containing more than 100 clusters, are needed to characterize the evolution. Two such samples currently exist. HEN shows a factor of 2 evolution compared to REFLEX at L(0.5,2.0) $= 2\\times10^{44}$ $\\mathrm{erg\\:s^{-1}}$ over the redshift interval 0.11 to 0.45 ($\\Omega_{m0} = 1.0$, $\\Omega_{\\Lambda0}= 0.0$). MACS shows a factor of 10 evolution compared to eBCS at L(0.1,2.4) $> 10^{45}$ $\\mathrm{erg\\:s^{-1}}$ over the redshift interval 0.05 to 0.55 ($\\Omega_{m0} = 1.0$, $\\Omega_{\\Lambda0}= 0.0$).  A cosmology with $\\Omega_{m0} = 0.4$, $\\Omega_{\\Lambda0}= 0.6$ is consistent with cluster X-ray temperature evolution, the cosmic microwave background and supernovae at the $\\sim1\\sigma$ level. The systematics of cluster temperature evolution constraints on cosmology are a minor concern at present, i.e. they are about the same size as the statistical errors.  However, this method is so statistically powerful that systematics will be the dominate effect as soon as the sample sizes become only a factor of two or three larger, at least for the determination of $\\sigma_8$.  A more positive statement would be that an X-ray cluster survey of 10,000 deg$^2$ to a flux of $\\mathrm{F(0.5,2.0) > 5\\times10^{-14} \\:erg\\:cm^{-2}\\:s^{-1}}$ would yield $\\sim18,000$ clusters to z $\\sim1.5$ and provide constraints of similar statistical quality as the upcoming Planck and the proposed SNAP missions (Petre et al., 2001). In fact the great complementarity of cluster evolution, cosmic microwave background, and supernovae constraints exhibited in Figure 10 would provide a check on the systematics of all three methods."
    },
    "0207/astro-ph0207654_arXiv.txt": {
        "abstract": "{We  extend  our study  of  the  nuclei  of 3CR  FR~II  radio galaxies through HST optical images up  to $z = 0.3$.  In the majority of them  an unresolved nucleus  (central compact core, CCC)  is found. We analyze their position in the plane formed by the radio and optical nuclear luminosities in relation to their optical spectral properties. The  broad--lined objects (BLO)  have the  brightest nuclei:  they are present  only  at  optical  luminosities  $\\nu  L_{\\nu}  \\gta  4\\times 10^{42}$ erg s$^{-1}$ which we  suggest might represent a threshold in the radiative  efficiency combined  with a small  range of  black hole masses.  About  $40 \\%$ of the  high and low  excitation galaxies (HEG and LEG)  show CCC  which resemble those  previously detected  in FR~I galaxies,  in  apparent  contrast   to  the  unification  model.   The equivalent  width of  the [OIII]  emission line  (with respect  to the nuclear  luminosity) reveals  the nature  of these  nuclei, indicating that the  nuclei of  HEG are obscured  to our  line of sight  and only scattered radiation is observed.   This implies that the population of FR~II is  composed of objects  with different nuclear  properties, and only   a   fraction   of   them   can   be   unified   with   quasars. ",
        "introduction": "In  the  framework  of  the  AGN unification  scheme  for  radio--loud sources,   powerful  radio   galaxies   with  FR~II   edge--brightened morphology (Fanaroff \\& Riley \\cite{fr}) are believed to be misaligned quasars, while lower power, edge--darkened FR~I are associated with BL Lac objects.  This  basic unification picture (Barthel \\cite{barthel}, Orr \\& Browne \\cite{browne}, see Urry \\& Padovani \\cite{urrypad} for a review)  is  mainly  supported  by  the  comparison  of  the  extended properties (radio  morphology and linear dimensions,  host galaxy type and environment,  and narrow emission  line luminosity) as well  as by the  number counts  of the  two  classes.  However,  the detection  of polarized broad  emission lines, although  only for a small  number of FR~II, is the  most direct evidence of the  unification scheme and the presence  of  absorbing ``tori''  in  high  power  radio loud  sources (Antonucci   \\&    Barvainis   \\cite{antonucci90},   Cohen    et   al. \\cite{cohen99}).  How these  results extend  to low power  sources is still  unclear. In fact, on large scales,  the morphological FR~I/FR~II dicothomy appears to be also  associated with other (large scale)  properties.  From the optical point of view, FR~II are associated with different sub-classes of  bright elliptical  galaxies.   On average,  FR~II  hosts are  less luminous with respect  to FR~I ones (Owen \\cite{owen},  Ledlow \\& Owen \\cite{ledlowowen}), and  belong to lower  density groups, at  least at low redshifts (e.g.  Zirbel \\cite{zirb96}, \\cite{zirb97}).   However,  how  these large--scale  properties  relate  to the  nuclear structure  and  central  activity   is  still  a  debated  issue.   In particular,  the properties  of  the emission  lines  observed in  the spectrum  of radio  galaxies and  plausibly connected  to  the nuclear activity, have revealed a phenomenology richer than the radio one. For example,  the role  of a sub-class  of low-ionization  FR~II (e.g. Laing et al.  \\cite{laing94}) has still to be assessed.  These objects have an  FR~II morphology, but  their optical spectral  properties are similar to those of  FR~I.  Wall \\& Jackson (\\cite{jacksonwall97}) and Jackson  \\&  Wall (\\cite{jacksonwall99})  proposed  that such  objects constitute, together with FR~I, a single population of radio galaxies. Furthermore, Willott et al.  (\\cite{willott}) have recently found that the fraction of  objects with observed broad lines (in  the 6C, 7C and 3CRR  samples and  having  excluded FR~I)  decreases with  luminosity. They indicate as  a possible explanation for this  lack of quasars the rise of  a distinct population of  radio galaxies which  have an FR~II radio morphology but lack a well-fed quasar nucleus.  Our  aim is to  investigate these  issues by  directly looking  at the nuclear  continuum  emission  in  the optical  band,  identifying  its physical origin  and relating it to  both the radio  and emission line properties.   HST  optical images  of  the  nuclear  regions of  radio galaxies are suited to this goal as their high resolution allows us to separate the AGN emission from the stellar host galaxy background.  In particular, the  optical snapshot surveys of 3CR  objects (e.g. Martel et al.  \\cite{martel},  De Koff et al.  \\cite{dekoff})  has provided a wealth of high quality data for this purpose.  We have started this study by considering complete samples of FR~I and FR~II from this catalog (Chiaberge et al. \\cite{pap1}, hereafter Paper~I; Chiaberge et al.  \\cite{pap2}, hereafter Paper~II).  We analyzed the properties of unresolved optical nuclei, which have been found to be present in the great majority of the objects.  The  optical  nuclei  of  3CR  FR~I behave  similarly,  and  are  best explained as  non thermal  synchrotron emission from  the base  of the relativistic jet.  Furthermore, the  high detection rate ($\\sim 85\\%$) directly  implies  that geometrically  thick  obscuring  tori are  not present in  FR~I radio galaxies  (or, alternatively, they  are present only in a  minority of them).  Given this, the  lack of broad emission lines in these objects cannot be due to obscuration.  The behavior of FR~II (at redshift below $z=0.1$) appears to be more complex, although their properties are clearly related to their spectral classification. Broad line radio galaxies show an optical excess with respect to the expected non thermal emission level, which might indicate a contribution from the thermal disk.  Several radio galaxies in which broad lines are absent do not show any nuclear source, and they can be interpreted as obscured nuclei, as expected in the framework of the current AGN unification scheme.  Most importantly, 5 sources of the sample have a core with radio-optical properties that are completely consistent with those found in FR~I.  Although not all of them belong to the low--ionization subclass, it is tempting to consider them as FR~I--like.  In all cases, this nuclear emission sets an upper limit to any radiation from the accretion flow. In particular, in the non--thermally dominated nuclei (FR~I or FR~II) this seems to imply that accretion might take place on a low efficiency radiative regime.  Since it is important both to establish whether these findings hold only for nearby sources, and to analyze a larger sample of objects in order to improve the statistics, in this paper we extend the sample up to a redshift of $z=0.3$. Furthermore we will show that in order to address the nature of the nuclei, a crucial parameter is the equivalent width of the [OIII] emission line, which we calculate with respect to the nuclear continuum emission.  The organization of the paper is as follows: in Sect.  \\ref{sample} we describe our sample of FR~II and the HST observations; in Sect.  \\ref{fr2ccc} we present the results of the photometry of the nuclei and we analyze the relation between the optical and radio core luminosity; in Sect. \\ref{discussion} we discuss our results for the different spectral subclasses, also considering their radio properties and, most importantly, the [OIII] emission line luminosity.  In Sect. \\ref{conclusions} we present a summary of our findings and we draw conclusions and future perspectives.  $H_0= 75$ km s$^{-1}$  Mpc$^{-1}$ and $q_0=0.5$ are adopted throughout the paper.  The  spectral index $\\alpha$ is defined  as $F_\\nu \\propto \\nu^{-\\alpha}$. ",
        "conclusions": "\\label{conclusions}  We have analyzed  the optical nuclear properties of  a complete sample of 65  FR~II radio galaxies up  to $z=0.3$ from the  3CR catalog.  The overall  scenario  basically  confirms  the  findings  for  the  lower redshift   sample   ($z<0.1$)    presented   in   Chiaberge   et   al. (\\cite{pap2}).   However,  the larger  number  of  sources (more  than double) allows us to reveal  a richer and more complex behavior, which turns  out  to  be   closely  associated  with  the  optical  spectral properties of the different objects.  The nuclear properties of FR~II, as inferred from our analysis, can be summarized as follows.  While the  great majority of  FR~I radio galaxies (optical  and radio) nuclei  are dominated  by  non-thermal synchrotron  emission from  the relativistic jet,  the FR~II population is  not homogeneous.  Although optical Central  Compact Cores appear to  be a common  feature also in FR~II  galaxies, their origin  can be  ascribed to  different physical processes.  BLO typically  have the  brightest nuclei which  show a  large optical excess with  respect to the radio-optical correlation  found for FR~I. This is  readily explained  if the dominant  component in  the optical band is due to thermal emission from an accretion disc.  We found that optical nuclei of BLO are present only for $L_{{\\rm o}} > 10^{28}$ erg s$^{-1}$  Hz$^{-1}$.    A  rather  different  behavior   is  seen  in radio-quiet AGN,  in which broad lines  are seen in  objects that span many orders of magnitude in nuclear optical luminosity, from LINERS to powerful  QSO. As we  argued that  this effect  cannot be  ascribed to obscuration  or  selection  effects,  we  suggest  that  this  is  the manifestation  of  a threshold  in  the  efficiency  of the  accretion process,  from  the   standard  optically  thick,  geometrically  thin accretion disk to low radiative accretion flows.  Note that the presence of any limit in luminosity requires a well defined behavior of the accretion rate and radiative efficiency but also a narrow distribution in black hole masses for radio-loud AGNs. This conjecture seems indeed to be supported by the direct measurements available to date.  The nature of  the nuclei of HEG is  certainly more complex.  Although most sources  lie along the FR~I  correlation, there are  at least two clear exceptions, showing  a significant optical excess.  Furthermore, there  are  at  least  two  sources  (3C~234  and  3C~109)  for  which spectropolarimetric studies  clearly showed that they  harbor a hidden QSO nucleus.  The amount  of scattered continuum light matches exactly the flux of their nuclear component in the HST images, indicating that their  CCC  is  a   compact  scattering  region.   Nonetheless,  their representative  points would  lie on  the correlation  by coincidence. From this  analysis it turns out  that in general it  is impossible to definitively  address the  nature of  HEG nuclei  by  only considering their position in the  optical--radio plane.  A fundamental advance is achieved with  the inclusion in  our analysis of a  further parameter, i.e. the luminosity of their emission lines.  In the plane formed by the ``nuclear'' EW of the [OIII] line vs the optical excess with respect to the non-thermal jet emission (a new ``fundamental'' diagnostic  plane?), the different classes of sources clearly separate according to the nature of their nuclei.  For low EW values ($\\sim 10^{2.5}$ \\AA) we find all QSO, WQ, LEG and FR~I, which differ only by the amount of the optical excess.  On the other hand, all but two of the HEG have much larger equivalent widths \\gta $10^{3.5}$ \\AA. The separation of sources depending on their line equivalent width is indeed expected from the unified models, as obscuration reduces the nuclear continuum emission while the line emission is less affected or unaffected.  In sources with very high values of EW a strong ionization source, obscured to our viewing angle, must be present.  We can then argue that all sources with high EW are hidden BLO.  The low EW region would then contains objects in which we see directly the source  of ionization.  A ramification  of this result  is that in FR~I the most likely dominant  source of gas ionization is synchrotron emission from the jet.  Only   two  HEG   are  located   among  the   ``unobscured''  sources: interestingly, these galaxies (3C~18 and  3C~349) also lie on the FR~I correlation.   Therefore, due  to their  position in  both  planes, we identify these objects with true unobscured ``FR~I--like'' FR~II.  According to the scenario proposed  here, the non--detection of CCC in galaxies of  different classes should have different  origins. For LEG this  might be  due to  either a  low contrast  with the  stellar host galaxy  emission or  to a  (moderate) amount  of  absorption, randomly oriented with respect to the jet.  On the other hand, for HEG this can be only attributed to a lower amount of scattered nuclear radiation.  The  picture which  emerges is  that  radio galaxies  manifest in  two types,  which are  not  directly related  to  the extended  FR~I/FR~II dicothomy.   In  the framework  of  the  unification  scheme, BLO  and obscured  HEG  appear to  have  the  same  nuclear structure:  intense thermal  disk  (ionizing) emission,  substantial  broad emission  line region,  torus-like absorber and,  of course,  powerful jets.   On the other  hand,  LEG,  FR~I  and  unabsorbed HEG  constitute  a  distinct population, characterized  by low radiative  efficient accretion, weak or absent broad line emission, lack of a significant nuclear absorbing structure. Unfortunately we  have complete information (radio, optical and emission line) for only roughly half of the sources in our sample. Although we can  {\\it only} make the assumption that  there is no bias in their  selection, we can estimate  that the population  of FR~II is composed of  $\\sim 50\\%$ obscured sources harboring  a quasar nucleus, $\\sim  25\\%$ BLO,  $\\sim  20\\%$ LEG  and  $\\sim 5  \\%$ FR~I--like  HEG (although so far  this is assessed only for  two sources).  The latter two classes can account for BL Lac objects with FR~II radio morphology and extended  radio power (e.g. Kollgaard  et al.  \\cite{kollgaard92}, Cassaro  et  al.   \\cite{cassaro}).   This scenario  apparently  poses problems  for  the  simplest  unification  models,  in  particular  in identifying  the  beamed  counterparts  of FR~I--like  HEG.   However, strong and  high excitation narrow emission lines  are indeed observed in  few BL  Lacs (Landt  et  al.  \\cite{landt}).   Moreover, the  line equivalent width  in the  beamed objects will  be reduced by  a factor $\\sim 10^4$ (the typical  ratio between the nuclear optical luminosity between   radio   galaxies  and   BL   Lacs,   see   Capetti  et   al. \\cite{capetti02})   producing  values   consistent  with   a   BL  Lac classification.  A possible way to test the  overall picture is to look at the spectral properties of the nuclei. BLO  and scattered nuclei are expected to be different  from   the  FR~I   and  FR~I--like  synchrotron   ones:  in particular, we expect to  observe flatter spectral indices, typical of quasars, in BLO, indicating the  presence of a thermal blue bump.  Due to  the  large  uncertainties,   optical  observations  (even  in  two different bands) are not enough to determine the spectral slope, which instead  could  be better  measured  by  taking  advantage of  the  UV information (Chiaberge et al. \\cite{papuv}).  In addition, an infrared nuclear excess is expected in  the obscured radio galaxies, while this has  to be  absent  in  FR~I--like objects,  as  promisingly shown  by Whysong \\& Antonucci  (\\cite{whysong}) for the case of  3C~405 (a true obscured quasar)  and 3C~274 (M~87).  A further diagnostic tool  is of course  the observations  in the  X--ray  band, which  can reveal  the presence of different amounts of nuclear absorption.  Finally,  it would  be  particularly interesting  to  analyze how  the properties  of  the  newly   discovered  quasars  showing  FR~I  radio morphology  (Blundell   \\&  Rawlings  \\cite{blundraw},   Lara  et  al. \\cite{lara}) fit in  our picture. In a sense  they might represent the analogous (but  oposite) case of LEG,  where an FR~II  harbors an FR~I nucleus.   The spectral  information available  in the  literature for such objects  is, to our best  knowledge, not yet  sufficient to state where  these   sources  are  located   in  the  diagnostic   plane  of Fig.~\\ref{ew}. However, if they  are indeed broad--lined FR~I (showing optical  thermal emission  from  the accretion  disk)  they should  be placed in the  lower--right end of the plane, among  the BLO.  If this is  the case,  the properties  of these  peculiar sources  might bring further support to the models that claim that the nuclear structure is not  directly  connected  to   the  extended  radio  morphology  (e.g. Bicknell   \\cite{bicknell84,bicknell94};    Gopal-Krishna   \\&   Wiita \\cite{wiita})."
    },
    "0207/hep-ph0207157_arXiv.txt": {
        "abstract": "Prompted by recent solar and atmospheric data, we re-analyze the four-neutrino description of current global neutrino oscillation data, including the LSND evidence for oscillations. The higher degree of rejection for non-active solar and atmospheric oscillation solutions implied by the SNO neutral current result as well as by the latest 1489-day Super-K atmospheric neutrino data allows us to rule out (2+2) oscillation schemes proposed to reconcile LSND with the rest of current neutrino oscillation data. Using an improved goodness of fit (\\gof) method especially sensitive to the combination of data sets we obtain a \\gof\\ of only $1.6\\times 10^{-6}$ for (2+2) schemes. Further, we re-evaluate the status of (3+1) oscillations using two different analyses of the LSND data sample. We find that also (3+1) schemes are strongly disfavoured by the data. Depending on the LSND analysis we obtain a \\gof\\ of $5.6\\times 10^{-3}$ or $7.6\\times 10^{-5}$. This leads to the conclusion that all four-neutrino descriptions of the LSND anomaly, both in (2+2) as well as (3+1) realizations, are highly disfavoured. Our analysis brings the LSND hint to a more puzzling status. \\begin{keyword} neutrino oscillations \\sep sterile neutrino \\sep four-neutrino models \\PACS 14.60.P \\sep 14.60.S \\sep 96.40.T \\sep 26.65 \\sep 96.60.J \\sep 24.60 \\end{keyword} ",
        "introduction": "The atmospheric neutrino data~\\cite{atm-exp,skatm,macro}, including the most recent 1489 Super-K data sample provide strong evidence for $\\nu_\\mu$ oscillations into an active neutrino (mainly $\\nu_\\tau$) with $\\dma\\sim 2 \\times 10^{-3}~\\eVq$~\\cite{solat02}. On the other hand, apart from confirming, once again, the long-standing solar neutrino problem~\\cite{sksol,chlorine,sage,gallex_gno,sno01}, the recent results from the Sudbury Neutrino Observatory (SNO) \\cite{sno02} have given strong evidence that solar neutrinos convert mainly to an active neutrino flavor.  This suggests that an extension of the Standard Model of particle physics is necessary in the lepton sector, capable of incorporating the required $\\nu_e$ conversion. Although certainly not yet unique \\cite{Miranda:2000bi,Guzzo:2001mi}, the most popular explanation of the solar neutrino data is provided by the active large mixing angle (LMA) neutrino oscillation hypothesis, characterized by a neutrino mass-squared difference $\\dms\\lesssim 10^{-4}~\\eVq$~\\cite{solat02}.  In contrast, reactor and accelerator neutrino data \\cite{KARMEN,bugey,CHOOZ,PaloV,CDHS} lead to negative results. However, the LSND experiment~\\cite{LSND,LSND2001} has provided positive results, which may or may not be confirmed by the forth-coming MiniBooNE experiment~\\cite{MiniBooNE}.  The required neutrino mass-squared difference $\\dml \\gtrsim 0.2$ eV$^2$ is in conflict with the neutrino mass-squared differences indicated by solar and atmospheric data in a minimal three-neutrino picture. Four-neutrino models \\cite{ptv92,pv93,cm93,4nuModels,4nuextra,Ioannisian:2001mu,Hirsch:2000xe} potentially account for all current oscillation data. The status of four-neutrino descriptions has been presented in Ref.~\\cite{Maltoni:2001bc}. An exhaustive list of four-neutrino references can be found in Ref.~\\cite{giuntiwebp}.  \\begin{figure}[t] \\centering \\includegraphics[width=0.65\\linewidth]{figures/schemes.eps} \\caption{\\label{fig:4spectra}% The six types of four-neutrino mass spectra. The spacings in the vertical axes correspond to the different scales of mass-squared differences required in solar, atmospheric and short baseline neutrino oscillations.} \\end{figure}  The six possible four-neutrino mass spectra are shown in Fig.~\\ref{fig:4spectra}.  For the case $\\dml \\gg \\dma, \\dms$, which we tacitly assume in order to reconcile the LSND anomaly with solar and atmospheric data, these schemes can naturally be divided into two very different classes, usually called (3+1) and (2+2) mass schemes~\\cite{barger00}. It is important to note that (3+1) mass spectra include the three-active neutrino scenario as limiting case. In this case solar and atmospheric neutrino oscillations are explained by active neutrino oscillations, with mass-squared differences $\\dms$ and $\\dma$, and the fourth neutrino state gets completely decoupled.  We will refer to such limiting scenario as (3+0).  In contrast, the (2+2) spectrum is intrinsically different, as there must be a significant contribution of the sterile neutrino either in solar or in atmospheric neutrino oscillations or in both.  In recent studies~\\cite{Maltoni:2001bc} it has been realized that there is considerably tension in four-neutrino fits to the global data. In such global four-neutrino analyses one is faced with the problem that there are different data sets, which all give very good fits if analyzed separately. Problems arise due to the {\\sl combination} of the different data sets in a four-neutrino framework.  In this letter we re-evaluate the status of four-neutrino interpretations of the LSND anomaly in the light of the recent solar \\cite{sksol,sage,sno02} as well as atmospheric \\cite{skatm} neutrino results.  To evaluate the quality of the fit we will apply appropriate statistical methods, which are especially sensitive to the combination of different data sets in a global analysis. We find that both the SNO NC result as well as the 1489-day Super-K atmospheric neutrino data strongly reject against sterile neutrino conversions. This essentially rules out (2+2) mass schemes. We also re-evaluate the status of (3+1) schemes by considering two different analyses of LSND data~\\cite{LSND2001,Church:2002tc}. We will further elaborate the result of previous studies \\cite{barger00,3+1early,BGGS,peres,carlo,GS,Maltoni:2001mt} that in (3+1) schemes the LSND signal is in serious disagreement with bounds from short-baseline (SBL) experiments reporting no evidence for oscillations \\cite{KARMEN,bugey,CDHS}. The net result is that all four-neutrino descriptions of the LSND anomaly, both in (2+2) as well as (3+1) realizations, are strongly disfavoured by the data. This brings the LSND anomaly to a theoretically more puzzling status. We note that also cosmology put strong constraints on four-neutrino schemes (for recent analyses see Ref.~\\cite{cosmology}).   The plan of the paper is as follows. In Sec.~\\ref{sec:four-neutr-oscill} we briefly describe our parameterization and approximations used in the global four-neutrino analysis.  In Secs.~\\ref{sec:solar} and \\ref{sec:atmosph} we summarize the solar and atmospheric neutrino data and their analysis \\cite{solat02}. In Sec.~\\ref{sec:sbl} we describe the SBL data we are using, and we compare the two different analyses of the LSND data~\\cite{LSND2001,Church:2002tc}. In Sec.~\\ref{sec:2+2} we show that (2+2) oscillation schemes are ruled out because of the tension between solar and atmospheric data, whereas in Sec.~\\ref{sec:3+1} we present our re-analysis of the disagreement in SBL data in (3+1) oscillation schemes in light of the two LSND samples.  In Sec.~\\ref{sec:comparing} we show that the goodness of fit is very bad in all four-neutrino cases. Furthermore, we compare the relative quality of the fit for the schemes (3+1) and (2+2), as well as the three active neutrino case (3+0).  The quantitative statistical criteria to analyze the conflict between different data sets, which are used in Secs.~\\ref{sec:2+2}, \\ref{sec:3+1} and \\ref{sec:comparing}, are formulated in the appendix \\ref{appendix}. In summary, we find that all four-neutrino descriptions of the LSND anomaly both in (2+2) as well as (3+1) realizations are highly disfavoured due to recent data, as mentioned in our conclusions, Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Prompted by the recent SNO neutral current result and recent atmospheric data, we have re-analyzed the global status of all current neutrino oscillation data in the framework of four-neutrino schemes, in a way similar to that of Ref.~\\cite{Maltoni:2001bc}. Our present update includes the global solar neutrino data, the most recent sample of atmospheric neutrino experiments, and a detailed treatment of experiments reporting no evidence (NEV) for oscillations (KARMEN, Bugey, CDHS as well as information from CHOOZ and Palo Verde). Besides the analysis of the global LSND data, we also perform an analysis based mainly on the decay-at-rest (DAR) data sample, which provides a higher sensitivity to the oscillation signal in LSND. We have identified the parameter consistency (PC) test and the parameter goodness of fit (PG) as useful statistical methods to evaluate the compatibility of different data sets in a global analysis.  We have found that in (2+2) mass schemes recent solar and atmospheric data are compatible only at the 3.5$\\sigma$ according to the PC test, and at 4.8$\\sigma$ according to the PG method.  In (3+1) mass schemes the disagreement of LSND with the rest of the oscillation data has been evaluated for both LSND analyses as a function of $\\dml$. For the global LSND data PC is achieved only at 2$\\sigma$ and the PG is below 1\\% for all values of $\\dml$. Using the LSND DAR sample the disagreement is always stronger: we find a PC only at 99\\% \\CL\\ and for most values of $\\dml$ the PG is worse than 4$\\sigma$.  We have evaluated the \\gof\\ of a global fit in terms of the PG by dividing the data in SOL, ATM, LSND and NEV samples.  In the best case we find a PG of only 0.56\\%. This value occurs for the (3+1) mass scheme and the global LSND data. Using the LSND DAR analysis we get a PG of $7.6\\times 10^{-5}$ for (3+1). For (2+2) oscillation schemes the situation is worse: we find a very bad PG of $1.6\\times 10^{-6}$ for LSND global and $3.5\\times 10^{-7}$ for LSND DAR.  Concerning the relative status of the hypotheses (3+1), (2+2) and the three-active neutrino case (3+0) we find that for the LSND global data (2+2) and also (3+0) are strongly disfavoured with respect to (3+1) with a $\\Delta\\chi^2= 17.8$ and $\\Delta\\chi^2= 20.0$, respectively. For LSND DAR we find for (2+2) a $\\Delta\\chi^2 = 11.5$ relative to (3+1), and the high significance of the oscillation signal condensed in the DAR sample leads to the huge value of $\\Delta\\chi^2 = 28.5$ of (3+0) relative to (3+1).  The exclusion of four-neutrino oscillation schemes of the (2+2)-type is based on the improved sensitivity of solar and atmospheric neutrino experiments to oscillations into a sterile neutrino, thanks to recent experimental data. This is a very robust result, independent of whether the LSND experiment is confirmed or disproved. The exclusion of (3+1) schemes depends somehow on the used LSND data sample\\footnote{We want to stress that, although (3+1) models themselves are not ruled out, they do not offer an acceptable framework for a combined description of current oscillation data, if LSND is included. In contrast, any model of the (2+2)-type is not viable if the gap separating the two pairs of neutrino states is large compared to $\\dma$.}. Furthermore, it heavily relies on the results of negative SBL experiments, especially on the Bugey and CDHS disappearance experiments.  Therefore, if LSND should be confirmed by MiniBooNE, it will be crucial to improve the experimental data on SBL $\\pnu{e}$ and/or $\\pnu{\\mu}$ disappearance. If the signal in LSND should indeed stem from oscillations in a (3+1) mass scheme the SBL $\\pnu{e}$ and $\\pnu{\\mu}$ disappearance amplitude must be just on the edge of the sensitivity of current experiments.  In summary, the interpretation of the global neutrino oscillation data including LSND in terms of four-neutrino mass schemes -- in either (3+1) or (2+2) realizations -- is strongly disfavoured.  In the best case we obtain a parameter \\gof\\ of only 0.56\\%. The analysis we have presented brings the LSND anomaly to a theoretically more puzzling status indeed. We want to note that also introducing more sterile neutrinos participating in the oscillations is unlikely to substantially improve the situation \\cite{peres}.  If LSND should be confirmed by the results of the MiniBooNE experiment the situation will become even more puzzling. Examples of more far-fetched alternative explanations are the possibility of lepton number violating muon decay~\\cite{Babu:2002ic} or large CPT violation in the neutrino sector~\\cite{CPT}. Such scenarios will be crucially tested at the upcoming experiments MiniBooNE~\\cite{MiniBooNE} and KamLAND~\\cite{kamland}.  \\begin{appendix}"
    },
    "0207/astro-ph0207132_arXiv.txt": {
        "abstract": "We calculate the screening corrections to the electron capture rates in dense stars by the relativistically degenerate electron liquid.  In order to calculate the screening corrections we adopt the linear response theory which is widely used in the field of solid state physics and liquid metal physics.  In particular, we use the longitudinal dielectric function for the relativistically degenerate electron liquid derived by Jancovici. We calculate the screening potential at the position of the nucleus. By using this screening potential one can calculate the screening corrections to the electron capture rates.  We will present accurate analytic fitting formulae which summarize our numerical results.  These fitting formulae will facilitate the application of the present results. The screening corrections to the electron capture rates are typically a few percent. ",
        "introduction": "Since the pioneering works of Fuller, Fowler, \\& Newman (1980, 1982a, 1982b, 1985), calculations of stellar weak-interaction rates have entered an era of precision science.  More recently an authoritative work of Langanke \\& Martinez-Pinedo (2000) on this subject appeared.  Since the presupernova stellar evolution and stellar nucleosynthesis critically depend on the details of the stellar weak-interaction rates (e.g., Wanajo et al. 2002), it is extremely important to calculate accurately the screening corrections to the electron capture rates in dense stars by the relativistically degenerate electron liquid.  This problem has been already addressed by many authors (Couch \\& Loumos 1974; Takahashi, El Eid, \\& Hillebrandt 1978; Gutierrez et al. 1996; Luo \\& Peng 1996; Bravo \\& Garcia-Senz 1999).  The plasma effects on the chemical potential of the nucleus and hence on the threshold energy for the electron capture, in particular, have been discussed by Couch \\& Loumos (1974), Gutierrez et al. (1996), as well as by Bravo \\& Garcia-Senz (1999).  In this paper we will address ourselves to the calculation of the effective potential energy felt by the relativistically degenerate electron.  We will use the linear response theory in order to calculate the screening potential caused by the relativistically degenerate electron liquid.  The present paper is organized as follows.  In \\S~2 we will calculate the effective potential energy felt by the electron using the longitudinal dielectric function of the relativistically degenerate electron liquid derived by Jancovici (1962).  We will thereby calculate the screening potential which will be used for the calculation of the screening corrections to the electron capture rates.  In \\S~3 we will summarize the numerical results in the form of analytic fitting formulae which will facilitate the application of the present results. We will give concluding remarks in \\S~4. ",
        "conclusions": "We have studied the screening corrections to the electron capture rates by the relativistically degenerate electron liquid.  In particular, we have calculated the screening potential caused by the relativistically degenerate electron by using Jancovici's (1962) static longitudinal dielectric function.  We have found that the screening potential is typically a few percent of the electron Fermi energy.  Hence we conclude that the screening corrections to the electron capture rates at high densities are not as great as anticipated by Takahashi, El Eid, \\& Hillebrandt (1978) and also by Fuller, Fowler, \\& Newman (1980).  We have presented accurate analytic fitting formulae which will be useful when one wishes to apply the present results to the calculations of the screening corrections to the electron capture rates at high densities."
    },
    "0207/astro-ph0207304_arXiv.txt": {
        "abstract": "Radio morphology data have been collected for a sample of radio galaxies from the Revised 3rd Cambridge (3CR) Catalog in the redshift range $0.15 < z < 0.65$.  Radio structure parameters including largest physical size, projected bending angle ($\\beta$),  lobe length asymmetry ($Q$) and hot spot placement (Fanaroff-Riley ratio) have been measured from the highest quality  radio maps available.   Combined with similar data for quasars in the same redshift range, these morphology data are used in conjunction with a quantification  of the richness of the cluster environment around these objects (the amplitude of the galaxy-galaxy spatial covariance function, $B_{gg}$) to search for indirect evidence of a dense intracluster medium (ICM). This is accomplished by searching for confinement and distortions of the radio structure that are correlated with $B_{gg}$.  Correlations between physical size and hot spot placement with $B_{gg}$ show evidence for an ICM  only at $z \\leq 0.4$,  but there are no correlations  at $z \\geq 0.4$, suggesting an epoch of $z \\sim 0.4$ for the formation of the ICM in these Abell richness class 0-1, FR2-selected clusters.  X-ray selected clusters at comparable redshifts, which contain FR1 type sources exclusively, are demonstrably richer than the FR2-selected clusters found in this study. The majority of the radio sources with high $B_{gg}$ values at $z\\leq$0.4 can be described as ``fat doubles'' or intermediate FR2/FR1s. The lack of correlation between $B_{gg}$ and $\\beta$  or $B_{gg}$ and $Q$ suggests that these types of radio source distortion are caused by something other than interaction with a dense ICM.  Therefore, a large $\\beta$ cannot be used as an unambiguous indicator of a rich cluster  around powerful radio sources. These results support the hypothesis made in Paper 1 that cluster quasars fade to become FR2s,  then FR1s, on a timescale of 0.9 Gyrs (for H$_0=$ 50 km s$^{-1}$ Mpc$^{-1}$). ",
        "introduction": "Just as the solar wind was first detected indirectly using comet tails, the intracluster medium (ICM) was first inferred to be present due to the morphology of extended radio sources associated with clusters. Fiducial studies of these cluster radio sources date from the 1970s with the work of \\cite{Owen}, \\cite{OwenRud} and \\cite{BOR79}. While some details concerning the exact relationship between the radio  galaxy and the ICM which surrounds it are still the subject of some debate \\citep{Roettiger, Eilek}, all current models for cluster radio source distortion invoke a dense ICM, in an amount consistent with its direct detection via thermal X-ray bremsstral\\\"ung \\citep{Sarazin}.  At low-$z$ ($\\leq$0.1) there is a good correlation (but with a large dispersion)  between the density of the ICM (as measured by its X-ray emission) and the cluster galaxy density \\citep[e.g.,][]{AK83, YE02}. However, since galaxies formed before the ICM, and at least partially created it through internal galaxy processes (i.e., supernovae and stellar winds), a cluster with a high galaxy density may not yet have formed a dense ICM at its observed epoch.  While sensitive X-ray observations have discovered and studied the ICM of  rich clusters out to $z\\approx$1 \\citep[e.g.,][]{GL94,Donahue,Rosati},  direct X-ray detection of a cluster ICM around radio galaxies and quasars is made much more difficult because the AGN itself is a strong X-ray emitter. Because quasars and powerful radio galaxies are found in clusters only  at higher redshifts (Hill \\& Lilly 1991, HL hereafter; Ellingson, Yee \\& Green 1991, EYG hereafter; Harvanek et al. 2001, Paper 1 hereafter), the presence or absence of a dense ICM surrounding luminous AGN in clusters has not been generally established through direct ICM detection (although a few recent detections have been made; see below). Thus, the use of the radio source morphology to trace indirectly  the presence (or absence) of a cluster ICM is still an important indirect method.  In general, radio sources associated with AGN can be divided into two basic types based upon morphology and radio power level: Fanaroff-Riley Type 1 and Type 2 \\citep[FR1 and FR2;][]{FR}. The distorted morphology of the FR1 type radio galaxies indirectly reveals the presence of a dense ICM in current epoch clusters. On the other hand, quasars, which are exclusively FR2 type sources, and powerful FR2 radio galaxies are found in clusters only at $z\\geq$0.2 (HL and Paper 1).  The FR classes were originally defined using the distance between the brightest flux points on opposite sides of the radio core divided by the total extent of the source measured from the faintest radio contour. If this ratio of distances (called herein the $FR$ ratio) is $<$  0.5, the source is classified as an FR1 type. If the $FR$ ratio $>$  0.5, the source is an FR2. Qualitatively, FR2s have a relatively weak radio core with an extended ``lobe'' of emission on each side.  The lobes tend to be fairly symmetric in both size and luminosity and are usually collinear (i.e., both lobes and the core lie on approximately the same line; although see Stocke, Burns \\& Christiansen 1985 for examples of ``bent'' FR2s).   The brightest regions of these sources tend to occur at or near the leading edges of the radio lobes and often these sources contain a weak, one-sided jet.  In contrast, FR1s are asymmetric and distorted and may bear little resemblance to a double-lobed structure.  These sources tend to have bright cores and/or bright, two-sided jets and the extended structure of the source dims with distance from the core. FR2s tend to have a larger physical size than FR1s. Thus, an FR1 type structure is small, distorted, asymmetric, edge-dimmed and (core+jet)-dominated.  FR2 type structures tend to be larger, collinear, symmetric, edge-brightened and lobe-dominated. Additionally, sources with 178 MHz radio power $P_{178} \\lesssim 2 \\times 10^{25}$ W Hz$^{-1}$ sr$^{-1}$ are FR1s while those with higher power are FR2s \\citep{FR}. More recent work on the FR1/FR2 dividing line has found a correlation between the radio power level of the dividing line and the optical host galaxy luminosity, such that more luminous optical galaxies can host more luminous FR1 type sources \\citep{OwenLaing}.  While the exact relationship between FR1s and FR2s remains uncertain, the work of EYG and Paper 1 has found evidence that cluster quasars evolve into radio galaxies by fading to become first FR2s, then FR1s. This ``fading AGN'' or ``evolutionary'' hypothesis  accounts for the presence of  quasars in moderately rich clusters at $z\\sim$0.5 and FR2s in similar richness clusters at $z\\sim$0.25; whereas only FR1s are found in such clusters in the current epoch. The ``e-fading'' timescale of the optical continuum emission from the AGN core suggested by EYG and Paper 1 is $\\sim$0.9 Gyrs (H$_0$=50 km s$^{-1}$ Mpc$^{-1}$).  Since the timescales associated with extended radio source outbursts are thought to be $\\sim$10$^8$ yrs \\citep{revmodphys}, the radio source power and morphology will ``track'' the fading of the AGN. By this hypothesis, the ICM surrounding the AGN host plays a significant role in this process (EYG and Paper 1) by affecting the triggering and/or fueling of the AGN \\citep{Roos, StocPerr}. This scenario also explains why quasars are found in clusters only at $z\\geq$0.4 (EYG) and only in poorer environments at lower-$z$. However, this hypothesis remains controversial, especially since another hypothesis \\citep{Barthel} relates quasars and FR2 radio galaxies entirely by orientation.  Indeed, there is much support for an ``obscuration based unification scheme'', as advocated by Barthel and others \\citep[e.g.,][]{Antonucci} in which quasars and radio galaxies are related not by evolution but by viewing angle; i.e., quasars are seen closer to their outflow axis than radio galaxies, causing radio galaxies to be preferentially lower in optical luminosity due to obscuring material perpendicular to the outflow axis. Evidence cited in favor of this hypothesis includes: (1). quasars have systematically smaller (factor of two at a $\\sim90\\%$ confidence level) extended radio structures on the plane of the sky than radio galaxies at the same redshift and radio power levels \\citep{Barthel}; (2). most luminous radio galaxies have only narrow emission lines in their optical spectra while all quasars have broad permitted lines \\citep{Antonucci}; (3). some narrow-line radio galaxies have broad permitted lines in polarized light \\citep{Antonucci}, including Cygnus A \\citep{ogle} (4). the luminosity of extended [OII] emission is comparable for radio galaxies and quasars with comparable radio power levels \\citep{Hes}; and (5). the far-infrared dust emission has comparable luminosities in quasars and radio galaxies at $z>0.8$ \\citep{meisenheimer}. But none of these results are without contradictory (or at least confusing) results from other investigators: (1). many studies on the extended radio size of quasars and radio galaxies have been conducted \\citep[see][and references therein]{Urry} with little agreement. It may be that the redshift range used by \\citep{Barthel} ($0.5<z<1.0$) is the only redshift range for which this test gives results in agreement with the unified scheme \\citep{Singal}; (2 \\& 3). some narrow line radio galaxies at lower redshift have only very weak emission lines \\citep[e.g., the ``low-excitation'' group of][]{laing94} and would be very unlikely to be ``hidden quasars''. The absence of luminous emission line gas in many FR2s could be an evolutionary effect since FR1s also have no broad line gas, very little narrow line gas \\citep{owen96} and no evidence for obscuration \\citep{falcke}; (4). the extended [OIII] emission is not comparable for radio galaxies and quasars \\citep{Jackson}. Advocates of unified schemes interpret this to mean that some of the [OIII] emission is obscured in the radio galaxies but not in the quasars, but this was originally taken to provide evidence against the unified hypothesis. Alternately, the [OII] emission, being lower ionization state gas, could be ionized by other sources than the AGN (e.g., star formation); (5). other infrared studies of quasars and radio galaxies in other redshift ranges \\citep{Hes2} find more infrared emission from quasars than radio galaxies. Thus, like the extended narrow line gas and the radio source size tests, this test seems to give results that are dependent upon redshift and thus inconclusive support for the unified schemes. So, while there is some significant support for a unified scheme, there seems to be significant contrary evidence as well. And because several of these tests yield different results in different redshift ranges, an evolutionary scenario is suggested.  In this work, we seek to test the evolutionary scenario for quasar fading by searching for indirect evidence of an ICM forming around some of the radio galaxies and quasars studied in Paper 1 using distortions in their extended radio structure. If the evolutionary hypothesis is correct, distorted radio structure  will be found only for the lower $z$ cluster sources in our sample. We use this indirect methodology because very few high-$z$ AGN-selected clusters have been detected directly by extended X-ray emission. Unambiguous detections exist for only five sources in our sample \\citep[e.g.,][]{GizLea,HW99, WB00},  which are insufficient to draw firm conclusions about the intermediate redshift quasar and FR2 radio galaxy population as a whole.  Several previous studies have been conducted on the relationship between environment and FR type at low and intermediate redshift. In an attempt to determine the relationship between FR1 and FR2 type radio galaxies, \\cite{OwenLaing} and \\cite{OwenWhite} examined very low redshift ($z < 0.2$) radio galaxies both in clusters \\citep{OwenWhite} and out of clusters \\citep{OwenLaing}.  They further subdivided the types of radio structure and found evidence for a third, transitional (FR1/FR2) type, which  they refer to as ``fat doubles''. The lower power FR1s and the transitional FR1/FR2s were found mostly in clusters while the higher power FR2s were found to avoid clusters. Similar results concerning the effects of environment on radio structure were found earlier by \\citet{Stoc78}, \\citet{LS79}, \\citet{Lilly87}, and \\citet{Prestage}.   In work used by us in Paper 1, HL extended these  studies of the environments of FR1 and FR2 radio galaxies to $z \\sim 0.5$.  Their results showed that higher power FR2 radio galaxies are found in richer galaxy environments at higher $z$ than at lower $z$, while the lower power FR1s show no change in environment with redshift. \\citet{chap5} confirm that FR1s  are found in rich clusters with a dense ICM regardless of redshift by observing with the VLA a sample of X-ray selected  clusters out to $z\\sim$0.8. While FR1s were found in distant clusters where a dense ICM is demonstrably present, Rector, Stocke \\& Ellingson (1995) found evidence against the  presence of a dense  ICM around a sample of quasars (FR2s) at these same redshifts. Because the \\citet{Rector} work  used the same methodology as employed herein, this paper will incorporate and expand upon  that work to include both quasars and radio galaxies.   To investigate the evolution of cluster AGN further, in Paper 1 of this series we conducted an optical imaging study for a large sample of 3CR/FR2 radio galaxies at 0.15$\\leq z \\leq$0.65, quantifying cluster richness for a very large percentage of the sample using the B$_{gg}$ formalism (see e.g., EYG). While most FR2s imaged in Paper 1 have B$_{gg}<$500Mpc$^{1.77}$ (Abell richness class $<$0), a comparable percentage ($\\sim$25\\%) of all FR2s studied at $z$=0.15-0.65 are in richer environments (B$_{gg}$= 500-1200 Mpc$^{1.77}$ or Abell class 0-1). From that study we concluded that the cluster richness data offer significant evidence against the \\citet{Barthel}  hypothesis that quasars are ``beamed'' or less-obscured radio galaxies, since many FR2 radio galaxies were found in clusters at epochs 0.15$\\leq z \\leq$0.45 when no quasars are found in clusters. Further, the cluster FR2s in this redshift range provide the ``missing link''  between cluster quasars at $z\\sim$0.5 and cluster FR1s in the current epoch.  In Paper 1 we interpreted the result that  FR2s are found in clusters in the past but not in the present epoch  as evidence for an evolving ICM in these AGN-selected clusters. The fading of both the AGN optical brightness and the radio power (and its associated morphological changes) are suggestive of evolution between the types of radio galaxies themselves. That is,  Paper 1 suggests that quasars at $z\\sim$0.5 fade to become FR2 radio galaxies at $z\\sim$0.25 and then FR1s at $z\\sim$0 on an ``e-fading'' timescale of 0.9 Gyrs (for H$_0$= 50 km s$^{-1}$ Mpc$^{-1}$).  In this paper we use the 3CR and quasar samples from Paper 1  and \\citet{Rector} to investigate the extended radio morphologies of these sources for confirmation of this ``evolutionary hypothesis''. That is, do we see changes in the radio morphology of the 3CR/FR2 sample used in Paper 1 due to the changing ICM density expected in the ``fading AGN'' hypothesis?  This paper is organized as follows: The AGN samples used in this study are discussed in Section 2. The radio structure parameters used herein are defined and the morphology data are presented in Section 3. These parameters are simple, quantitative measures of radio source size (and thus confinement)  and distortions. The results of the correlation analyses between  the radio structure parameters and the density of the surrounding galaxy environment ($B_{gg}$) are given in Section 4.  This correlation analysis shows that only some of the radio structure parameters are sensitive to the presence of a dense ICM. However, those parameters that do correlate with B$_{gg}$ do so only at $z\\leq$0.4, strongly suggesting that the ICM around these sources is just forming at $z\\approx$0.4. Section 5 contains a  comparison between the clusters found around the FR2s in the present sample  and rich, X-ray selected clusters at comparable redshifts, which contain only FR1 type radio sources \\citep{chap5}. Because the X-ray selected clusters are demonstrably richer than the AGN-selected clusters in our sample, it is not surprising that both a dense ICM and also  FR1 radio galaxies are present in them. A summary of the findings is given in Section 6, in which we report that this investigation provides substantial, additional support for the ``evolutionary'' hypothesis of EYG and Paper 1.  Future work that can further test this hypothesis   also is proposed.  Values of H$_0$ = 50 km s$^{-1}$ Mpc$^{-1}$ and q$_0$ = 0 are assumed throughout this work to be consistent with earlier work;  however, none of the results are sensitive to these choices. ",
        "conclusions": "\\subsection{A Summary of the Findings}  In Paper 1 of this series, evidence was presented that quasars and FR2 radio galaxies are primarily related, not by orientation as previously proposed by \\cite{Barthel}, but by evolution (EYG); i.e., quasars in clusters at $z\\sim$0.5 fade to become FR2 radio galaxies at $z\\sim$0.25, and continue to fade to become FR1 radio galaxies at $z\\sim$0. In order to further test the ``evolutionary hypothesis'', we have scrutinized the radio morphologies for a large sample of quasars and radio galaxies at $z$=0.15-0.65 to search for indirect evidence of a dense ICM around some of these sources.  This was accomplished using various morphology parameters that measure the confinement and distortion of the extended radio source. We then determined whether these morphology parameters were correlated with a quantitative measurement of galaxy richness ($B_{gg}$) from Paper 1. These morphology  parameters included: projected  physical size, hot spot location (Fanaroff-Riley or $FR$ ratio), projected bending angle ($\\beta$) and lobe length asymmetry ($Q$). Using the radio structure measurements and $B_{gg}$ values for a combined sample of 63 radio galaxies and 35 quasars in the above redshift range we find the following results:   1.  When the radio structure data for the radio galaxies and quasars are divided into low-$z$ ($<$0.4) and high-$z$ ($>$0.4) subsamples, correlations between environment and projected physical size and between environment and $FR$ ratio show evidence for the presence of a dense ICM in the richer environments ($B_{gg} > 500$ Mpc$^{1.77}$) only at $z \\lesssim 0.4$.  There are no correlations and, thus no evidence for an ICM around the sources at $z \\gtrsim 0.4$. This suggests that the formation of a dense ICM in environments of $B_{gg} \\sim 500-1000$ Mpc$^{1.77}$ occurs at $z \\sim 0.4$.  2.  There is no correlation between projected bending angle and environment or between lobe length asymmetry and environment at any redshift within the sample.  This indicates that bending angle distortions and asymmetry distortions in the radio structure are not caused by an interaction with the ICM. Instead, a collision between the radio jet/lobe and a nearby galaxy or dense intergalactic cloud may be responsible \\citep[e.g.,][] {SBC}.  3. The lack of a correlation between bending angle and environment also demonstrates that for FR2 type radio sources, a large bending angle is not a reliable predictor of a rich galaxy environment.  Thus, it should not be used as a indicator of a cluster of galaxies around an FR2 at high-$z$, as has been previously proposed by \\citet{Hintzen} and recently reproposed by \\citet{Blanton}.   4.  No significant correlations between radio structural properties and radio power are found within the range of radio powers used in this study: 26.8 $\\le$ log P$_{178}$ (W Hz$^{-1}$) $\\le$ 29.0.   Specifically,  all sources studied are FR2s, significantly above the FR2/FR1 dividing line in radio power. Thus, the differences in radio structure seen in our sample are not due to variations in radio power  but rather due to interaction with a dense ICM.  5. At $z <$0.4, the sources found in clusters with $B_{gg}\\geq$500Mpc$^{1.77}$ almost exclusively have morphologies that can be described as ``fat doubles'' \\citep{OwenLaing}, with brightest spots within the extended lobes which are well back from the leading edges of the lobes, and often in luminous jets. Three of these sources (3C\\,28, 3C\\,346 \\& 3C\\,348) have detected extended X-ray emission surrounding them \\citep[]{HW99,GizLea,WB00} so that for these cases, there is both direct and indirect evidence for a dense ICM. Additional observations of AGN clusters with CHANDRA will be able to test directly the inferences made here (see Section 6.2).  6.  Correlations between redshift and $FR$ ratio and between redshift and projected physical size support an ICM formation at $z \\sim 0.4$ in environments of $B_{gg} \\sim$ 500-1000 Mpc$^{1.77}$.  7. Numerous X-ray selected clusters are known to exist in the redshift range  of our sample. These clusters have radio galaxy populations which are exclusively FR1 in morphology and power level \\citep[e.g.,][]{chap5}. However, these clusters are substantially richer than the environments of the sources in our radio galaxy and quasar samples at the same redshift, indicating that the formation epoch of a dense ICM depends, as expected theoretically, on cluster richness.   In fact, the poorest of the EMSS cluster environments are comparable to the richest galaxy environments in the radio galaxy and quasar samples.    Thus, the results of this study are consistent with the predictions of the ``evolutionary hypothesis'' for radio-loud AGN first suggested by EYG. Specifically, in EYG and in Paper 1, evidence was presented that quasars found in clusters at $z\\sim$0.5 fade quickly (``e-fading'' timescale of 0.9 Gyrs for H$_0$=50 km s$^{-1}$ Mpc$^{-1}$) to become, first FR2 radio galaxies at $z\\sim$0.25, and then FR1 radio galaxies at the current epoch. While the physical mechanism for this fading is unknown, in Paper 1 we speculated that if radio-loud AGN are powered by a rapidly spinning Black Hole \\citep[e.g.,][]{BZ77,BBR80,WC95}, the development of a deep gravitational potential and a dense ICM around these AGN would prevent further  ``spin up'' by preventing the formation of new, supermassive Black Hole binaries. These  binaries would not form because any galaxy-galaxy collisions in this cluster would be both at much higher relative velocities than previously and also would be relatively gas free, so that bound binary formation would be much less likely. In the absence of additional mechanisms to spin up the supermassive Black Hole, each radio outburst would extract spin energy that could not be replaced and so each consecutive outburst would be less powerful. Since the suggested timescale for radio outbursts ($\\sim$10$^8$ years; Begelman et~al.\\ 1984) is short compared to the AGN fading timescale measured in Paper 1, the radio source power and structure would ``track'' the fading. So, while this scenario is consistent with a duty cycle of $\\sim10\\%$, our data do not address directly the question of AGN duty cycle. And, while the spin down of a supermassive Black Hole seems to us to be the most attractive scenario to account for the fading of cluster AGN, it is also possible to imagine that an accretion powered AGN could have its fueling stifled by the development of a dense ICM \\citep{StocPerr}.  \\subsection{Prediction for Future Observations}  The important aspects of our conclusions which are amenable to test currently include the presence or absence of a dense ICM based upon the FR2 radio source morphology and the possibility that FR1s in clusters were FR2s in the past. In both cases these tests involve CHANDRA imaging spectroscopy.  The first prediction is in two parts. First, we predict that a dense X-ray emitting ICM will  be found around the ``fat doubles'' at $z\\leq$0.4, some of which are shown in Figure 6. In the cases of 3C\\,348, 3C\\,346 and 3C\\,28 there is already considerable evidence for extended X-ray emission around these three AGN \\citep[e.g.,][] {GizLea}. Two other AGN with small $FR$ ratios at slightly higher $z$ (the quasar 3C\\,215 with $FR$=0.65 at $z$=0.411 and the radio galaxy 3C\\,295 with $FR$=0.78 at $z$=0.461) also have X-ray  evidence from ROSAT for a dense ICM as their radio morphologies and $B_{gg}$ values (1000 and 1030 Mpc$^{1.77}$ respectively; see Paper 1) predict. In fact, \\citet{Harris} have already detected  extended cluster X-ray emission around 3C\\,295 with CHANDRA.  The other ``fat doubles'' should also be imaged with CHANDRA to make sure that the presence of an ICM is generic to this class of sources.  While necessary, this first prediction is  insufficient to test completely our use of radio morphology to locate a dense ICM. It is also important to verify that there is no ICM present around quasars and FR2s of ``classical double'' morphology (i.e., high $FR$ ratios). In our survey this type of source is found both in low-$B_{gg}$ regions at all redshifts and in high-$B_{gg}$ regions at $z>$0.4. Of particular importance are the quasars with high-$B_{gg}$ at high-$z$, which are candidates for clusters whose ICM has yet to form. Examples of this class include the quasars 3C\\,263 ($z$=0.646; $B_{gg}$=993 Mpc$^{1.77}$), 3C\\,275.1 ($z$=0.557; $B_{gg}$=1125 Mpc$^{1.77}$)  and PKS 0155-109 ($z$=0.616; $B_{gg}$=777 Mpc$^{1.77}$). 3C\\,275.1 is particularly important to observe with CHANDRA both because it has a large $\\beta$ \\citep[see map in][]{SBC}, and because there is a tentative ICM detection made with the ROSAT HRI. Our analysis in this paper suggests that, despite the large $\\beta$ and the large $B_{gg}$ value (Abell richness class 1), CHANDRA observations will fail to confirm the tentative HRI detection of a dense cluster ICM around 3C\\,275.1. Since one definite CHANDRA detection already has been made (Worrall et~al.\\ 2001) of a dense ICM around a ``classical double'' FR2 at high-$z$ (3C\\,220.1, $z$=0.620, $B_{gg}$=418 Mpc$^{1.77}$), a few others would call into question the methodology used herein to infer the presence or absence of a dense ICM from the radio source morphology.  Perhaps the most controversial hypothesis put forward in this paper and Paper 1 is the idea that FR2s in clusters at high-$z$, fade and become FR1s in current epoch clusters. However, the recent  CHANDRA discovery of ``holes'' in the ICM X-ray emission in some clusters  may offer a means of testing this hypothesis by discovering ``fossil'' evidence of FR2s around FR1s.  In at least one case of an X-ray ``hole'' around the FR1 radio galaxy in Abell 4059 \\citep{Heinz}, the inferred power required to evacuate  the ``hole'' of X-ray emitting gas by $pdV$ work  is much larger than the power inferred to be present in the FR1 ($\\sim$5$\\times$10$^{42}$ ergs s$^{-1}$) averaged over its 10$^8$ yr lifetime (Reynolds, Heinz \\& Begelman 2001, which uses the prescriptions in Bicknell, Dopita \\& O'Dea 1997). The X-ray cavity walls also do not correspond with the current boundaries of the radio source lobes. Both the anomalously large size and power requirements of the X-ray cavity  in Abell 4059 suggest that the radio source was both larger and more powerful in the recent past; i.e., the previous outburst was an FR2. Also, on the basis of the interpretation put forward in this paper, the ICM of Abell 4059 could have been much less dense in the recent past than now, also making the cavity easier to create.  On the basis of the current work,  we predict that this one case is not unique, but that other, similar examples will be found with CHANDRA. Indeed, the ``fat doubles'' in our sample are ideal targets to search for such evidence since, from the ``evolutionary hypothesis'',  the most recent outburst of these sources is the first after a dense ICM has formed around them."
    },
    "0207/astro-ph0207074_arXiv.txt": {
        "abstract": "Laser guide stars created by Rayleigh scattering provide a reasonable means to monitor atmospheric wavefront distortions for real-time correction by adaptive optics systems.  Because of the $\\lambda^{-4}$ wavelength dependence of Rayleigh scattering, short wavelength lasers are a logical first choice for astronomical laser guide star systems, and in this paper we describe the results from a sustained experimental effort to integrate into an adaptive optics system a 351 nm Rayleigh laser guide star created at an altitude of 20 km (above MSL) at the Mt. Wilson 2.5-m telescope. In addition to providing obvious scientific benefits, the 351 nm laser guide star projected by UnISIS is \"Stealth qualified\" in terms of the FAA and airplane avoidance.   Due to the excellent return signal at the wavefront sensor, there is no doubt that future applications will be found for short-wavelength Rayleigh scattered laser guide stars. ",
        "introduction": "UnISIS (University of Illinois Seeing Improvement System) is a laser guided adaptive optics system operating at the Coude focus of the 2.5-m telescope at Mt. Wilson Observatory.  It is the first astronomical system to employ a Rayleigh laser guide star at 351 nm.  While several descriptions of UnISIS have been published during system development and construction \\citep{tho94,tho95,tho98}, this paper provides a complete over-view of the laser guide star system with detailed information on its design and its \"as-built\" configuration. The lessons learned in the UnISIS development effort will be of interest to those who are planning current generation laser guide star systems.  To keep this paper to a reasonable length but still describe key experimental issues in depth, only the UnISIS laser guide star system is discussed here.   Subsequent papers will describe the UnISIS adaptive optics system and its closed-loop performance characteristics including the cone effect (i.e., focal anisoplanatism) caused by the laser guide star's location in the near field when compared to astronomical objects.  Soon after the publication of the seminal paper describing the laser guide star concept by \\citet{foy85}, \\citet{tho87} reported on experimental work with laser guide stars and began detailed engineering design studies for laser guided adaptive optics systems (\\citet{gar90} and references therein).  A key outcome of this early work was the realization that sodium laser guide stars, while conceptually attractive, would remain difficult to implement in the near-term especially if the goal is to successfully operate an adaptive optics system for scientific observations at wavelengths less than 2.2 microns.  Sodium wavelength lasers with sufficient power to adequately excite the sodium resonance line at 589.3 nm were not available at that time, and fifteen years later, these lasers are still difficult to obtain. To date, the only operational sodium laser guide star system is that at Lick Observatory \\citep{max97} where a 15 W sodium laser is used on a regular basis. The Lick Observatory laser -- and its ``sister'' laser at the Keck Observatory -- are one-of-a-kind systems built by Lawrence Livermore National Laboratory that will probably never be duplicated again.  Rayleigh laser guide star techniques at 351 nm were discussed from the earliest times (\\citet{tho89}; \\citet{san94}) because commercial-quality excimer lasers capable of producing significant levels of pulsed power at 351 nm were already in production at that time.  This led to the experimental work of \\citet{tho92} who reported on the creation and initial calibration of the return flux from a 351 nm laser guide star to altitudes up to $\\sim$33 km.  After the Thompson and Castle experimental work had begun, the U.S. Air Force declassified information on the laser guided adaptive optics system at Starfire Optical Range \\citep{fug91}.  The Starfire group had independently chosen to develop a Rayleigh laser guide star system that used the backscattered light from a copper-vapor laser (at 511 nm and 578 nm).  The Starfire system design \\citep{fug94} -- even though independently devised -- closely matched the published design work by \\citet{gar90}.  Subsequent visits to Starfire Optical Range provided information that improved the conceptual design of UnISIS. Every laser guided adaptive optics system employs a unique set of tools to multiplex the optical system, to project the laser beacon into the sky, to reject scattered light, etc. Our purpose here is not to present a comprehensive review of these methods. Those interested in examples of other systems can refer to \\citet{gre92}, \\citet{fug94}, and \\citet{san94}, and references therein.  Determining the ideal laser for Rayleigh scattered laser guide stars is the subject of another paper \\citep{tho02a}, but it is worth noting here that short wavelength Rayleigh laser guide star systems are a viable alternative for astronomical purposes and that market forces continue to drive technological improvements for short wavelength commercial lasers.  UnISIS relies in a 30-Watt excimer laser originally built by Questek Incorporated.  This laser is no longer in production, its basic structure having morphed into the laser systems now used by VISX for LASIK eye surgery.  About two years ago Lambda Physik placed in production an excimer laser system more powerful ($\\sim$150 Watts) and better suited for astronomical use called Lambda Steel (developed for the manufacture of flat panel displays).  It is also worth noting that diode pumped and frequency tripled Nd:YAG and YLF lasers, which operate at 355nm and 349 nm, respectively, are also available commercially, and both are viable systems for short wavelength Rayleigh laser guide star systems.  This paper contains a diverse collection of both design information and practical experience obtained during the development of UnISIS.  It will provide assistance to those who are beginning the process of designing and implementing a laser guide star system. The field of laser guided adaptive optics is still in an early phase of development, and there are many new ideas to discover and to exploit.  In this paper, the most notable achievements are the \"Stealth\" characteristic of the UnISIS laser guide star system (Sec. 5) and the successful acquisition of the laser wavefront return signal from the Rayleigh laser guide star (Sec. 8). ",
        "conclusions": "The UnISIS Rayleigh laser guide star system has been commissioned and operated in open loop with satisfactory return signal for a 13 x 13 set of subapertures across the pupil of the Mt. Wilson 2.5-m telescope.  A tightly focused laser guide star return signal with a FWHM $\\sim$1 arcsec was received from a $\\Delta$z = 2.5 km range gate centered at 18.2 km above the telescope ($\\sim$20 km above mean sea level).  This successful demonstration of the UnISIS laser guide star system sets the stage for closed loop operation with the full UnISIS adaptive optics system.  The 351 nm Rayleigh laser guide star technique with full- aperture broadcast provides a reasonable method to monitor wavefront perturbations in the Earth's atmosphere, and it is especially attractive because of its \"Stealth\" characteristics.  Rayleigh laser guide stars are likely to be the basis for many other laser guided adaptive optics applications in the future."
    },
    "0207/astro-ph0207597_arXiv.txt": {
        "abstract": "We have studied the host galaxies of a sample of radio-loud AGN spanning more than four decades in the energy output of the nucleus. The core sample includes 40 low-power sources (BL Lac objects) and 22 high-power sources (radio-loud quasars) spanning the redshift range $0.15\\lesssim z\\lesssim 0.5$, all imaged with the high spatial resolution of HST. All of the sources are found to lie in luminous elliptical galaxies, which follow the Kormendy relation for normal ellipticals. A very shallow trend is detected between nuclear brightness (corrected for beaming) and host galaxy luminosity. Black hole masses are estimated for the entire sample, using both the bulge luminosity--black hole mass and the velocity dispersion--black hole mass relations for local galaxies. The latter involves a new method, using the host galaxy morphological parameters, $\\mu_e$ and $r_e$, to infer the velocity dispersion, $\\sigma$, via the fundamental plane correlation. Both methods indicate that the entire sample of radio-loud AGN are powered by very massive central black holes, with $M_{\\bullet}\\sim 10^8$ to $10^{10} M_\\odot$. Eddington ratios range from $L/L_{Edd} \\sim 2\\times10^{-4}$ to $\\sim 1$, with the high-power sources having higher Eddington ratios than the low-power sources. Overall, radio-loud AGN appear to span a very large range in accretion efficiency, which is all but independent of the mass of the host galaxy. ",
        "introduction": "Whether there is a link between the intrinsic power of Active Galactic Nuclei (AGN) and their host galaxies is not known. It seems plausible that more massive host galaxies might form in high density regions that also support the formation of more massive nuclear black holes \\citep{Small,Haehnelt,Kauffmann} and/or that more massive host galaxies could support an increased rate of fuelling. For nearby, non-active galaxies there is observational evidence that the mass of the central supermassive black hole is correlated with bulge mass \\citep{Magorrian,vanderMarel} and with bulge velocity dispersion \\citep{Ferrarese,Gebhardt}. This could lead to an observable link between emission from the region around the black hole and the luminosity of the hosting galaxy, for example as observed by \\citet{vanderMarel99} for local spheroids.  Several studies have indeed suggested nuclear luminosity might be related to host galaxy mass in AGN \\citep{McLeod,Schade,Hooper}; however, other studies find no such relation \\citep{Urry,McLure,Bahcall,Smith2}. Certainly among radio-loud AGN alone, for which the host galaxies are generally luminous ellipticals with well-defined morphologies, no trend has been detected in previous studies. This may be due to the fact that only a small range of (high) nuclear power has been probed in the past.  To make a clean comparison among AGN that differ only in nuclear output (rather than host galaxy morphology, dust content, star formation history, etc.), we restrict the present study to radio-loud AGN ($F_{\\rm 5~GHz}/F_B > 10$; Kellerman {\\it et al.} 1989). These are known to have relativistic jets formed near the central supermassive black hole \\citep{UrryP}, and so should be governed by similar physics near the black hole. Additionally, the early-type spectral energy distributions typical of the host galaxies \\citep{Ridgway, Scarpa2, McLure2, Pentericci} make the K corrections straightforward.  Our goal is to probe the connection between AGN power (processes near the black hole) and environment (host galaxy properties) for radio-loud AGN over the full range of intrinsic nuclear power. Using BL Lac objects and Radio-Loud Quasars (RLQs) to represent the extremes of this range, it is possible to define redshift-matched samples that span more than four orders of magnitude in intrinsic nuclear power. This range reflects a continuum of accretion powers, which in turn arise from some variation in the process of fuelling and/or jet formation near the black hole. ",
        "conclusions": "We find that the host galaxies of radio-loud AGN are luminous ellipticals, occupying the low surface-brightness tail of the Kormendy relation for normal elliptical galaxies, and are statistically consistent with this relation. Comparing the host galaxies of low-power and high-power radio-loud AGN, we find general overlap, with a slight difference in median absolute Cousins R magnitudes, $-23.75$~mag  and $-24.2$~mag, respectively. After correcting the (highly beamed) low-power AGN for Doppler beaming, we find a significant positive trend between nuclear and host galaxy luminosity, but with a very shallow slope --- a factor of 1.3 in host galaxy brightness over at least four orders of magnitude in nuclear luminosity --- ruling out a close relation between host galaxy and nuclear luminosity in radio-loud AGN.  We find that the central black holes of luminous radio-loud AGN are universally large, with median black hole mass $\\sim 10 \\time 10^9 M_\\odot$ for this sample. This is found to be the case using either the $M_{\\bullet}$---$L_{bulge}$ relation and the $M_{\\bullet}$---$\\sigma_e$ relations to derive black hole masses. This supports the view that a high central black hole mass is an important factor in generating a powerful radio source.  No correlation is found between black hole mass and energy output from the nucleus. Rather, the black hole masses derived span a surprisingly small range compared to the range in intrinsic power of this sample. Eddington ratios for radio-loud AGN span more than four orders of magnitude, with $\\frac{L_{bol.}}{L_{Edd}} \\lesssim 2\\times 10^{-4}$ in the lowest-power sources to $\\frac{L_{bol.}}{L_{Edd}} \\sim 1$ in the highest. Across this range, the host galaxies luminosities are tightly constrained, all within one magnitude of brightest cluster galaxies. Thus, although the properties of the host galaxy may have a strong influence the mass of its central black hole, they have at most a very weak influence on the mass accretion rate in radio-loud AGN."
    },
    "0207/astro-ph0207242_arXiv.txt": {
        "abstract": "A torus around a stellar mass Kerr black hole can emit about 10\\% of the spin-energy of the black hole in gravitational radiation, potentially associated with a gamma-ray burst. Wide tori may develop buckling modes by the Papaloizou-Pringle instability and gravitational radiation-reaction forces in the Burke-Thorne approximation. Gravitational wave experiments may discover these emissions in a fraction of nearby supernovae. This provides a test for Kerr black holes, and for GRB inner engines by comparison with the de-redshifted durations of long GRBs. ",
        "introduction": "Stellar mass black holes surrounded by a compact torus may represent catastrophic events such as core-collapse in a massive stars and black hole-neutron star coalescence. These scenarios have been considered as sources of cosmological gamma-ray bursts \\citep{woo93,pac91}. We may consider black hole-torus systems and their emissions more generally, especially when the black hole is rapidly spinning. Their potential association with GRBs provides observational constraints on their evolution.  A torus around a rapidly rotating black hole may develop a state of suspended accretion for the lifetime of rapid spin of the black hole \\citep{mvp01a}. This points towards major energetic output in ``unseen\" emissions, in gravitational radiation, magnetic winds, thermal emissions  and neutrino emissions \\citep{mvp02}. The energy $E_{gw}$ emitted in gravitational radiation is expected to be about $10\\%E_{rot}$, i.e., \\begin{eqnarray} E_{gw} \\simeq 6\\times 10^{53}\\mbox{erg} \\label{EQN_E} \\end{eqnarray} for a $10M_\\odot$ black hole. This output (\\ref{EQN_E}) may be detected by the upcoming Laser Interferometric Gravitational Wave Observatory LIGO \\citep{abr92} and the French-Italian counter part VIRGO \\citep{bra92}, possibly in combination with any of the bar or sphere detectors currently being developed. This provides a calorimetric compactness test for Kerr black holes \\citep{mvp01}, and a means of identifying the inner engine of GRBs by comparison with de-redshifted durations of long GRBs.  In the Woosley-Paczynski-Brown scenario of hypernovae \\citep{woo93,pac98,bro00,bro02}, core-collapse in rotating massive stars forms a Kerr black hole surrounded by a compact disk or torus. Long GRBs correlate with star-forming regions \\citep{blo00} and, hence, young massive stars, possibly in binaries. The rotating black hole may produce wide-angle ejecta back into the interstellar medium leaving behind a soft X-ray transient with a chemically enhanced companion star \\citep{bro00}, such as GRO J1655-40 \\citep{isr99} and V4641Sgr \\citep{oro01}. This potential supernova and SXT association is important in identifying progenitors to GRBs and their inner engines. The beamed output of true GRB energies of $E_\\gamma=10^{50-51}$ergs \\citep{fra01} represents a minor energetic output for a long GRB from a rotating black hole. The potential for long gamma-ray bursts from rotating black holes suggests that GRB associated supernovae may emit bursts of gravitational radiation, e.g., GRB 980425/1998bw \\citep{gal98}, GRB 011121 \\citep{blo02} and GRB 011211 \\citep{ree02}.  In this {\\em Letter}, we suggest LIGO/VIRGO searches for bursts of gravitational radiation from back hole-torus systems using upcoming continuous all-sky supernovae surveys. Focusing on supernovae may serve to reduce data analysis by their well-determined coordinates, which includes distances. The expected gravitational wave-spectrum is here identified with multipole moments in a wide torus due to a Papaloizou-Pringle instability, by extension of the theory for slender tori \\citep{pap84,gol86} and including the secular effect of gravitational radiation backreaction-forces in the Burke-Thorne approximation \\citep{tho69}.  \\begin{figure} \\plotone{f1} \\caption{Shown is the histogram of redshift corrected durations of 27 long bursts with individually determined redshifts from their afterglow emissions (sample from Djorgovski et al., 2001, astro-ph/0107535, and references therein, and updated (Djorgovski, 2002; Hurley, 2002). The mean value of the durations of all bursts is 38s; the average is 23s without the long bursts GRB 980703 $(T_{90}/1+z=400$s) and GRB 000911 $(T_{90}/1+z=243$s), which consist of two (in BATSE; one in Ulysses), respectively, three well-separated sub-bursts. GRBs from rotating black holes are expected to be accompanied by ``unseen\" emissions in gravitational radiation from the torus. This predicts a similar distribution of durations for their bursts of gravitational radiation with an expectation value of about one-half minute.} \\end{figure} ",
        "conclusions": ""
    },
    "0207/hep-ph0207139_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.1in}} Ultra-high energy cosmic neutrinos are incisive probes of both astrophysical sources and new TeV-scale physics.  Such neutrinos would create extensive air showers deep in the atmosphere. The absence of such showers implies upper limits on incoming neutrino fluxes and cross sections. Combining the exposures of AGASA, the largest existing ground array, with the exposure of the Fly's Eye fluorescence detector integrated over all its operating epochs, we derive 95\\% CL bounds that substantially improve existing limits.  We begin with model-independent bounds on astrophysical fluxes, assuming standard model cross sections, and model-independent bounds on new physics cross sections, assuming a conservative cosmogenic flux. We then derive model-dependent constraints on new components of neutrino flux for several assumed power spectra, and we update bounds on the fundamental Planck scale $M_D$ in extra dimension scenarios from black hole production. For large numbers of extra dimensions, we find $M_D > 2.0\\ (1.1)~\\tev$ for $\\mbhmin = M_D\\ (5M_D)$, comparable to or exceeding the most stringent constraints to date. ",
        "introduction": "Cosmic neutrinos provide a unique window on astrophysical processes because they escape from dense regions and typically propagate to the Earth unhindered~\\cite{Sigl:2001th}.  At ultra-high energies, they also provide an important probe of new ideas in particle physics.  In contrast to all other standard model (SM) particles, their known interactions are so weak that new physics may easily alter neutrino properties, sometimes drastically.  This is especially relevant for neutrinos with energies above $10^6~\\gev$, which interact with nucleons with center-of-mass energies above 1 TeV, where the SM is expected to be modified by new physics.  The signal for ultra-high energy neutrinos is quasi-horizontal giant air showers initiated deep in the atmosphere~\\cite{Anchordoqui:2002hs}.  This signal is well-studied and easily differentiated from air showers initiated by hadrons.  The Earth's atmospheric depth rises from about $1000~\\g/\\cm^2$ vertically to nearly $36000~\\g/\\cm^2$ horizontally.  For all but the most extreme (and typically problematic~\\cite{Burdman:1997yg,Kachelriess:2000cb}) neutrino cross sections, the mean free path of neutrinos is larger than even the horizontal atmospheric depth. Neutrinos therefore interact with roughly equal probability at any point in the atmosphere and may initiate showers in the volume of air immediately above the detector. These will appear as ``normal'' showers, with large electromagnetic components, curved fronts (a curvature radius of a few km), and signals well spread over time (of the order of microseconds).  On the other hand, hadrons have interaction lengths of the order of $40~\\g/\\cm^2$ and so always interact at the top of the atmosphere. The electromagnetic component of an air shower has mean interaction length $\\sim 45-60~\\g/\\cm^2$.  For a quasi-horizontal shower initiated by an ordinary hadron, then, this component is absorbed long before reaching the ground, as it has passed through the equivalent of several vertical atmospheres --- 2 at a zenith angle of $60^\\circ$, 3 at $70^\\circ$, and 6 at $80^\\circ$. In these showers, only high energy muons created in the first few generations of particles survive past 2 equivalent vertical atmospheres. The shape of the resulting shower front is therefore very flat (with curvature radius above $100~\\km$), and its time extension is very short (less than $50~\\ns$).  These shower characteristics are exploited by both ground arrays and fluorescence detectors in searches for primary cosmic ray neutrinos. At present, no ultra-high energy neutrino signal has been reported. Here we determine the total exposure for neutrino detection from existing facilities and derive both model-independent and model-dependent bounds on astrophysical neutrino fluxes and new neutrino interactions.  The outline of the paper is as follows. In Sec.~\\ref{sec:exposure} we examine acceptances for neutrino detection and compute the current combined total exposure using all available data from the Akeno Giant Air Shower Array (AGASA)~\\cite{Chiba:1991nf} and Fly's Eye~\\cite{Baltrusaitis:mx} experiments.  In Sec.~\\ref{sec:fluxes} we determine model-independent bounds on the total neutrino flux, assuming SM cross sections. To derive model-independent results, we assume only that fluxes are confined to a small window around some central neutrino energy and obtain bounds as a function of this central energy. After that, we assume a power law neutrino flux $d\\Phi/dE_{\\nu} \\propto E_\\nu^{-\\gamma}$ to obtain stronger, but more model-dependent, bounds on the total neutrino flux from integrating over all energies.  In Sec.~\\ref{sec:interactions} we derive model-independent bounds on high energy neutrino cross sections, assuming a conservative cosmogenic flux.  These significantly improve existing limits~\\cite{Tyler:2001gt}. We then derive model-dependent bounds on cross sections, focusing on the example of TeV-scale gravity scenarios in Sec.~\\ref{sec:gravity}.  We improve existing constraints~\\cite{Anchordoqui:2001cg} on the fundamental Planck scale from the non-observation of microscopic black hole production by cosmic neutrinos~\\cite{Feng:2001ib,Anchordoqui:2001ei,% Emparan:2001kf,Ringwald:2001vk,Kowalski:2002gb}. For large numbers of extra dimensions, these bounds are comparable to or exceed all existing bounds on extra dimensions. Sec.~\\ref{sec:conclusions} contains our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In the first part of this paper, we derived new limits on the cosmic neutrino flux striking the Earth's atmosphere. This was accomplished by searching for quasi-horizontal deeply developing showers in ultra-high energy cosmic ray data, taking into account the combined exposures of the AGASA and Fly's Eye experiments. Our results significantly strengthen existing limits and present serious problems for models where exotic elementary $X$ particles cascade decay to cosmic ray particles. In particular, models where topological defects are responsible for the events detected with energies $\\agt 10^{11}~\\gev$ are severely constrained, because neutrinos are typically a significant component in $X$ decays, and have a hard spectrum extending up to $M_{\\rm GUT} \\sim 10^{16}~\\gev$. The bounds obtained in this paper will also challenge any attempt to normalize the observed spectrum to the proton flux as predicted by top down models.  In the second part of the paper, we used the atmosphere as a giant calorimeter to probe neutrino-nucleon cross sections at $\\sqrt{s} \\agt 1~\\tev$. We first combined the complete neutrino exposure of the above-mentioned facilities with the flux of cosmogenic neutrinos, to derive model-independent upper bounds on the neutrino-nucleon cross section.  These bounds strengthen existing limits by roughly one order of magnitude.  We then considered TeV-scale gravity models to study BH production. The upper bounds on the neutrino-nucleon cross section implied lower limits on the fundamental Planck scale, which represent the best existing limits on TeV-scale gravity for $n \\geq 5$ extra spatial dimensions."
    },
    "0207/astro-ph0207256_arXiv.txt": {
        "abstract": "We present numerical hydrodynamical models of the effects of planets or brown dwarfs orbiting within the extended atmosphere and wind formation zone of Mira variables. We find time-dependent wake dynamics and episodic accretion phenomena which may give rise to observable optical events and affect SiO maser emission. ",
        "introduction": "When stars like the Sun evolve to become AGB stars, their outer envelopes swell to sizes of order an astronomical unit (au). During the AGB phase the star becomes unstable to large amplitude radial pulsations, and becomes a Mira or semi-regular variable (e.g., \\citet{han94}).  These pulsations drive strong shock waves down the steep pressure gradient outside the photosphere, producing an extended atmosphere, and a dense, warm stellar wind (\\citet{bow88}, also see \\citet{wil00}).  In the most luminous stars the wind driving is also enhanced by radiation pressure on dust grains that form a few stellar radii from the photosphere (e.g., \\citet{bow88}, \\citet{gai99}, \\citet{wil00}).  During the Mira phase it is likely that the winds are primarily radial since the massive, extended stellar atmosphere cannot have significant rotation, and at this stage it is unlikely that a globally ordered magnetic field strongly affects the dynamics of the wind (see \\citet{sok02}).  In recent years, about 80 extrasolar planetary systems have been discovered \\citep{but01}.  Generally, these systems have a single gas giant planet, with a mass of up to $\\simeq 15$ Jupiter masses ($M_J$), and an orbital semi-major axis of commonly less than $1.0\\  au,$ but ranging up to $\\simeq 4.0\\  au$ (e.g., \\citet{hat00}, \\cite{fis02}). Increasing numbers of brown dwarf stars have also been discovered recently.  Most of these brown dwarfs are not companions of main sequence Mira progenitors (but see \\cite{fis02}, \\cite{liu02}). The more massive ``planets'' are in fact low mass brown dwarfs, and they make up a significant fraction of the extrasolar 'planetary' systems.  The interaction between Mira winds and giant planets or low-mass brown dwarfs with small orbital radii is an intriguing topic. This is especially true in light of the distortion observed in the outflow in the {\\it{o Ceti}} system, apparently as a result of its (more distant) white dwarf companion \\citep{kar97}. The class of symbiotic systems containing a Mira and a dwarf companion are also similar. Jura and collaborators \\citep[and refs. therein]{jur99} have studied other interesting giant star binary systems, with molecular reservoirs and dust rings.  We cannot hope to observe close planet systems in the same detail as these related systems. As we will describe below, observational signatures are likely to be indirect (see \\citet{wil01}).  Because planetary orbital periods (e.g., 3-30 yr.)  are roughly similar to wind crossing times over radial scales of about 1.0 au (.5 - 3.0 yr.), massive planets will create substantial wakes, and perturb the local wind flow. Within about 5 au, the motion of the extended stellar atmosphere is still primarily oscillatory.  The oscillating atmospheric gas elements and the propagating shocks interact with the planetary bow shock and wake, so the resulting flow is complex and very time-dependent. The planet accretes gas out of the stellar wind, but the amount of mass accreted over the AGB lifetime is very small compared to the planet's mass \\citep[and below]{wil01}. Gas drag on Jovian planets and brown dwarfs will also be small, unless they orbit very close to the star, i.e., within the atmosphere \\citep{wil02}.  Nonetheless, the accretion may release enough energy, in a suitable form, to give rise to observable events.  E.g., \\citet{wil01} argue that episodic optical flashes like those reported in the literature (e.g., \\citet{del98}, \\citet{sch91}, \\citet{ste01}) may be produced. \\citet{mou01} have also suggested that wake-like cometary 'exospheres' of giant planets, orbiting near their parent stars, might be observable. The planetary perturbation may affect other still observables, like the SiO maser emission that arises from the same region \\citep{str88}. To further understand both the observational signatures and the basic dynamics of the interaction, we have undertaken a program of hydrodynamic modeling of such systems, and present the first results below.  Ultimately, the observational signatures of planets or brown dwarfs in Mira winds could provide a new means of discovering such companions, and of learning more about their fates in the late stages of stellar evolution. ",
        "conclusions": ""
    },
    "0207/astro-ph0207583_arXiv.txt": {
        "abstract": "\\vskip-5.5in \\begin{flushright} UMN-TH-2106/02 \\\\ TPI-MINN-02/21 \\\\ astro-ph/0207583 \\\\ July 2002 \\end{flushright} \\vskip+4.4in  Recent observations by Bania \\etal (2002) measure \\he3 versus oxygen in Galactic \\hii\\ regions, finding that \\he3/H is within a factor of 2 of the solar abundance for [O/H] $\\ga -0.6$. These results are consistent with a flat behavior in this metallicity range, tempting one to deduce from these observations a primordial value for the \\he3 abundance, which could join D and \\li7 as an indicator of the cosmic baryon density. However, using the same data, we show that it is not possible to obtain a strong constraint on the baryon density range. This is due to (i) the intrinsically weak sensitivity of the primordial \\he3 abundance to the baryon density; (ii) the limited range in metallicity of the sample; (iii) the intrinsic scatter in the data; and  (iv) our limited understanding of the chemical and stellar evolution of this isotope. Consequently, the \\he3 observations correspond to an extended range of baryon-to-photon ratio, $\\eta = (2.2 - 6.5)\\times 10^{-10}$, which diminishes the role of \\he3 as a precision baryometer. On the other hand, once the baryon-to-photon ratio is determined by the CMB, D/H, or \\li7/H, the primordial value of \\he3/H can be inferred. Henceforth new observations of Galactic \\he3, can in principle greatly improve our understanding of stellar and/or chemical evolution and reconcile the observations of the \\hii\\ regions and those of the planetary nebulae. ",
        "introduction": "As the sole parameter of the underlying theory of big bang nucleosynthesis (BBN), the baryon-to-photon density ratio, $\\eta \\propto \\Omega_{\\rm B} h^2$, is one of the holy grails of cosmology.  In the past, $\\eta$ was best determined by the concordance of the four light element isotopes produced by BBN, D, \\he3, \\he4, and \\li7 (Walker \\etal 1991; Schramm \\& Turner 1998; Olive, Steigman, \\& Walker 2000; Nollett \\& Burles 2000; Cyburt, Fields, \\& Olive 2001; Coc \\etal 2002; Fields \\& Sarkar 2002).  As the systematic uncertainties in the abundance determination of each of these isotopes are becoming better understood, it seems that our ability to `predict' a precise value of $\\eta$ diminishes. Perhaps, our best hope for an accurate determination of $\\eta$ lies with the analysis of the microwave background anisotropy power spectrum. From this independent determination, we can certainly expect to gain a substantial amount of insight in the systematic effects involved in the abundance measurements (Kneller \\etal 2001; Cyburt, Fields, \\& Olive 2002). Of course, the concordance of BBN (within known uncertainties) remains a critical test of the standard cosmological model up to temperature scales of order $\\ga 1$ MeV. Needless to say, BBN continues to provide countless constraints on particle physics models which affect the evolution of the Universe at that epoch.  It is in this context that Bania \\etal (2002) have recently measured \\he3/H in about 20 \\hii\\ regions in the Galactic disk. These data show almost no variation over the metallicity range [O/H] = $-0.6$ to +0.2, i.e. a `plateau' in [\\he3/H] vs  [O/H]. Using recent developments of the stellar evolution theory, these authors have set an upper limit to the primordial \\he3 abundance of ${\\rm \\he3/H} \\le (1.1 \\pm 0.2) \\times 10^{-5}$, corresponding to $\\Omega_{b}h^{2}$ of about 0.02.  Bania \\etal\\ argue that the upper limit is robust despite the uncertainties in the details of \\he3 evolution; this robustness is argued to restore \\he3 as a BBN baryometer.  However, in the nineties severe doubt was cast regarding the use of \\he3 as a baryon density indicator due to the large uncertainties in its production in low mass stars. Standard stellar theory for low mass stars (see e.g. Iben \\& Truran, 1978) predicts a significant amount of production in these stars.  When incorporated into simple models of Galactic chemical evolution, one would expect \\he3 abundances in great excess from those observed (Vangioni-Flam \\etal 1994; Olive \\etal 1995; Galli \\etal 1995; Scully \\etal 1996, 1997; Dearborn, Steigman, \\& Tosi 1996). While it is quite possible that additional \\he3 destruction mechanisms (Charbonnel 1994, 1995, 1996; Hogan 1995; Wasserburg, Boothroyd, \\& Sackmann 1995) can lead to a consistent picture for the evolution of \\he3 (Olive \\etal 1997; Galli \\etal 1997), one must argue that the new process is not operative in all stars in order to avoid a contradiction between a few planetary nebulae showing high \\he3 abundances (from 2 to $10 \\times10^{-4}$, Balser \\etal 1997; Balser, Rood, \\& Bania 1999) and \\hii\\ regions with small \\he3 content (about $2\\times10^{-5}$, Balser \\etal 1999). The \\hii\\ region observations are in good agreement with the protosolar abundance value ${\\rm \\he3/H} = (1.5 \\pm 0.2) \\times 10^{-5}$ (Geiss \\& Gloeckler 1998, Gloeckler \\& Geiss 1998). The problem of \\he3 therefore seemed to be one  of stellar and/or Galactic in nature. For this reason, it was deemed to be a poor cosmological tracer.  Recently, the conclusion that \\he3 is not significantly produced in stars has been corroborated by a wealth of observations of \\c13 anomalies in low mass RGB stars and in some planetary nebulae. The RGB $\\c12/\\c13$ anomalies point out the existence of extra mixing process in low mass stars (Charbonnel \\& do Nascimento 1998; Sackmann \\& Boothroyd 1999a, 1999b and references therein). In this context, the mechanism responsible for the low $\\c12/\\c13$ observed in most of the RGB stars should lead to destruction of \\he3 in the external layers. The very large fraction of stars experiencing this phenomenon (of order 90 \\%) seems to prevent the overproduction of \\he3 in the course of Galactic evolution. In addition, the \\c12/\\c13 ratio has been recently observed in planetary nebulae. While Palla et al.\\ \\cite{pgmst} find high \\c12/\\c13 inconsistent with mixing in one object, Palla et al.\\ \\cite{pbstg} find that most of \\c12/\\c13 ratios in a sample of 14 planetary nebulae are low, implying that extra mixing has occurred. Balser, McMullin, \\& Wilson (2002) reach a similar conclusion in a study of 11 planetary nebulae. These theoretical and observational results thus represent impressive progress in resolving the \\he3 problem, and further studies are likely to clarify the situation more. This rapid progress has reopened the question of whether \\he3 can be restored as a probe of the cosmic baryon density. It will be shown below that the present results are not sufficient by themselves to constrain the Galactic evolution of \\he3 and hence its primordial abundance.  Consequently, it appears necessary to take a careful look at this problem, integrating the available data in a Galactic evolutionary model to reevaluate quantitatively the cosmological status of \\he3. In section 2, we briefly review the most recent developments in BBN that will be used in the present work. In section 3, the evolution of \\he3 is traced in the framework of a Galactic evolutionary model in order to analyze the potential production/destruction of \\he3 in stars. We find that present observations and theory cannot yet sufficiently constrain the  primordial abundance of \\he3 at the needed precision. In section 4, we take advantage of recent CMB observations (and the derived value of $\\Omega_B h^2$) to draw some consequences on primordial \\he3 and its evolution. Our summary and conclusions are given in section 5. ",
        "conclusions": "Considerable observational and theoretical progress has recently been made in understanding the stellar and Galactic evolution of \\he3. These studies have gone far to address the apparent dichotomy between the high \\he3 values measured in some planetary nebulae, and the much lower values seen \\hii\\ regions. Detailed stellar evolution models, and observations of carbon isotopes in red giants and planetary nebulae, strongly suggest that an extra mixing process is responsible for the destruction of \\he3 in 90\\% of low mass stars; thus the apparent observational inconsistency is removed. A key new contribution to this emerging picture are the Bania et al.\\ (2002) measurements \\he3 over a range of metallicity; the observed lack of strong \\he3 evolution with oxygen confirms that the production of \\he3 by low mass stars is very limited.  This progress in stellar physics notwithstanding, the remaining observational and theoretical uncertainties surrounding \\he3 evolution leave this isotope poorly-suited to be the precision baryometer found in \\li7 and especially D. On the other hand, \\he3 nevertheless remains extremely interesting as an astrophysical probe. More data will increase its value.  For example, an unbiased survey of \\he3 in planetary nebulae will shed light on whether \\he3 is increasing or decreasing from its primordial value. Moreover, the CMB results on $\\eta$ allow one to predict primordial abundances quite accurately for \\he3 and the other light isotopes. This will cast the \\hii\\ region \\he3 data in a new light, allowing these data to constrain both stellar and chemical evolution. Thus, we foresee a promising future for \\he3, and urge that the longstanding \\he3 observational program continue with renewed vigor."
    },
    "0207/astro-ph0207060_arXiv.txt": {
        "abstract": "We present results from the imaging portion of a far-ultraviolet (FUV) survey of the core of 47~Tucanae. We have detected 767 FUV sources, 527 of which have optical counterparts in archival HST/WFPC2 images of the same field.  Most of our FUV sources are main-sequence (MS) turn-off stars near the detection limit of our survey. However, the FUV/optical color-magnitude diagram (CMD) also reveals 19 blue stragglers (BSs), 17 white dwarfs (WDs) and 16 cataclysmic variable (CV) candidates. The BSs lie on the extended cluster MS, and four of them are variable in the FUV data. The WDs occupy the top of the cluster cooling sequence, down to an effective temperature of $T_{eff} \\simeq 20,000$~K. Our FUV source catalog probably contains many additional, cooler WDs without optical counterparts. Finally, the CV candidates are objects between the WD cooling track and the extended cluster MS.  Four of the CV candidates are previously known or suspected cataclysmics. All of these are bright and variable in the FUV. Another CV candidate is associated with the semi-detached binary system V36 that was recently found by Albrow et al. (2001). V36 has an orbital period of 0.4 or 0.8~days, blue optical colors and is located within 1~arcsec of a Chandra x-ray source. A few of the remaining CV candidates may represent chance superpositions or SMC interlopers, but at least half are expected to be real cluster members with peculiar colors. However, only a few of these CV candidates are possible counterparts to Chandra x-ray sources. Thus it is not yet clear which, if any, of them are true CVs, rather than non-interacting MS/WD binaries or Helium WDs. ",
        "introduction": "Globular clusters (GCs) are fantastic stellar crash test laboratories. Violent encounters between binaries and single stars in dense cluster cores give rise to exotic stellar populations, such as blue stragglers (BSs), cataclysmic variables (CVs), low-mass x-ray binaries and milli-second pulsars. All of these objects have considerably bluer spectral energy distributions than the non-degenerate stars that make up the bulk of the GC mass. Observations at short wavelengths are thus ideally suited to the detection and characterisation of these relatively rare stellar species. This is nicely illustrated by the recent Chandra survey of 47~Tuc (Grindlay et al. 2001a), which finally showed that the theoretically predicted, extensive population of interacting close binaries really does exist there. Nevertheless, important puzzles still remain. For example, even though Grindlay et al. (2001a) estimated that there are 30 CVs with $L_x \\gtappeq 10^{30}$~ergs~s$^{-1}$~in 47~Tuc --  far more than found in any previous study -- this is at best 1/3 of the number predicted by tidal capture theory (Di Stefano \\& Rappaport 1994).  Given the striking differences between optical and Chandra x-ray images of 47 Tuc, it seems clear that observations at intermediate wavelengths -- i.e. the far-ultraviolet (FUV) -- would be invaluable in identifying and characterising both the Chandra sources and other exotic objects in the cluster. For example, it is well known that CVs and young white dwarfs (WDs) radiate much of their luminosity in the FUV waveband. We have therefore embarked on a program to study the cores of nearby GCs at FUV wavelengths, using both imaging and slitless spectroscopy. The extraordinary spatial resolution and FUV sensitivity of the Space Telescope Imaging Spectrograph (STIS) aboard the Hubble Space Telescope (HST) make this the instrument of choice for our study. Here, we  present first results from a campaign on 47~Tuc, which clearly demonstrate the value of FUV observations to the study of GCs. ",
        "conclusions": "We have presented first results from the imaging part of a FUV survey of 47~Tuc. The great benefit of moving to the FUV is that most ordinary cluster members are too cool to show up in this bandpass. This dramatically reduces crowding and makes it easy to find hot and/or exotic objects such as CVs, BSs and young WDs. The same lack of crowding even allows us to carry out {\\em slitless}, multi-object FUV spectroscopy of the dense cluster core.  In this first analysis of our FUV observations of 47~Tuc, we have focused on the imaging aspect of our survey. By matching our FUV source catalog to a list of sources in a deep WFPC2/F336W ($\\simeq$~U) image of the same field, we have been able to find 17 WDs, 19 BSs and 16 candidate CVs. All four previously known CVs in our FoV -- including one which was only recently found on the basis of Chandra x-ray imaging (Grindlay et al. 2001a) -- are among these 16 candidates. Another one of our CV candidates is associated with the semi-detached binary system V36, which was recently disovered by Albrow et al. (2001). V36 has an orbital period of 0.4 or 0.8~days, blue optical colors and is located within 1~arcsec of a Chandra x-ray source. A few of the 11 remaining CV candidates may represent chance superpositions or SMC interlopers, but at least half are expected to be real cluster members with peculiar colors. However, only a few of these are located close to known x-ray sources. Our existing FUV spectroscopy and/or additional HST/WFPC2 images in different bandpasses should soon allow us test which, if any, of them are true CVs.  In summary, this first, deep FUV survey of a GC core highlights the great benefits of FUV observations to the study of rare and exotic stellar populations, not just in GCs, but also in other environments.  {\\em Note added:} The referee of this paper, Peter Edmonds, has informed us that unpublished astrometry carried out by himself and Ron Gilliland suggests a different identification for the Chandra source close to V36. We await to see details of this work in the literature.At the present time, we think even the remaining evidence -- V36's  blue FUV/optical colors, its orbital period of 0.4~or 0.8~days and Albrow et al.'s suggested classification of V36 as a semi-detached system -- is sufficient to support our classification of V36 as a strong CV candidate. We hope that our FUV spectroscopy will soon shed additional light on the nature of this intriguing system."
    },
    "0207/astro-ph0207310_arXiv.txt": {
        "abstract": "We present radiation-driven wind models for Luminous Blue Variables (LBVs) and predict their mass-loss rates. A comparison between our predictions and the observations of AG\\,Car shows that the variable mass loss behaviour of LBVs is due the recombination/ionisation of Fe\\,{\\sc iv}/{\\sc iii} and Fe\\,{\\sc iii}/{\\sc ii}. We also derive a present-day mass of 35 \\msun\\ for AG\\,Car. ",
        "introduction": "The strong winds of LBVs show a wide variety of mass-loss behaviour. During their S\\,Dor-type variations they expand in radius at approximately constant luminosity. In some cases the mass loss increases while the star expands (e.g. R\\,71), whereas for others (e.g. R\\,110) the behaviour is the exact opposite: as the star expands, its mass-loss rate drops. Recent radiation-driven wind models of OBA supergiants show that stars change their wind characteristics at spectral types B1 and A0, where \\mdot\\ jumps upwards by factors of five, due to Fe recombinations. In this poster, we investigate whether these ``bi-stability jumps'' can also explain $\\mdot(\\teff)$ of LBVs. ",
        "conclusions": "\\begin{itemize}  \\item LBV winds are driven by radiation pressure. \\noindent \\item The mass loss behaviour of LBVs (of up to over 0.5 dex) during their S\\,Dor-type variation cycles can be explained by the ionisation and recombination of Fe\\,{\\sc iv}/{\\sc iii} and Fe\\,{\\sc iii}/{\\sc ii}. noindent \\item The \\mdot(\\teff) behaviour of AG\\,Car can be matched when we adopt a mass of 35 \\msun.  \\end{itemize}"
    },
    "0207/astro-ph0207126_arXiv.txt": {
        "abstract": "We show that in a system of two planets initially in nearly circular orbits, an impulse perturbation that imparts a finite eccentricity to one planet's orbit causes the other planet's orbit to become eccentric as well, and also naturally results in a libration of their relative apsidal longitudes for a wide range of initial conditions. We suggest that such a mechanism may explain orbital eccentricities and apsidal resonance in some exo-planetary systems. The eccentricity impulse could be caused by the ejection of a planet from these systems, or by torques from a primordial gas disk. The amplitude of secular variations provides an observational constraint on the dynamical history of such systems. ",
        "introduction": "Orbital eccentricities in exo-planetary systems\\footnote{http://www.exoplanets.org} discovered thus far are often surprisingly large and have proven to be a major puzzle in understanding these systems. At least two systems with multiple planets in eccentric orbits ($\\upsilon$ Andromedae and HD 83443) are suspected of exhibiting secular apsidal resonance (Chiang, Tabachnik \\& Tremaine 2001, Wu \\& Goldreich 2002). Apsidal resonance is the phenomenon of phase-locking of the apsidal longitudes of two orbits, such that the two planets have a common average rate of apsidal precession and the angular difference of their apsidal longitudes, $\\Delta\\varpi=\\varpi_1-\\varpi_2$, librates around 0.  We describe here a dynamical mechanism for establishing apsidal resonance in a pair of planets that are initially on nearly circular orbits. We use classical analysis from celestial mechanics to show that an impulse perturbation that imparts an eccentricity to one of the orbits, excites the other planet's eccentricity on a secular timescale and also results in the libration of the relative apsidal longitude for a wide range of initial conditions.  A plausible cause of an impulse perturbation is the ejection of a planet from the system; torques from an exterior primordial disk may also cause such an eccentricity impulse. ",
        "conclusions": "We have shown that in a two planet system, an impulse perturbation on the eccentricity of the outer planet can excite the eccentricity of the inner planet on a secular timescale, and will result in apsidal libration with a $\\sim\\!50\\%$ probability. There are two significant characteristics of the secular dynamics following an eccentricity impulse, whether or not it results in apsidal resonance: (i) the inner planet's eccentricity has large amplitude variation, and it drops to its initial small value periodically; (ii) only a small fraction of time is spent at large values of $|\\Delta\\varpi|$ and small values of $e_1$ (see Fig.~1). Therefore, observing small values of $|\\Delta\\varpi|$ and large values of $e_1$ is not surprising if the system has suffered an impulse perturbation in its history and the secular dynamics indicates large amplitude variations.  We note that in the opposite limit of a slow adiabatic perturbation which increases one planet's eccentricity on a timescale much longer than the secular timescale, apsidal resonance will occur with nearly 100\\% probability and the resulting libration amplitude will be very small (for initially circular orbits). It is rather remarkable that the secular dynamics of two planets admits a high probability of apsidal resonance in both the impulse and the adiabatic limit of eccentricity perturbation of one planet. We defer detailed analysis to a future study, but we note here that the apsidal libration amplitude provides an observational diagnostic: large amplitudes favor the impulse mechanism whereas small amplitudes favor the adiabatic mechanism.  One possible mechanism for the eccentricity impulse is {\\it planet-planet scattering}. Highly eccentric orbits can be produced by gravitational interactions of planets initially on circular orbits that are close to the threshold of orbital instability. This has been proposed as a possible explanation for the eccentric orbits of exo-planets (Rasio \\& Ford 1996, Weidenschilling \\& Marzari 1996, Lin \\& Ida 1997, Ford et al.~2001, Marzari \\& Weidenschilling 2002). There are three possible outcomes of such interactions: a final stable configuration different from initial, the merger of two planets, or the ejection of one planet. Numerical experiments thus far have explored only a small fraction of the parameter space of two and three giant-planet systems. In systems of two equal mass planets, comparable branching ratios for the three possible outcomes are found; the two former outcomes do not result in eccentricity excitation, but the third -- ejection of a planet -- does, with typical eccentricity of the surviving planet in the range 0.4--0.8 (Ford et al.~2001). One problematic aspect is that two planet systems must have initial orbits fine-tuned close to the threshold of dynamical instability, i.e. orbital separation close to $2\\sqrt{3}$ times their mutual Hill radius (Gladman 1993). But this problem does not exist in systems of three or more planets, where chaotic instabilities can occur over a wider range of orbital separations; such systems can persist in quasi-stable orbital configurations for long periods of time, possibly exceeding $10^8$ years, before becoming dynamically unstable; then the most common outcome is the ejection of one planet (Marzari \\& Weidenschilling 2002).  Also, in the more realistic case of unequal masses, the probability of planet ejection is likely to be enhanced further.  We consider the ejection into an unbound orbit of a third planet by gravitational scattering from the outer of the two planets whose secular dynamics we have analyzed in section 2. (For comparable masses, ejection by scattering from the outer of the two planets is more likely than from the inner planet because the escape velocity is smaller at larger distance from the star: $v_{\\rm esc}\\sim r^{-1/2}$; it is also likely to produce the least direct perturbation to the inner members of the system.) Assuming that the ejected planet leaves on a nearly parabolic orbit, energy and angular momentum conservation yield the following estimate for the final eccentricity of the bound planet: \\begin{equation} 1 - e_{\\rm 2f}^2 \\simeq (1+{m_3\\over m_2} {a_2\\over a_3})\\Big[ 1+ {m_3\\over m_2}\\sqrt{{m_\\star+m_2\\over m_\\star+m_3}} \\Big(\\sqrt{a_3\\over a_2}-\\sqrt{2q_3\\over a_2}\\Big) \\Big]^2, \\end{equation} where $m_3,a_3,q_3$ are the mass, initial semimajor axis and final periastron distance of the ejected planet; $a_2$ is the initial semimajor axis of $m_2$, and is related to its final semimajor axis by energy conservation: $a_{\\rm 2f} \\approx m_2a_2a_3/(m_2a_3+m_3a_2)$. Unfortunately, there is no additional dynamical constraint or simple argument to constrain the final periastron distance of the ejected planet. So the final eccentricity of the bound planet cannot be predicted analytically and must be determined numerically.  Still, we can make the following estimate: since the periastron distance of the ejected planet is likely to be not greatly in excess of the apoastron distance of the surviving planet, then, for $q_3/a_2\\simeq$ 1--1.5, the final eccentricity of the bound planet will be in the range of 0.2--0.7 for $m_3/m_2\\simeq$ 0.1--0.7.  Once a dynamical instability sets in, the subsequent evolution is highly chaotic and unpredictable in detail. A simplified description, based upon the presently available numerical simulations, is as follows: at first, close encounters between $m_2$ and $m_3$ lead to a rapid increase in their orbital eccentricities, over $\\sim\\!10-10^2$ orbital periods; this is followed by a longer period, $\\gtrsim 10^4$ orbits, during which many weak distant encounters gradually and stochastically increase the orbital period and apoastron of $m_3$, the planet-to-be-ejected. In the present context, the initial eccentricity excitation of the surviving planet, $m_2$, satisfies the impulse approximation as the timescale of the first stage is much shorter than the secular timescale (which is of order $(m_{1,2}/m_\\star)^{-1}$ times the orbital period, or $\\sim\\!10^3$ orbital periods for jovian mass planets). Thus, the secular dynamics of the two surviving planets, $m_1$ and $m_2$, would be as described in Section 2. The weak distant encounters that eventually lead to the ejection of the third planet would produce small perturbations to the secular solution, as would the gravitational interactions between $m_1$ and $m_3$. (Neglecting the latter is justified, in the lowest order, as the $m_1$,$m_3$ separation is much greater than the $m_2$,$m_3$ separation.) The effects of these perturbations on the secular dynamics presented above will be considered in a future study. \\medskip  \\noindent{\\it The case of $\\upsilon$ Andromedae}\\quad The analysis given in the previous section can be used to estimate the pre-impulse orbital eccentricity of at least one of the planets. A good candidate amongst the known exo-planetary systems to apply this theory is the $\\upsilon$ Andromedae system, which represents perhaps the most (dynamically) constrained of the known exo-planetary systems. A recent orbital solution for this system (as quoted in Chiang et al.~2001) gives present eccentricities of planets C and D of 0.25 and 0.34, respectively, and yields secular dynamics for these two planets similar to that shown in Fig.~1(e,f), with apsidal libration amplitude of 25$^\\circ$--35$^\\circ$ (for relative orbital inclinations $\\la 20^\\circ$). An independent orbital solution obtained previously by Stepinski et al.~2000 (from an older and therefore smaller set of observational data) has orbital parameters within $\\sim\\!2\\sigma$ uncertainties of the more recent solution; it yields secular dynamics similar to that shown in Fig.~1(a,b), with apsidal libration amplitude of 80$^\\circ$--90$^\\circ$. The libration period for this system is $\\sim7\\times10^3$ yr. Now, if we allow that planet D suffered an eccentricity impulse, $e_D\\approx0.35$, in its history, we can estimate that the initial eccentricity of planet C could have been as small as $\\sim\\!0.07$ if the apsidal libration amplitude is 25--35 deg, but even smaller, 0--0.01, if the amplitude is 80--90$^\\circ$.  After submitting this paper, we became aware of a recent preprint by Chiang \\& Murray (2002) who propose an adiabatic eccentricity perturbation to explain the apsidal resonance in the $\\upsilon$ Andromedae planetary system. They invoke torques from an exterior massive primordial disk to provide the adiabatic eccentricity excitation to the outermost planet D. However the adiabaticity of the perturbation is owed to a specific choice of disk parameters, which are not well constrained; a different, and perhaps equally plausible, choice of disk parameters may instead provide an impulse perturbation (E.~Chiang 2002, personal communication); then the secular dynamics described here would apply.  As mentioned previously, the apsidal libration amplitude provides an observational diagnostic: large amplitudes favor the impulse mechanism whereas small amplitudes favor the adiabatic mechanism. In the case of $\\upsilon$ Andromedae, there remain significant uncertainties in the orbital parameters, and at present it appears difficult to discriminate between the adiabatic and the impulse mechanisms. The 25--35 deg apsidal libration amplitude found by Chiang et al.~(2001) is neither clearly small nor clearly large; in both the adiabatic mechanism and the impulse mechanism, an initial eccentricity near 0.06--0.07 of planet C can account for this libration amplitude. On the other hand, the large libration amplitude of 80--90 deg found by Stepinski et al.~(2000) would be explained by initial $e_C\\lesssim 0.01$ in the impulse limit, but cannot be explained naturally in the adiabatic limit. We urge further observations and analysis to improve the accuracy and fidelity of orbital solutions of exo-planetary systems; this would help to constrain their dynamical history."
    },
    "0207/nucl-th0207004_arXiv.txt": {
        "abstract": "\\noindent Renormalization group methods can be applied to the nuclear many-body problem using the approach proposed by Shankar. We start with the two-body low momentum interaction $V_{\\text{low k}}$ and use the RG flow from the particle-hole channels to calculate the full scattering amplitude in the vicinity of the Fermi surface. This is a new straightforward approach to the many-body problem which is applicable also to condensed matter systems without long-range interactions, such as liquid $^3$He. We derive the one-loop renormalization group equations for the quasiparticle interaction and the scattering amplitude at zero temperature. The RG presents an elegant method to maintain all momentum scales and preserve the antisymmetry of the scattering amplitude. As a first application we solve the RG equations for neutron matter. The resulting quasiparticle interaction includes effects due to the polarization of the medium, the so-called induced interaction of Babu and Brown. We present results for the Fermi liquid parameters, the equation of state of neutron matter and the $^1$S$_0$ superfluid pairing gap.  \\vspace{0.5cm}  \\noindent{\\it PACS:} 21.65.+f;          % 71.10.Ay;          % 11.10.Hi \\\\        % \\noindent{\\it Keywords:} Neutron matter; Fermi liquid theory; Renormalization Group; Polarization Effects; Superfluidity ",
        "introduction": "Fermi liquid theory is a prototype effective theory, invented by Landau in the late 1950's~\\cite{Landau1,Landau2,Landau3}. In this theory the properties of normal Fermi liquids at zero temperature are encoded in a few parameters, which are related to the effective two-fermion interaction. Over the years, Fermi liquid theory has proven to be an extremely useful tool for studying normal Fermi liquids, in particular liquid $^3$He~\\cite{BaymPethick}. The main assumption is that the elementary, low-lying excitations of the interacting system are relatively long-lived quasiparticles, with properties resembling those of free particles. The remaining strength in the spectral function is distributed over modes that add up incoherently. In the language of the renormalization group, the incoherent background is integrated out into the quasiparticle interaction, which can be determined by comparison with experiment or by microscopic calculations.  Babu and Brown later realized~\\cite{BB} that, in order to satisfy the Pauli principle in microscopic calculations, one has to take into account not only the particle-hole channel considered by Landau, which gives rise to the propagation of zero sound, but also the exchange channel thereof. This is taken into account in the induced interaction, which incorporates the response of the many-body medium to the presence of the quasiparticles.  Almost a decade ago, Shankar revived the interest in Fermi liquid theory by developing a renormalization group (RG) approach to interacting Fermi systems (for an introduction see Shankar~\\cite{FLTandRG1} and the lecture notes of Polchinski~\\cite{FLTandRG2}). Since Fermi liquid theory is restricted to low-lying excitations in the vicinity of the Fermi surface, normal Fermi systems are amenable to the renormalization group, where one has precisely such a separation of modes in mind. Moreover, by taking the loop contributions from both particle-hole channels into account, the RG flow remains antisymmetric, i.e., at any step of the renormalization, the scattering amplitude obeys the Pauli principle and consequently the solution obeys the Pauli principle sum rules. Thus, the antisymmetry of the scattering amplitude, which was the guiding principle of Babu and Brown, is realized naturally in the RG approach. Recently, we obtained additional RG invariant constraints~\\cite{IIpaper}. These approximate constraints are analogous to those obtained by Bedell and Ainsworth for paramagnetic Fermi liquids~\\cite{BA}.  In this paper, we solve the one-loop RG equations in the particle-hole channels at zero temperature and in three dimensions. We work in the approximation that both particle-hole momentum transfers are small compared to the Fermi momentum. The justification for this approximation is two-fold: In Fermi liquid theory the long-wavelength excitations play a central role. Hence, our primary aim is to treat these correctly. Furthermore, for the case of nuclear or neutron matter, the dependence of the induced interaction on Landau angle is fairly weak~\\cite{BBfornuclmat,bigreport}. Thus, we expect an expansion in momentum transfers to be accurate at least for the Fermi liquid parameters.\\footnote{This may change when effects of the tensor interaction are included in the induced interaction.} In this approximation the treatment of the particle-hole phase space with cutoffs is enormously simplified. An RG analysis of a schematic model, including the complete particle-hole phase space in two dimensions, is discussed in~\\cite{BF}. Improved RG equations, beyond one-loop order, are presented in~\\cite{IIpaper}.  Here we apply the RG techniques to neutron matter, where the tensor force does not contribute in the S-wave. In this system, a further physical motivation for expanding in small momentum transfers over the Fermi momentum is provided by the unique low momentum nucleon-nucleon interaction $V_{\\text{low k}}$~\\cite{Vlowk,Vlowkflow}, which we employ as the starting point for the RG flow. As shown in Figs.~2 and~4 of Bogner {\\em et al.}~\\cite{Vlowkflow}, $V_{\\text{low k}}$ is significant only for relative momenta $k < 1.3\\,\\text{fm}^{-1}$. Moreover, from the evolution of $V_{\\text{low k}}$ shown in Fig.~5 of~\\cite{Vlowkflow} we deduce that the important momentum modes are around the pion mass $m_\\pi \\approx 0.7\\,\\text{fm}^{-1}$, since higher momenta do not renormalize $V_{\\text{low k}}$ considerably. Therefore, the typical momentum transfer is small compared to the Fermi momentum of neutron matter at nuclear matter density $\\kf = 1.7\\,\\text{fm}^{-1}$.  The main idea of the RG approach to the many-body problem is to ``adiabatically'' include the in medium corrections to the effective interaction by solving the RG flow equations in the relevant channels. In this exploratory calculation we include the particle-hole channels, which play a special role in Fermi liquid theory. The main effect of scattering in the particle-particle channel, the removal of the short-range repulsion, is taken care of by using $V_{\\text{low k}}$ as the starting point for the RG flow. The low-lying excitations in this channel, which are responsible e.g., for superfluidity, are not included. These are then treated explicitly in Section~\\ref{gapsec}, where we compute the superfluid gap by employing BCS theory for the particle-hole reducible scattering amplitude.  In the RG approach, the width of the momentum shell that is integrated out in one iteration is a ``small parameter''. The full scattering amplitude is an RG invariant quantity. In a conventional approach, this is calculated from a ``set of diagrams'' while in the RG approach it is continuously evolved from the starting interaction by gradually including the many-body corrections from the narrow momentum shells one at a time. In this sense the RG provides a method for dealing with strongly interacting systems, where perturbation theory fails. We stress that this application of the renormalization group to many-body systems is only indirectly related to the well known RG treatment of critical phenomena.  We derive the one-loop renormalization group equations for the quasiparticle interaction and the scattering amplitude at zero temperature. The evolution of the effective mass is included in the RG flow, as well as a simplified treatment of the renormalization of the quasiparticle strength. In the forward scattering limit, the role of the energy transfer is in the RG taken over by the cutoff in momentum space. We find that, in the long-wavelength limit, the dependence of the effective four-point vertex on $\\Lambda/q$ corresponds to the behavior with $\\omega/q$ in the microscopic derivation of Fermi liquid theory by Landau.  Finally, we present the solution of the RG equations for neutron matter. These are obtained by employing the unique low momentum nucleon-nucleon interaction as the initial condition of the RG. Our results include the Fermi liquid parameters in the density range of interest for neutron stars as well as the full scattering amplitude for general (non-forward) scattering on the Fermi surface. Using the Fermi liquid parameters, we compute the equation of state, including polarization effects, by integrating the incompressibility~\\cite{Kaellman}. Finally, we compute the $^1$S$_0$ superfluid pairing gap using weak coupling BCS theory. This is an application of our approach, which probes the angular dependence of the scattering amplitude. We generally find very good agreement with the results obtained in the polarization potential model by Ainsworth, Wambach and Pines~\\cite{WAP1,WAP2}. However, it is worth noting that much of the model dependence, inherent in the work of Ainsworth {\\em et al.}, can be avoided in the RG approach. ",
        "conclusions": "In this work, we have developed a practicable framework for renormalization group calculations of strongly interacting Fermi systems. We employed the RG approach to compute the quasiparticle interaction and the scattering amplitude on the Fermi surface for neutron matter, which for many purposes is a good approximation to neutron star matter, where there is a small admixture of protons. Our approach follows the ideas of Shankar in applying RG techniques to the interacting fermion problem~\\cite{FLTandRG1,Shankar}.  RG methods used to solve multi-channel problems have several advantages. First, all channels and the corresponding momentum scales can be treated on equal footing in a systematic way, which can accommodate fundamental symmetries, such as the Pauli principle. In this study, we considered the RG flow in the particle-hole channels at the one-loop level. An expansion in momentum transfers was performed, which enabled us to compute the beta functions analytically. We note that it is fairly straightforward to include the particle-particle channel in the RG equations for the scattering amplitude and the quasiparticle interaction. This opens the possibility to explore the interference of the particle-hole and particle-particle channels. A calculation of this kind corresponds to solving the ``parquet equations at the one-loop order''. The expansion in momentum transfers can be removed by computing the complete particle-hole phase space~\\cite{BF}.  In contrast to the traditional approach to the many-body problem, the RG adapts the starting vacuum interaction to include the in medium effects due to the fastest particle-hole excitations. In the subsequent iteration, this renormalized interaction is modified to account for the next fastest modes. Therefore, at every step, the flow of the effective interaction is expanded around the current system, which highlights the efficacy of the approach. Moreover, the RG flow is physically very transparent, since the effects of the high momentum modes are readily tractable. We believe that the RG method is a promising tool for studying a wide range of nuclear many-body problems.  At the one-loop level and at zero temperature the RG equations are particularly transparent. We found an ambiguous long-wavelength limit, analogous to that appearing in the original treatment of Landau, which can be used to distinguish between the quasiparticle interaction and the forward scattering amplitude. This implies that, at the one-loop level, it is necessary to solve only the RG equation for the scattering amplitude. The quasiparticle interaction is then obtained in the corresponding limit. A test for the consistency of the solution is provided by general relations between the forward scattering amplitude and the Fermi liquid parameters.  We solved the RG equations in the particle-hole channels for neutron matter, where tensor interactions are absent in relative S-states. The RG flow includes the exchange channel, and thus it builds up the polarization effects of the induced interaction of Babu and Brown. Our results for neutron matter are very encouraging and provide strong motivation for further developments of the RG approach to the nuclear many-body problem as well as for condensed matter systems.  In the calculation, polarization and self-energy effects have been treated self-consistently. The latter have been included in a simple approximation, which nevertheless led to good quantitative results at intermediate densities.  Results for the Fermi liquid parameters and the scattering amplitude for the density range of interest to neutron stars were presented. The Fermi liquid parameters satisfy the Pauli principle sum rules by construction. Moreover, polarization effects lead to an enhancement in the pressure and the energy per neutron below nuclear matter density. As a first application, which probes the momentum dependence of the scattering amplitude, we computed the $^1$S$_0$ superfluid gap using weak coupling BCS theory. A reduction from a direct gap of $3.3 \\, \\text{MeV}$ to $0.8 \\,\\text{MeV}$ was found. Our results confirm the importance of particle-hole polarization effects for superfluidity. A further application of the scattering amplitude at finite momentum transfers is in calculations of transport processes.  For a similar analysis of nuclear matter, it is necessary to extend the flow equations to incorporate tensor interactions. These are important for a complete discussion of nuclear matter, and we expect rather large renormalization effects from the arguments given in~\\cite{IIpaper}. Moreover, tensor correlations and the tensor Fermi liquid parameters play an important role in the physics of dense matter~\\cite{Olsson}. As noted in the introduction, tensor forces are less pronounced in neutron matter, since they are not operative in relative S-states. However, for specific quantities the tensor force may be of crucial importance.  To date, polarization effects on the $^3$P$_2$--$^3$F$_2$ pairing of neutrons have not been addressed and a quantitative analysis is very much needed. Pethick and Ravenhall argue that both density and spin-density fluctuations would increase the gap in this state and thus possibly extend the superfluid regime to lower densities~\\cite{3pf2}. Once the RG approach is adapted to allow for tensor forces, it should be straightforward to address this problem~\\cite{inprep}.  \\begin{ack} We thank Tom Kuo and Scott Bogner for providing us with the $V_{\\text{low k}}$ code and Kevin Bedell, Vincent Brindejonc and Janos Polonyi for rewarding discussions. AS thanks the Theory Group at GSI for their kind hospitality. The work of AS and GEB is supported by the US-DOE grant No. DE-FG02-88ER40388. \\end{ack}  \\appendix"
    },
    "0207/astro-ph0207195_arXiv.txt": {
        "abstract": "{We have derived the galaxy luminosity function (GLF) in the cluster of galaxies Abell 496 from a wide field image in the I band. A single Schechter function reproduces quite well the GLF in the 17$\\leq {\\rm I_{AB}} \\leq$22 ($-19.5\\leq {\\rm M_I} \\leq -14.5$) magnitude interval, and the power law index of this function is found to be somewhat steeper in the outer regions than in the inner regions. This result agrees with the idea that faint galaxies are more abundant in the outer regions of clusters, while in the denser inner regions they have partly been accreted by larger galaxies or have been dimmed or even disrupted by tidal interactions. ",
        "introduction": "\\label{sec:intro}  Galaxy luminosity functions (hereafter GLF) are fundamental to analyse the properties of galaxies in clusters. In a number of cases, it is impossible to fit the entire GLF with a single Schechter function: there appear to be two components in the GLF, one for the bright galaxies - a gaussian distribution, and another for fainter galaxies - a power law or a Schechter function (see e.g. Godwin \\& Peach 1977, Biviano et al. 1995, Durret et al. 1999a). This suggests that there are at least two populations of galaxies in clusters, which do not vary strongly from one cluster to another, since the dip between both curves falls roughly at the same absolute magnitude in several clusters (see e.g. Tab.~2 in Durret et al. 1999a).  Besides, at faint magnitudes, the slope of the GLF can be steeper in the outskirts of clusters i.e. in less dense environments, and flatter near the cluster center (Lobo et al. 1997, Driver et al. 1998, Adami et al. 1998, 2000, Andreon 2002, Beijersbergen et al. 2002). This can be interpreted as due to the fact that in dense environments, small (and faint) galaxies are more likely to be accreted by larger ones. Moreover, they have also probably suffered repeated tidal interactions on their way towards the cluster center, consequently being dimmed or even disrupted (see e.g. Moore et al. 1996, Phillipps et al 1998, Kajisawa et al. 2000).  We have performed a first analysis of the GLF of Abell 496 (Durret et al.  2000) and intend to reobserve Abell 496 spectroscopically with the VLT and VIRMOS; as a preparation, we asked R. Ibata and C. Pichon to obtain for us a wide image of this cluster with the CFH12K camera at CFHT. We present below our analysis of the GLF in different regions at various distances from the cluster center. ",
        "conclusions": "We have derived the GLF in various regions of Abell 496.  The slope of the Schechter function fit is always found to be steep (between $-1.60$ and $-2.05$). Since such a steep slope could be due to several artefacts, we will discuss the validity of our results. First, the background counts could have been underestimated.  However, the good agreement of the various background counts (VIRMOS, Postman, Cabanac), and the fact that the subtraction is mainly that of the VIRMOS counts (in the interval 18.5$< {\\rm I_{AB}} \\leq$22), made with the same instrument, filter and magnitude system as ours, tends to suggest that this is not the case. Second, the number of faint galaxies may have been overestimated; for example, we may have confused globular clusters with galaxies at faint magnitudes, as explained in detail by Andreon \\& Cuillandre (2002). However, we limit our sample to ${\\rm I_{AB}}$=22, where such effects should not be too strong. Third, our ${\\rm I_{AB}}$ magnitudes may be too bright by 0.25 magnitude, as suggested by the difference with the Moretti et al. data (see Sect. 2.1). We tried to fit the GLF in several regions after shifting the ${\\rm I_{AB}}$ magnitudes by 0.25 (before subtracting the background) and find slopes $\\alpha= -1.68\\pm0.05, -1.85\\pm 0.03$ and $-1.94\\pm 0.04$ for regions CDIJ, ABK and EFGH respectively, instead of the previous values of $-1.75, -1.93$ and $-1.98$. Therefore, although the values change a little, the slope remains flatter in CDIJ.  We then used our ${\\rm I_{AB}}$ catalogue limited to the region in common with Molinari et al. (1998) and made a Schechter fit as described above. The Molinari et al. zone partially covers our CCDs J,I,C,D,E and F, with a main concentration towards CCD E. A Schechter fit in the same magnitude interval gives a slope $\\alpha=-1.71 \\pm 0.06$, close to the value of $-1.68\\pm 0.13$ found in CCD E. A shift of ${\\rm I_{AB}}$ by 0.25 magnitude as above gives $\\alpha=-1.68\\pm 0.09$, in perfect agreement with Molinari.  Fits in broader magnitude intervals give: $\\alpha=-1.66\\pm 0.05, -1.64\\pm 0.04$ and $-1.57\\pm 0.03$ for the magnitude intervals 17$\\leq {\\rm I_{AB}} \\leq$22.5, 17$\\leq {\\rm I_{AB}} \\leq$23, and 17$\\leq {\\rm I_{AB}} \\leq$23.5 respectively. Molinari et al. give a slope $\\alpha=-1.49\\pm 0.04$ in the i band, but mention that a magnitude correction allows them to reach a slope as steep as $\\alpha=-2.0$. We therefore believe that our results are consistent with theirs.  Note that such a slope is not much steeper than found e.g. in Coma (Lobo et al. 1997) or in Abell 665 (De Propris et al. 1995). This could indicate an excess of faint red galaxies, but to ascertain this hypothesis it would be necessary to derive the GLF in Abell 496 in other filters from samples of comparable quality (covered area and depth).  As still another test on the robustness of our results, we reanalyzed the GLF in the CDIJ region. For this, we reduced the number counts by 10, 20, 30 40 and 50\\% for ${\\rm I_{AB}} >$20 (below ${\\rm I_{AB}}$=20 the counts remained unchanged) and made fits of these new GLFs.  Results are given in Table 2. They show that the slope changes strongly only if the counts are reduced by at least 30\\%, an unrealistic number. Besides, in order to account for the difference in slope of 0.2 that we observe for example between regions CDIJ and EFGH, we would need to make an unrealistically large error of 50\\% on the counts.  We are therefore confident that both the absolute values of $\\alpha$ and their variations from one zone to another are robust.  Our second result is that the slope of the Schechter function fit tends to be steeper in the outer regions of the cluster, as already observed in other clusters (see references in Sect.~\\ref{sec:intro}). Such a variation of $\\alpha$ can be interpreted as due to the fact that faint galaxies are accreted by larger ones preferentially in the inner parts of clusters, where the galaxy density is higher, therefore inducing a lack of faint galaxies and a flattening of the GLF in the inner regions.  Moreover, galaxies are likely to have suffered repeated tidal interactions on their way towards the cluster center, consequently being dimmed or even disrupted in a scenario of galaxy harassment (Moore et al. 1996).  The next step is obviously to confirm these results through deep multiband imaging and/or spectroscopy that would make the background subtraction more secure and would allow us to compare the GLFs in various filters."
    },
    "0207/astro-ph0207530_arXiv.txt": {
        "abstract": "The distribution and kinematics of neutral hydrogen have been studied in a wide area around  the supernova remnant W28. A 2\\rlap{$^\\circ$}.5 $\\times$ 2\\rlap{$^\\circ$}.5 field centered at $l$ = 6\\rlap{$^\\circ$}.5, $b$ = 0$^\\circ$ was surveyed using the Parkes 64-m radio telescope (HPBW 14\\rlap{$^\\prime$}.7 at $\\lambda$ 21 cm).  Even though W28 is located in a complex zone of the Galactic plane, we have found different \\HI\\ features which are evidence of the interaction between W28 and its surrounding gas. An extended cold cloud with about 70  M$_{\\odot}$ of neutral hydrogen was detected at the location of W28 as a self-absorption feature, near the LSR velocity + 7 \\ks. This \\HI\\ feature is the atomic counterpart of the molecular cloud shown by previous studies to be associated with W28. From this detection, we can independently confirm  a kinematical distance of about 1.9 kpc for W28. In addition, the neutral hydrogen observed in emission around the SNR displays a ring-like morphology in several channel maps over the velocity interval [--25.0, +38.0] \\ks. We propose that these features are part of an interstellar HI  shell that has been swept-up by the SN shock front. Emission from this shell is confused with unrelated gas. Hence, we derive an upper limit for the shell mass of 1200 -- 1600  M$_{\\odot}$, a maximum radius of the order of 20 pc, an expansion velocity of $\\sim$ 30 \\ks, an initial energy of about 1.4 -- 1.8 $\\times 10^{50}$ ergs and an age of $\\sim$ 3.3 $\\times 10^4$ yrs. The pre-existing ambient medium has a volume density of the order of 1.5 -- 2 cm$^{-3}$. W28 is probably in the radiative evolutionary phase, although it is not possible to identify the recombined thin neutral shell expected to form behind the shock front with the angular resolution of the present survey. ",
        "introduction": "Each supernova remnant (SNR) is the unique product of its own history (the progenitor and the explosion mechanism) and the characteristics of the environs in which it evolves. The study of the interstellar medium around SNRs can be used to understand  the appearance of a remnant in different spectral regimes (distorted shapes, local brightness enhancements, filamentary emission, etc.). Such studies also allow the analysis of the  temporal evolution of SNRs. In addition, the investigation of the gaseous matter around SNRs can lead to an understanding of the Galactic interstellar medium. These studies are important in understanding the response of the interstellar gas to the large injection of energy and momentum that a supernova (SN) explosion represents.  Numerous investigations of interaction of SNRs with the surrounding ISM have been made using atomic and molecular lines (Routledge et al. 1991, Pineault et al. 1993, Wallace et al. 1994, Frail et al. 1994, 1996 and 1998, Reynoso et al. 1995,  Dubner et al. 1998a and b). These investigations show the manner in which the expansion of a SN shock front modifies the surrounding environment and the effect that the surrounding gas has, in turn,  on  the shape and dynamics of the SNR. In the present study, we report the results of an \\HI\\ study around the SNR W28 (G 6.4--0.1).  The SNR W28 is located in a very complex region of the Galaxy, near the large HII regions M8 and M20, and the young clusters NGC 6530, NGC 6514, and Bo 14. It has a number of prominent morphological characteristics. In the radio continuum, there is diffuse emission together with thin filaments and small bright regions, as seen in Figure 1. This image is the result of combining 50 VLA pointings into a 20 cm mosaic (Dubner et al. 2000). In X-rays, diffuse thermal emission fills the interior of W28, although ear-shaped segments of a limb-brightened shell can also be observed toward the NE and NW (Rho et al. 1996). In the optical, there are bright narrow filaments strongly correlated with radio features and diffuse H$\\alpha$ nebulosities, apparently anti-correlated with the radio synchrotron emission (Long et al. 1991, Dubner et al. 2000).   A number of observations support the existence of a physical interaction between W28 and an adjacent molecular cloud: $(1)$ the existence of shocked CO and other molecular species (Wootten 1981, Frail \\& Mitchell 1998, Arikawa et al. 1999) $(2)$ the detection of over forty 1720 MHz OH masers distributed along the brightest synchrotron features (Claussen et al.  1997 and 1999), and $(3)$ the coincidence of the molecular gas with the brightest synchrotron filaments which, are the features with the flattest spectral index in the SNR (as expected for high Mach number shocks; from Dubner et al. 2000). All these indicators point to the existence of an interaction between the SNR and the molecular cloud.  In what follows, we analyze the distribution of the neutral hydrogen around W28, based on a survey of the $\\lambda$21 cm \\HI\\ line carried out for a 2\\rlap{\\arcdeg}.5 $\\times$ 2\\rlap{\\arcdeg}.5 field with the Parkes 64--m radio telescope. ",
        "conclusions": "We have carried out a study of the neutral hydrogen in the environs of the SNR W28.  Our analysis of the kinematics and distribution of the \\HI\\  has revealed several \\HI\\ features that are most probably with W28, revealing signatures of the interaction of this SNR with the interstellar medium. We have detected, as a self-absorption feature around $\\sim$ 7 \\kms,  the neutral gas counterpart of the molecular cloud detected by Arikawa et al. (1999) in unshocked CO gas. Based on the presence of this cold cloud, we can independently confirm a kinematical distance of 1.9$\\pm$0.3 kpc for W28.  Portions of an incomplete \\HI\\ shell are also observed in emission at different positive and negative LSR velocities, with a maximum angular size ($\\sim 0^\\circ.6$) at V$_{LSR}$=+17.5\\kms. An \\HI\\ cloud is detected near V$_{\\rm LSR} \\sim +37$ \\kms overlapping the center of W28. We interpret this last feature as the ``cap'' of the irregular expanding interstellar shell swept-up by the W28 shock front. The mass of this shell has been estimated to be between 1200 and  1600 M$_{\\odot}$.  Based on the present results, the following scenario for W28 can be proposed:  (a) A SN explosion of energy $\\sim$ 1.6 $\\times 10^{50}$~ergs occurred about 3.3 $\\times 10^4$~yr ago, at the position ({\\it l, b}) = (6\\arcdeg30\\arcmin , $-$0\\arcdeg 12\\arcmin) and at distance of $\\sim$~1.9 kpc. At this location, the ambient density of the ISM was $\\sim$1.5 -- 2 cm$^{-3}$.  (b) The expanding shock wave has collided with a cold gas concentration, observed as an absorption \\HI\\ feature and as molecular clouds around  the LSR velocity of 7 \\kms.   The mass of this cold cloud ($\\sim $ 70 M$_\\odot$), is only a small fraction of the total mass estimated for  molecular hydrogen. Most of the atomic hydrogen is detected in emission as an \\HI\\ shell, as mentioned below.  (c) The interaction of the SN shock front with the surrounding \\HI\\ gas has swept-up a thick interstellar HI shell, presently expanding at $\\sim$ 30 \\kms. W28 has entered into the radiative stage of evolution about 2$ \\times 10^4$ yrs ago. However, the thin neutral shell expected to form by recombination behind the shock front could not be identified because of the relatively low angular resolution of the present study."
    },
    "0207/astro-ph0207139_arXiv.txt": {
        "abstract": "{ We investigate the possibility of measuring the Hubble constant, the fractional energy density components and the equation of state parameter of the ``dark energy'' using lensed multiple images of high-redshift supernovae.  With future instruments, such as the SNAP and NGST satellites, it will become possible to observe several hundred lensed core-collapse supernovae with multiple images. Accurate measurements of the image separation, flux-ratio, time-delay and lensing foreground galaxy will provide complementary information to the cosmological tests based on, e.g., the magnitude-redshift relation of Type Ia supernovae, especially with regards to the Hubble parameter that could be measured with a statistical uncertainty at the one percent level. Assuming a flat universe, the statistical uncertainty on the mass density is found to be $\\sigma^{\\rm stat}_{\\om} \\lsim 0.05$.  However, systematic effects from the uncertainty of the lens modeling are likely to dominate. E.g., if the lensing galaxies are extremely compact but are (erroneously) modeled as singular isothermal spheres, the mass density is biased by $\\sigma^{\\rm syst}_{\\om} \\sim 0.1$.  We argue that wide-field near-IR instruments such as the one proposed for the SNAP mission are critical for collecting large statistics of lensed supernovae. ",
        "introduction": "Gravitational lensing of high-$z$ objects has been used in the past as a tool for deriving cosmological parameters with mixed success. While the fraction of quasars with multiple images and the distribution of image separations may be used to probe the vacuum energy density, $\\Omega_\\Lambda$ \\citep{turner90,sef}, the method suffers from severe systematic uncertainties related to the lens modeling and the results are as of yet inconclusive \\citep{fukugita92,maoz,kocha96,park,chiba}.  On the other hand, the method proposed by \\citet{refsdalb,refsdalc} using time delay measurements of multiple imaged quasars to constrain $H_0$ has provided results that are in good agreement with independent techniques, \\citep[see e.g.][]{koopmans99,brown}. Thus, distances derived from geometrical measurements of lensed high-redshift sources may be a viable way to gain further knowledge of cosmological parameters although there are still unresolved issues concerning the modeling of the lenses. E.g. in a recent paper \\citep{kocha02}, five gravitational lens systems are analyzed and it is suggested that, unless the dark matter distributions in the lensing galaxies are rather compact, the derived value of $H_0$ is too low in comparison with the local measurements.  In this note we investigate how to use multiply imaged high-redshift supernovae (SNe) to constrain the Hubble parameter, the mass and dark energy density of the universe, $\\om$ and $\\olx$, as well as the equation of state parameter, $w_0 = {p_X \\over \\rho_X}$, which we assume is constant for $z\\lsim 5$. SNe are well suited for this technique because of the expected high rate at high-redshifts and most importantly because of their well known lightcurves. In particular, the (rest-frame) optical lightcurves of core-collapse (CC) SNe show fast rise-times, typically about 1 week long. Thus, the time difference between images can be measured to better than one day's precision.  Another possibility is the use of the UV shock-breakout. As has been seen in SN 1987A and modeled in \\citet{ensman} the shock breakout can serve as a time stamp with a precision of just minutes. The entire UV flash occurs over a period of minutes to several hours in the rest-frame of the SN depending on the nature of the progenitor.  Wide field optical and NIR deep surveys, such as the planned SNAP satellite \\citep{snap}, have the potential to discover $\\sim10^6$ CC SNe. While these SNe are sometimes regarded as a ``background'' for the primary cosmology program based on Type Ia SNe, we argue that lensed SNe of {\\em any} kind may also provide useful information on cosmological parameters.  In \\citet{holz} the number of multiply imaged CC SNe up to $z < 2$ were estimated by simply scaling the Type Ia rate by a factor of 5. In this work we extend the considered redshift up to $z=5$ using a SN rate calculation derived from the star formation history. Further, we take into account the NIR wide field instrumentation in the current design of the SNAP satellite.  Using simple toy-models, we investigate the accuracy of the strong lensing technique to improve our knowledge of cosmological parameters. Our observables are the source redshift, the redshift of the lensing galaxy, the image-separation, $\\Delta\\theta$, the time-delay, $\\Delta t$, and the flux-ratio, $r$.  We derive the relation of our observables and cosmological parameters for different matter distributions in order to investigate the sensitivity to the choice of lens model.  We speculate, that if the measurement is done using a very large sample of lensed systems, useful bounds on the cosmological parameters can be found in spite of large uncertainties in the lens model.  We have used the SNOC Monte-Carlo simulation package \\citep{snoc} to estimate the rate and measurable quantities of multiple image SNe, e.g. the distribution of time-delays, image-separations and flux-ratios. We also simulate extinction by dust both in the host galaxy of the SN and in the foreground lensing galaxy.  While about half of the known lensed systems exhibit more images \\citep[see e.g.][]{keeton98}, this work is limited to spherically symmetric lensing systems producing only two or ring-like images. Systems with more lensed images are potentially very interesting as they provide more measurables which can be used to constrain the lens model. ",
        "conclusions": "Strongly gravitationally lensed SNe could be detected in large numbers in planned wide field, deep, SN search programs, probably on the order of several hundred. In particular, we argue in favor of large NIR imagers in space missions like SNAP.  Lensed SNe are potentially interesting as they provide independent measurements of cosmological parameters, mainly $H_0$, but also the energy density fractions and the equation of state of dark energy. The results are independent of, and would therefore complement the Type Ia program. At the faint limits considered in this note, several quasars per square arcminute are expected \\citep{quasars}. Thus, we expect that an instrument like SNAP would find several hundred multiply imaged QSOs, in addition to the strongly lensed SNe. Thus, the statistical uncertainty could become smaller than what we have considered here.  While the systematic uncertainties remain a source of concern, we show that the simplest spherically symmetric models introduce moderate biases ($\\sigma^{\\rm syst}_{\\om} \\lsim 0.1$) , at least as long as multiple images with similar fluxes are considered.  While projections of different SNe could in principle be interpreted as a lensed SN the two scenarios may be distinguished. The signatures of CC SNe are unique. Crudely, the lightcurves are a product of the progenitor mass, the mass loss during its evolution off the main sequence, the amount of radioactive Ni synthesized during the explosion and the kinetic energy imparted to the ejecta. In addition, the environments the SNe explode in (the density and structure of their local interstellar medium), often play a significant role in what we eventually see of the CC event. These differences have given rise to all the different classifications of these events we currently have; Type IIP, IIL, Ib, Ic, IIn, etc. Given all this diversity it makes it quite easy to distinguish one event from another and to not confuse a lensing event with a coincident CC event along the same line of sight.  In our simulations we find that extinction in the foreground galaxies does not severely affect the detectability of multiple lensed events nor the ability to derive the flux-ratio between the images. Clearly, multi-band observations of the SNe will be important in order to correct for the different amounts of extinction of the images. At the same time, the data could provide important results on the dust properties of the foreground lensing galaxies, similar to the studies done with multiple imaged quasars at $z_d\\lsim 1$ \\citep{falco}.  Unlike quasars, CC events have a very deliberate color evolution along their lightcurves. In general, from the moment of shock-breakout onwards, the atmospheres of CC SNe expand and become cooler and redder. This signature not only helps in the relative timing of these events, but also allows us to measure the differential extinction due to the varying amounts of dust along the lensed paths to the SN quite well. With 3 or more filters one could make a measurement of both the total extinction relative to the bluest event in addition to the ratio of the total to selective extinction due to differences in the dust properties. Furthermore, one could enhance this method by taking a spectrum of the SN (any of the lensed events would do) at a given epoch and through spectrum synthesis derive the true, unextinguished spectral energy distribution of the event [see e.g. \\citep{mitchell,baron}]."
    },
    "0207/astro-ph0207249_arXiv.txt": {
        "abstract": "Black hole mass, along with mass accretion rate, is a fundamental property of active galactic nuclei. Black hole mass sets an approximate upper limit to AGN energetics via the Eddington limit. We collect and compare all AGN black hole mass estimates from the literature; these 177 masses are mostly based on the virial assumption for the broad emission lines, with the broad-line region size determined from either reverberation mapping or optical luminosity. We introduce 200 additional black hole mass estimates based on properties of the host galaxy bulges, using either the observed stellar velocity dispersion or using the fundamental plane relation to infer $\\sigma$; these methods assume that AGN hosts are normal galaxies. We compare 36 cases for which black hole mass has been generated by different methods and find, for individual objects, a scatter as high as a couple of orders of magnitude. The less direct the method, the larger the discrepancy with other estimates, probably due to the large scatter in the underlying correlations assumed. Using published fluxes, we calculate bolometric luminosities for 234 AGNs and investigate the relation between black hole mass and luminosity. In contrast to other studies, we find no significant correlation of black hole mass with luminosity, other than those induced by circular reasoning in the estimation of black hole mass. The Eddington limit defines an approximate upper envelope to the distribution of luminosities, but the lower envelope depends entirely on the sample of AGN included. For any given black hole mass, there is a range in Eddington ratio of up to three orders of magnitude. ",
        "introduction": "Black holes have been the leading candidate to power the central engines in AGN for over three decades (Lynden-Bell 1969), but direct evidence for their presence has been elusive. In nearby galaxies, spatially resolved kinematics have provided strong evidence for the ubiquity of nuclear black holes, with dynamical black hole detections reported for 37 galaxies (Kormendy \\& Gebhardt 2001). Such observations are available only for a handful of the nearest AGN (Harms et al. 1994, Miyoshi et al. 1995, Greenhill et al. 1996).  Black hole mass, along with mass accretion rate, is a fundamental property of AGN. Via the Eddington limit, a maximum luminosity for the idealized case of spherical accretion ($L_{Edd} = 1.25 \\times 10^{38} \\times M_{BH}/M_{\\odot}$ ergs s$^{-1}$), the black hole mass sets an approximate upper limit to AGN energetics. It is also the integral of the accretion history of the AGN. However, direct kinematic observations of the black hole mass are limited by finite spatial resolution (a typical AGN at redshift 2 would require nano-arcsecond resolution to probe the sphere of influence of the black hole), not to mention that scattered light from the bright central source dilutes any kinematic signal from orbiting material.  For these reasons, various less direct methods for estimating black hole mass have been devised. One set of methods (\\S\\S~2.1,2.2) assumes the broad-line region (BLR) is gravitationally bound by the central black hole potential, so that the black hole mass can be estimated from the orbital radius and the Doppler velocity. The reverberation mapping technique utilizes the time lag between continum and emission lines to derive the distance of the BLR from the black hole (Blandford \\& McKee 1982, Peterson 1993). About three dozen AGN black hole masses have been measured using this technique. A less costly alternative is to infer the BLR size from the optical or ultraviolet luminosity (McLure \\& Dunlop 2001, Vestergaard 2002), with which it is correlated, at least over a limited range of luminosities (Kaspi et al. 2000).  A different approach to estimating black hole mass is to exploit the correlation, seen in nearby normal galaxies, between black hole mass and stellar velocity dispersion, $\\sigma$ (Ferrarese \\& Merritt 2000, Gebhardt et al. 2000a). If AGN host galaxies are similar to non-active galaxies, this correlation should hold also for them. Since stellar velocity dispersion measurements are still difficult for higher redshift AGN, the stellar velocity dispersion can possibly be inferred from effective radius and central surface brightness assuming AGN host galaxies occupy the same fundamental plane as ordinary ellipticals (O'Dowd et al. 2002).  Some previous studies have found a tight relation between mass and luminosity in AGN (Dibai 1981, Wandel \\& Yahil 1985, Padovani \\& Rafanelli 1988, Koratkar \\& Gaskell 1991, Kaspi et al. 2000); however, the scatter is large when the black hole masses are restricted to the most reliable estimates (from reverberation mapping). One might have expected a correlation between AGN black hole mass and luminosity since the Eddington luminosity is proportional to black hole mass, but if there is a range in accretion rates and/or efficiencies, the relation will be weaker.  In this paper, we collect and compare all AGN black hole mass estimates from the literature, and we make new black hole mass estimates from stellar velocity dispersions (\\S~2). We calculate bolometric luminosities for these same AGN to investigate their mass--luminosity relation, and look for trends of Eddington ratio with luminosity (\\S~3). Table~\\ref{T_SUM} summarizes the number of black hole mass estimates from the various methods. We use $H_0=75$~km s$^{-1}$ and $q_0=0.5$ throughout this paper. ",
        "conclusions": "We estimated and/or collected from the literature black hole masses for 377 AGN, obtained with various methods. These span a range of nearly 4 orders of magnitude, from $10^6~M_\\odot$ to $7 \\times 10^{9}~M_\\odot$. Direct comparisons suggest that reverberation mapping and stellar velocity dispersion give reliable black hole mass estimates --- within factors of a few --- while using optical luminosity to infer broad-line size or using the fundamental plane to infer velocity dispersion leads to somewhat larger uncertainties. In the case of virial estimates (reverberation mapping, optical luminosity, or other), additional uncertainties enter through the unknown orbits and the possible non-virial motions of the line-emitting gas.  We estimated bolometric luminosities for most of the AGN, apart from those affected strongly by beaming or by obscuration of the nuclear emission. Comparing bolometric luminosity to black hole mass for 234 AGN, we find little or no correlation. Gaps in coverage of the $L_{bol}$--$M_{BH}$ plane are due at least in part to high-mass, low-luminosity objects like the BL Lac objects and radio galaxies for which we have no good bolometric luminosity estimates.  For a given black hole mass, bolometric luminosities range over as many as four orders of magnitude. The Eddington ratios span nearly as large a range, 2--3 orders of magnitude at most luminosities. These are much larger than any uncertainties in the estimates of either black hole mass or luminosity. There are no strong trends of Eddington ratio with luminosity, contrary to long-held preconceptions. The absence of low Eddington ratios at high redshifts (high luminosities) can be explained at least in part by selection effects in flux-limited surveys wherein highly sub-Eddington AGN disappear progressively at higher redshifts.  We also do not confirm previously reported trends of radio luminosity with black hole mass, and while our results indicate a modest dependence of radio loudness on black hole mass, selection effects may exaggerate or even produce this trend. On the whole, black hole mass seems to have remarkably little to do with the appearance of active nuclei, either their luminosities or radio power.  Of course, the present sample includes a randomly selected mix of AGN, with black hole masses estimated in different ways, by different people, from different data sets. There may be real trends dependent on other variables not taken into account here (e.g., AGN type). It is obviously of interest to apply the more robust black hole mass estimation methods --- reverberation mapping and stellar velocity dispersion --- to a large sample of AGN, at as high a redshift as possible, although these methods will probably not work for the typical AGN at $z\\sim 2$--3. In practice, such a study would start with measurements of stellar velocity dispersions at $0.05 \\lesssim z \\lesssim 0.4$, which require 4- to 10-m class telescopes."
    },
    "0207/astro-ph0207280_arXiv.txt": {
        "abstract": "s{In the past five years observations with the {\\em Rossi X-ray Timing Explorer} have revealed fast quasi-periodic oscillations in the X-ray flux of about 20 X-ray binaries. Thought to originate close to the surface of a neutron star, these oscillations provide unique information about the strong gravitational field in which they are produced.}   The dynamical timescale close to an object of mass $M$ and radius $R$ is $\\tau \\sim R / v_{\\rm ff} = \\sqrt{ R^3 / 2 G M}$, where $v_{\\rm ff}$ is the local free-fall velocity. For typical masses and radii of neutron stars ($M \\sim 1-2$\\,$\\Msun$ and $R \\sim 10-20$\\,km), $\\tau \\simless 10^{-3}$\\,s. It was only after NASA's {\\em Rossi X-ray Timing Explorer (RXTE)} was launched in December 1995 that we could finally probe these short timescales, both in neutron-star and black-hole X-ray binaries. Groups at Goddard Space Flight Center and the University of Amsterdam, almost simultaneously announced the detection of submillisecond quasi-periodic variability in the X-ray flux of two binary systems with neutron star primary\\cite{iauc6319,iauc6320}. Since then, so-called kilohertz quasi-periodic oscillations (kHz QPOs) have been observed in more than 20 X-ray binary systems\\cite{vdk00}.   KHz QPOs usually appear in pairs\\cite{ford97a}, with frequencies $\\nu_{1}$ and $\\nu_{2} (> \\nu_{1})$ between $\\sim 300$ Hz\\cite{jonker00} (the lowest value for $\\nu_{1}$) and $\\sim 1330$ Hz\\cite{straaten00} (the highest value for $\\nu_{2}$). For a given source both $\\nu_{1}$ and $\\nu_{2}$ can drift by $\\sim 100$ Hz on timescales of a few hundred seconds\\cite{strohmayer96}, with a typical dynamic range of $\\sim 200 - 300$ Hz; up to $\\sim 800$ Hz variations have been observed in one source\\cite{straaten00}. Despite these rather large $\\nu_{1}, \\nu_{2}$ shifts, for each source $\\Delta \\nu = \\nu_{2} - \\nu_{1}$, remains more or less constant\\cite{ford97a,strohmayer96,wijnands98a}. In the best studied sources, $\\Delta \\nu$ has been observed to decrease by $50-100$ Hz as $\\nu_{1}, \\nu_{2}$ increase\\cite{vanderklis97,mendez98c}. In all the other sources of kHz QPO, $\\Delta \\nu$ is consistent with this trend\\cite{psaltis98}  It is widely accepted that at least one of the two QPO peaks is produced by matter in Keplerian orbits around the neutron star at the inner edge of the accretion disc$^{12-14}$. If this is so, QPO frequencies should depend sensitively upon the mass and angular momentum of the compact star as well as upon the orbital radius at which the QPOs are produced\\cite{mlp98a,sv99}. In principle it is possible to measure some of these parameters independently$^{23-25}$, which would allow us to set simultaneous constraints to the mass and radius of the star, and would provide means of knowing the high-density equation of state of neutronic matter\\cite{mlp98a}, an issue of central importance in nuclear physics research.  On timescales of a few hours QPO frequencies correlate with X-ray intensity\\cite{ford97a}, but this correlation breaks down on timescales of a day or more\\cite{zhang98a}; on longer timescales QPO frequencies correlate much better with the X-ray spectral properties of the source$^{6,10-16}$. It is thought that mass accretion rate is the mechanism behind this behavior\\cite{mlp98a}.  The Fourier power spectra of X-ray binaries often show a broad-band component that is roughly flat below a break frequency at $\\nu_{\\rm b} \\sim 1 - 10$ Hz and decreases above $\\nu_{\\rm b}$, and a low-frequency QPO at $\\sim 10 - 50$ Hz. Studies with RXTE have recently shown that both $\\nu_{\\rm b}$ and the frequency of the low-frequency QPO are strongly correlated to the frequency of the kHz QPOs$^{16,26-28}$ and to each other\\cite{wk99}, and that these correlation encompasses systems both with neutron star and black hole primaries. This result is quite significant, as it seems to indicate that the presence of a solid surface, a magnetic field, or an event horizon plays no significant role in the mechanism that produces the rapid variability observed in these objects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207555_arXiv.txt": {
        "abstract": "The Millennium Galaxy Catalogue (MGC) is a $37.5$~deg$^2$, medium-deep, $B$-band imaging survey along the celestial equator, taken with the Wide Field Camera on the Isaac Newton Telescope. The survey region is contained within the regions of both the Two Degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey Early Data Release (SDSS-EDR). The survey has a uniform isophotal detection limit of $26$\\mpass\\ and it provides a robust, well-defined catalogue of stars and galaxies in the range $16 \\le \\bmgc < 24$~mag.  Here we describe the survey strategy, the photometric and astrometric calibration, source detection and analysis, and present the galaxy number counts that connect the bright and faint galaxy populations within a single survey. We argue that these counts represent the state of the art and use them to constrain the normalizations ($\\phi^*$) of a number of recent estimates of the local galaxy luminosity function. We find that the 2dFGRS, SDSS Commissioning Data (CD), ESO Slice Project, Century Survey, Durham/UKST, Mt Stromlo/APM, SSRS2, and NOG luminosity functions require a revision of their published $\\phi^*$ values by factors of $1.05\\pm0.05$, $0.76\\pm0.10$, $1.02\\pm0.22$, $1.02\\pm0.16$, $1.16\\pm0.28$, $1.75\\pm0.37$, $1.40\\pm0.26$ and $1.01\\pm0.39$, respectively. After renormalizing the galaxy luminosity functions we find a mean local $\\bj$ luminosity density of $\\overline{j_{\\bj}} = (1.986 \\pm 0.031) \\times 10^8 \\; h \\; L_{\\odot}$~Mpc$^{-3}$.\\footnotemark ",
        "introduction": "\\label{introduction} Our understanding of the local universe and the local galaxy population originates primarily from the all-sky photographic Schmidt surveys and the established catalogues of bright galaxies derived from them, such as the Catalogue of Galaxies and Clusters of Galaxies \\citep{Zwicky68}, the Morphological Catalogue of Galaxies \\citep{Vorontsov74}, the Uppsala General Catalogue of Galaxies \\citep{Nilson73}, the ESO/Uppsala Catalogue \\citep{Lauberts82,Lauberts89}, the Southern Galaxy Catalogue \\citep*{Corwin85}, the Catalogue of Principal Galaxies \\citep{Paturel89}, the Edinburgh/Durham Southern Galaxy Catalogue \\citep*{Heydon89}, the APM catalogue \\citep{Maddox90}, the Third Reference Catalogue of Bright Galaxies \\citep{deVaucouleurs91} and the SuperCOSMOS Sky Survey \\citep{Hambly01}. While these catalogues have provided invaluable information and insight, uncertainty remains as to their completeness, particularly for low surface brightness and compact galaxies \\citep{Disney76,Sprayberry97,Impey97,Drinkwater99}. In addition there are concerns as to the photometric accuracy \\citep*[e.g.][]{Metcalfe95b}, the susceptibility to scale errors \\citep{Bertin97}, plate-to-plate variations \\citep{Cross03} and dynamic range.  These photographic-based catalogues have been the starting point for numerous spectroscopic surveys aimed at measuring the local space density of galaxies (i.e.\\ the local galaxy luminosity function). The space density of galaxies is our fundamental census of the local contents of space and therefore a crucial constraint for models of galaxy formation \\citep[e.g.][]{White91,Cole00,Pearce01}. If the imaging catalogues are in omission and/or photometrically inaccurate then regardless of the completeness of the spectroscopic surveys our insight into the galaxy population will be incomplete and most likely biased against specific galaxy types.  Over the past two decades there have been numerous estimates of the local galaxy luminosity function (e.g.\\ EEP, \\citealt*{Efstathiou88}; Mt Stromlo/APM, \\citealt{Loveday92}; Autofib, \\citealt{Ellis96}; ESP, \\citealt{Zucca97}; SSRS2, \\citealt{Marzke98}; Durham/UKST, \\citealt{Ratcliffe98}; SDSS-CD, \\citealt{Blanton01}; 2dFGRS, \\citealt{Norberg02}) and of the three-parameter Schechter function used to represent it \\citep{Schechter76}. Typically the surveys agree broadly on the faint end slope ($\\alpha$, $\\Delta \\alpha \\approx \\pm 0.15$) but show a marked variation in the characteristic luminosity ($L^*$, $\\Delta L^* \\approx 40$ per cent) and normalization ($\\phi^*$, $\\Delta \\phi^* \\approx 50$ per cent). The uncertainties in the Schechter parameters result in an uncertainty of $> 60$ per cent in the local luminosity density, $j = \\phi^*  L^*  \\Gamma(\\alpha+2)$.  This uncertainty is usually expressed as the normalization problem which has been somewhat overshadowed by the more notorious faint blue galaxy problem \\citep{Koo92,Ellis97}. The latter describes the inability of basic galaxy number count models to predict the numbers of galaxies seen at faint magnitudes ($22 < B < 28$~mag) in the deep pencil beam CCD-based surveys \\citep[e.g.][]{Tyson88,Metcalfe95,Metcalfe01}. The lesser known normalization problem describes the inability of number count models to explain the galaxy counts even at bright magnitudes ($18 < B < 20$~mag) by as much as a factor of $2$ \\citep*[see discussions in][]{Shanks84,Driver95,Marzke98,Cohen03}. In many ways the normalization problem is the more fundamental: while luminosity evolution, cosmology and/or dwarf galaxies can be, and have been, invoked in varying mixtures to explain the faint blue galaxy problem \\citep*[e.g.][]{Broadhurst88,Babul92,Phillipps95,Ferguson98}, none of these can be used to resolve the normalization problem.  \\begin{figure} \\psfig{file=fig1.ps,width=\\columnwidth} \\caption{The magnitude ranges and survey areas spanned by some previous number count publications. Surveys based on photographic material are shown with a dashed line. The vertical lines show various transition regions where various effects start to dominate the galaxy counts. Key: SSS \\citep{Hambly01}, APM \\citep{Maddox90b}, SDSS-EDR \\citep{Yasuda01}, MAMA \\citep{Bertin97}, EDSGC \\citep{Heydon89}, MGC (this work), G96 \\citep{Gardner96}, J91 \\citep{Jones91}, EIS \\citep{Prandoni99}, KW01 \\citep{Kummel01}, I86 \\citep*{Infante86}, K86 \\citep{Koo86}, A97 \\citep{Arnouts97}, CADIS \\citep{Huang01}, M91 \\citep{Metcalfe91}, T88 \\citep{Tyson88}, WHDF \\citep{Metcalfe01}, M95 \\citep{Metcalfe95}, NTTDF \\citep{Arnouts99}, HDF \\citep{Williams96}.} \\label{surveys} \\end{figure}  In the past the problem was typically circumvented by renormalizing the number count models to the range $18 < B < 20$~mag \\citep*[e.g.][]{Driver94,Metcalfe95,Driver95,Driver98,Marzke98,Metcalfe01}. The justification was that the bright galaxy catalogues, on which the luminosity function measurements are based, are shallow and therefore susceptible to local clustering. However the crucial normalization range typically occurs at the faint limit of the photographic surveys (where the photometry and completeness are more likely to be a problem) and at the bright end of the pencil beam CCD surveys (where statistics are poor). While convenient, the clustering explanation overlooks two more worrisome possibilities: gross photometric errors and/or gross incompleteness in the local catalogues. If either of these two latter explanations play a part this will have important consequences for the new-generation spectroscopic surveys, namely the 2dFGRS \\citep{Colless01} and the SDSS \\citep{York00}. The input catalogue of the 2dFGRS is an extensively revised version of the photographic APM survey (which is known to show a peculiar steepening in its galaxy counts at bright magnitudes, \\citealt{Maddox90b}), with zero-point and scale-error corrections from a variety of sources including the 2MASS $K$-band survey and the data presented in this paper (see \\citealp{Norberg02} for details). In the case of the SDSS -- which leaps forward in terms of dynamic range, uniformity and wavelength coverage -- the effective exposure time is relatively short ($54$~s) and the isophotal detection limit is comparable to that of the photographic surveys. Hence while issues of photometric accuracy should be resolved the question mark of completeness may remain.  \\begin{figure} \\centerline{\\psfig{file=fig2.ps,angle=-90,width=\\columnwidth}} \\caption{Outline of MGC fields 1, 2 and 3 and the arrangement of the CCDs.} \\label{outline} \\end{figure}  To address the above problems within a single, well-defined dataset we require a survey that is reasonably deep and yet has a large enough solid angle to provide accurate and statistically significant galaxy counts over the crucial normalization range. Furthermore, the survey's photometry must be accurate and its completeness high, i.e.\\ it must probe to low surface brightnesses. Fig.~\\ref{surveys} shows a number of imaging surveys in terms of their sky coverage and magnitude range. Dashed and solid lines indicate photographic and CCD-based surveys, respectively. Typically the faint surveys are CCD-based while the local surveys are photographic (with the notable and recent exception of the SDSS-EDR, \\citealt{Stoughton02}). While the CCD surveys make significant improvements in surface brightness and magnitude limits their sky coverage is small. It is only very recently that large CCD mosaics such as the Wide Field Camera \\citep[WFC,][]{Irwin01} and the SDSS instrument \\citep{Gunn98} have been constructed that now allow a large area of sky to be surveyed within a realistic time frame.  In this paper we present the Millennium Galaxy Catalogue (MGC, Sections \\ref{data}--\\ref{sellimits}). The MGC represents a new medium-deep, wide-angle galaxy resource, which firmly connects the local and distant universe within a single dataset (cf.\\ Fig~\\ref{surveys}). In Section \\ref{counts} we produce the galaxy number counts spanning the range $16 \\le \\bmgc < 24$~mag. We then focus on the normalization problem by comparing our counts over the range $16 \\le \\bmgc < 20$~mag to the predictions of a number of local luminosity function estimates in Section \\ref{lumfs}. Our counts provide stringent constraints on the normalization of the luminosity function and hence on the local luminosity density. Our conclusions are given in Section \\ref{conclusions}.  The 2dFGRS and SDSS will essentially supersede all previous redshift surveys and therefore it is important to verify their photometric accuracy and completeness on as large a scale as possible. We will provide a detailed comparison of the 2dFGRS and SDSS-EDR imaging catalogues with the MGC in a companion paper \\citep{Cross03}.  \\begin{figure} \\psfig{file=fig3.ps,width=\\columnwidth,angle=-90} \\caption{Summary of the data quality across the MGC survey strip. The dots in the uppermost panel indicate the location of the photometric calibration fields. The sky and sky noise parameters were calculated from the measured mode and rms of the background pixel value distribution.} \\label{obsstat} \\end{figure}  A more long-term aim of the MGC project is to provide structural information on the galaxy population around the crucial normalization point ($16 < B < 20$~mag). It is ironic that since the advent of the Hubble Space Telescope we have a greater understanding of the morphological mix of galaxies at faint magnitudes than at bright magnitudes. For example, \\citet{Driver98} and \\citet{Cohen03} published morphological galaxy counts spanning the range $21 < B_{\\rm F450W} < 26$~mag, yet no reliable morphological galaxy counts at brighter magnitudes exist. Consequently, no accurate local morphological luminosity functions exist (compare, for example, the conflicting results of \\citealt{Loveday92} and \\citealt{Marzke98}) and the evolution of the different morphological types cannot be accurately constrained. The MGC will enable us to remedy this situation as it allows morphological classification and the extraction of structural parameters to $\\bmgc = 20$~mag.  The data and catalogues presented in this paper are publically available at http://www.roe.ac.uk/$\\sim$jol/mgc/. ",
        "conclusions": "\\label{conclusions} Here we have presented a detailed description of the Millennium Galaxy Catalogue (MGC), a deep ($\\mu_{\\rm lim} = 26$\\mpass) wide ($37.5$~deg$^2$) survey along the equatorial strip from $9^{\\rm h} 58^{\\rm m}$ to $14^{\\rm h} 47^{\\rm m}$. We have demonstrated that the internal photometric accuracy of the MGC is $\\pm 0.023$~mag and that the astrometric accuracy is $\\pm 0.08$~arcsec in both RA and Dec. Using {\\sc SExtractor} we have derived a source catalogue containing over 1 million objects spanning the range $16 \\le \\bmgc < 24$~mag. All non-stellar detections brighter than $\\bmgc = 20$~mag have been visually inspected and the objects repaired where necessary. We have taken care to exclude objects from regions where the photometry is likely to be erroneous, resulting in a robust and clean estimation of the galaxy number counts over the range $16 \\le \\bmgc < 24$~mag. These data finally connect the faint pencil beam CCD surveys of the past decade to the local Universe. The selection boundaries of the MGC are well defined and to $\\bmgc = 20$~mag we are demonstrably robust to star--galaxy separation and low- and high-surface brightness concerns. We contest that the MGC galaxy number counts in this range are the state of the art, superseding all previous intermediate number count data.  We use the counts to test various estimates of the galaxy luminosity function and find that many of them predict counts where the normalizations are inconsistent with our observations. In Fig.~\\ref{phistar} we present the best-fitting value of $\\phi^*$ as a function of $M^*$ and $\\alpha$. In Table~\\ref{lumftab} we list the appropriate $\\phi^*$ values for a number of popular $B$-band luminosity function estimates.  From these revised values we constrain the $\\bj$-band luminosity density of the local Universe for each of these luminosity functions. We find $\\overline{j_{\\bj}} = (1.986\\pm 0.031) \\times 10^8 \\; h \\; L_{\\odot}$~Mpc$^{-3}$. The 2dFGRS and SDSS-EDR consistently give a slightly higher value of $j_{\\bj} = (2.035 \\pm 0.046) \\times 10^8 \\; h \\; L_{\\odot}$~Mpc$^{-3}$."
    },
    "0207/astro-ph0207233_arXiv.txt": {
        "abstract": "{Simple analytical models, such as the Hernquist model, are very useful tools to investigate the dynamical structure of galaxies. Unfortunately, most of the analytical distribution functions are either isotropic or of the Osipkov-Merritt type, and hence basically one-dimensional. We present three different families of anisotropic distribution functions that self-consistently generate the Hernquist potential-density pair. These families have constant, increasing and decreasing anisotropy profiles respectively, and can hence represent a wide variety of orbital structures. For all of the models presented, the distribution function and the velocity dispersions can be written in terms of elementary functions. These models are ideal tools for a wide range of applications, in particular to generate the initial conditions for $N$-body or Monte Carlo simulations. ",
        "introduction": "From a stellar dynamical point of view, the most complete description of a stellar system is the distribution function $F(\\bfr,\\bfv)$, which gives the probability density for the stars in phase space. In this paper, we will concentrate on the problem of constructing anisotropic equilibrium distribution functions that self-consistently generate a given spherical mass density profile $\\rho(r)$. In the assumption of spherical symmetry, the mass density of stellar system can easily be derived from the observed surface brightness profile, at least if we assume that the mass-to-light ratio is constant and that dust attenuation is negligible. And as the surface brightness of a galaxy (or bulge or cluster) is fairly cheap and straightforward to observe, compared to other dynamical observables which require expensive spectroscopy, the problem we will deal with is relevant and important.  The first step in the construction of self-consistent models is the calculation of the gravitational potential $\\psi(r)$, which can immediately be determined through Poisson's equation. The second step, the actual construction of the distribution function, is less straightforward. Basic stellar dynamics theory (see e.g.\\ Binney \\& Tremaine 1987) learns that steady-state distribution functions for spherical systems can generally be written as a function of binding energy and angular momentum. We hence have to determine a distribution function $F({\\cal{E}},L)$, such that the zeroth order moment of this distribution function equals the density, i.e.\\ we have to solve the integral equation \\begin{equation} \\rho(r) = \\iiint F({\\cal{E}},L)\\,\\txd\\bfv \\label{detdf} \\end{equation} for $F({\\cal{E}},L)$. Hereby we have to take into account that not every function $F({\\cal{E}},L)$ that satisfies this equation is a physically acceptable solution: an acceptable solution has to be non-negative over the entire phase space. In general, the problem of solving the integral equation (\\ref{detdf}) is a degenerate problem, because there are in general infinitely many distribution functions possible for a given potential-density pair.  Particularly interesting are models for which the distribution function and its moments can be computed analytically. Such models have many useful applications, which can roughly be divided into two classes. On the one hand, they can improve our general understanding of physical processes in galaxies in an elegant way. For example, they can serve as simple galaxy models, in which it is easy to generate the starting conditions for $N$-body or Monte Carlo simulations, or to test new data reduction or dynamical modelling techniques. A quick look at the overwhelming success of simple analytical models, such as the Plummer sphere (Plummer 1911; Dejonghe 1987), the isochrone sphere (H\\'enon 1959, 1960), the Jaffe model (Jaffe 1983) and the Hernquist model (Hernquist 1990), provides enough evidence. On the other hand, analytical models are also useful for the detailed dynamical modelling of galaxies. For example, in modelling techniques such as the QP technique (Dejonghe 1989), a dynamical model for an observed galaxy is built up as a linear combination of components, for each of which the distribution function and its moments are known analytically. As a result, the distribution function and the moments of the final model are also analytical, which obviously has a number of advantages.  Unfortunately, the number of dynamical models for which the distribution function is known analytically is rather modest. Moreover, most of them consist of distribution functions that are isotropic or of the Osipkov-Merritt type, and therefore basically one-dimensional. An exception is the completely analytical family of anisotropic models described by Dejonghe (1987). These models self-consistently generate the Plummer potential-density pair, a simple yet useful model for systems with a constant density core.  During the last decade, however, it has become clear that, at small radii, elliptical galaxies usually have central density profiles that behave as $r^{-\\gamma}$ with $0\\leq\\gamma\\leq2.5$ (Lauer et al.\\ 1995; Gebhardt et al. 1996). Such galaxies can obviously not be adequately modelled with a constant density core. This has stimulated the quest for simple potential-density pairs, and corresponding distribution functions, with a central density cusp. The first effort to construct such models was undertaken by Ciotti (1991) and Ciotti \\& Lanzoni (1997), who discussed the the dynamical structure of stellar systems following the $R^{1/m}$ law (S\\'ersic 1968), a natural generalization of the empirical $R^{1/4}$ law of de Vaucouleurs (1948). A major drawback of this family, however, is that the spatial density and the distribution function can not be written in terms of elementary functions (see Mazure \\& Capelato 2002). A more useful family is formed by the so-called $\\gamma$-models (Dehnen 1993; Tremaine et al.\\ 1994), characterized by a density proportional to $r^{-4}$ at large radii and a divergence in the center as $r^{-\\gamma}$ with $0\\leq\\gamma\\leq3$. The dynamical structure of models with this potential-density pair has been extensively investigated (e.g.\\ Carollo, de Zeeuw \\& van der Marel 1995; Ciotti 1996; Meza \\& Zamorano 1997), but only for isotropic or Osipkov-Merritt type distribution functions. Simple analytical models with a more general anisotropy structure are still lacking.  In this paper we construct a number of families of completely analytical anisotropic dynamical models that self-consistently generate the Hernquist (1990) potential-density pair. It is a special case of the family of $\\gamma$-models, corresponding to $\\gamma=1$. In dimensionless units, the Hernquist potential-density pair is given by \\begin{subequations} \\begin{gather} \\psi(r) = \\frac{1}{1+r} \\label{hernpot} \\\\ \\rho(r) = \\frac{1}{2\\pi}\\,\\frac{1}{r(1+r)^3}. \\label{hernrho} \\end{gather} \\end{subequations} As the density diverges as $1/r$ for $r\\rightarrow0$, the surface brightness $I(R)$ will diverge logarithmically for $R\\rightarrow0$. More precisely, the surface brightness profile has the form \\begin{equation} I(R) = \\frac{1}{2\\pi}\\, \\frac{(2+R^2)\\,X(R)-3}{(1-R^2)^2}, \\end{equation} with $X(R)$ a continuous function defined as \\begin{equation} X(R) = \\begin{cases} \\,\\,(1-R^2)^{-1/2}\\,\\arcsech R &\\qquad\\text{for $0\\leq R\\leq1$,} \\\\ \\,\\,(R^2-1)^{-1/2}\\,\\arcsecans R &\\qquad\\text{for $1\\leq R\\leq\\infty$.} \\\\ \\end{cases} \\label{defX} \\end{equation}  The paper is organized as follows. The general theory on the inversion the fundamental integral equation (\\ref{detdf}) is resumed in Section 2. Each of the subsequent sections is devoted to special cases of this inversion technique and the corresponding family of Hernquist models. Isotropic models are the most simple ones; Hernquist (1990) showed that, in the special case of isotropy, the distribution function and its moments can be calculated analytically. We repeat the most important characteristics of the isotropic Hernquist model in Section 3. In Section 4 we construct a one-parameter family of models with a constant anisotropy. In Section 5, a two-parameter family of Hernquist models is constructed by means of the Cuddeford (1991) inversion technique. These models have an arbitrary anisotropy in the center and are radially anisotropic at large radii. On the contrary, in Section 6, a two-parameter family is constructed that has a decreasing anisotropy profile, with arbitrary values for the anisotropy in the center and the outer halo. Finally, Section 7 sums up. ",
        "conclusions": "Three new families of anisotropic dynamical models have been presented that self-consistently generate the Hernquist potential-density pair. For all models, in particular for the Cuddeford models of Section~{\\ref{cudd.sec}}, we checked the conditions on the adopted parameters such that the distribution is positive, and hence physically acceptable, in phase space.  They host a wide variety of orbital structures: in general, the models presented can have an arbitrary central anisotropy, and a outer halo with the same anisotropy, a purely radial orbital structure, or an arbitrary, but more tangential, anisotropy. In order to produce models that have an arbitrary anisotropy in the central regions, and a more radial, but not purely radial, anisotropy at large radii, the most cost-effective way seems to construct a linear combination of a number of 'component' dynamical models, such as the ones presented here. This technique has been adopted for several years in the QP formalism (Dejonghe 1989, for an overview see Dejonghe et al. 2001), where most of the components in the program libraries have an intrinsically tangential orbital structure.  For all of the presented models, we have analytical expressions for the distribution function and the velocity dispersions in terms of elementary functions. They are hence ideal tools for a wide range of applications, for example to generate the initial conditions for $N$-body or Monte Carlo simulations. At this point, a number of remarks are appropriate.  First, very few elliptical galaxies are perfectly spherical; actually, various observational and theoretical evidence suggests that many elliptical galaxies are at least moderately triaxial (Dubinski \\& Carlberg 1991; Hernquist 1993; Tremblay \\& Merritt 1995; Bak \\& Statler 2000). Unfortunately, an extension of the presented techniques to construct analytical axisymmetric or triaxial systems is not obvious, because the internal dynamics of such stellar systems is much more complicated than in the spherical case. Nevertheless, our models can be used as a onset to construct numerical axisymmetric of triaxial distribution functions with different internal dynamical structures, for example by the adiabatic squeezing technique presented by Holley-Bockelmann et al. (2001).  Second, the models presented here are self-consistent models, whereas it is nowadays believed that most elliptical galaxies contain dark matter, either in the form of a central black hole (Merritt \\& Ferrarese 2001 and references therein) and/or a dark halo (Kronawitter et al.\\ 2000; Magorrian \\& Ballantyne 2001). When constructing dynamical models with dark matter, an extra component must be added to the gravitational potential. For example, Ciotti (1996) constructed analytical two-component models in which both the stellar and dark matter components have a Hernquist density profile and an Osipkov-Merritt type distribution function. The models presented in this paper can also be extended to contain a dark halo or a central black hole. Indeed, the adopted inversion techniques are perfectly suitable for this, because the augmented density functions $\\tilde{\\rho}(\\psi,r)$ do not necessarily need to satisfy the self-consistency condition (\\ref{condrho}). Adding an extra term to the potential does not conceptually change the character of the inversion, but it might complicate the mathematical exercise.  Third, we have not discussed stability issues for the presented models. The study of the stability of anisotropic stellar systems is difficult, and a satisfactory criterion can not easily be given. For stability against radial perturbations, we can apply the sufficient criterions of Antonov (1962) or Dor\\'emus \\& Feix (1973), but numerical simulations have shown that these criteria are rather crude (Dejonghe \\& Merritt 1988; Meza \\& Zamorano 1997). Moreover, the only instability that is thought to be effective in realistic galaxies is the so-called radial orbit instability, an instability that drives galaxies with a large number of radial orbits to forming a bar (H\\'enon 1973; Palmer \\& Papaloizou 1987; Cincotta, Nunez \\& Muzzio 1996). The behavior of galaxy models against perturbations of this kind can only be tested with detailed $N$-body simulations or numerical linear stability analysis. Meza \\& Zamorano (1997) used $N$-body simulations to investigate the radial orbit instability for a number of spherical models of the Osipkov-Merritt type, including the Hernquist model. They found that the models are unstable for $r_a\\lesssim1$, which significantly restricts the set of models that correspond to positive distribution functions (see Table \\ref{lambdamax.tab}). It would be interesting to extend this investigation to the three families of Hernquist models presented in this paper, but this falls beyond the scope of this paper."
    },
    "0207/astro-ph0207005_arXiv.txt": {
        "abstract": "Recent {\\it XMM--Newton} observations of the high-redshift, lensed, broad absorption line (BAL) quasi-stellar object APM 08279+5255, one of the most luminous objects in the universe, allowed the detection of a high column density absorber ($N_H \\approx 10^{23}$~cm$^{-2}$) in the form of a K-shell absorption edge of significantly ionized iron (Fe XV - XVIII) and corresponding ionized lower--energy absorption. Our findings confirm a basic prediction of phenomenological geometry models for the BAL outflow and can constrain the size of the absorbing region. The Fe/O abundance of the absorbing material is significantly higher than solar (Fe/O = 2--5), giving interesting constraints on the gas enrichment history in the early Universe. ",
        "introduction": "Absorption Line (BAL) Quasi-stellar Objects (QSOs) are a key to understand the geometry and physical state of the medium in the immediate vicinity of accreting supermassive black holes. A new unified model (Elvis 2000)  indicates that a significant fraction of the matter accreted into the region of the compact object is flowing out again. On either side of the accretion disk it should form a funnel- shaped shell, in which the outflowing gas is ionized and accelerated to velocities of 0.05-0.1$c$ by the powerful radiation force of the central object. If the observer's inclination is favorable, this flow intercepts the line of sight with a large column density and produces the blue-shifted broad UV absorption line features observed in about 10\\% of all QSOs. However, the UV/optical spectra sample only a minor fraction of the total column density of the flow, which is predicted to be highly ionized so that it mainly absorbs X--rays.   APM 08279+5255 is an exceptionally luminous Broad Absorption Line (BAL) QSO at redshift z=3.91 (Irwin et al., 1998), coincident with an IRAS Faint Source Catalogue object and was also detected at 850 $\\mu$m with SCUBA, implying an apparent far-infrared luminosity of $> 5 \\times 10^{15}~L_{\\odot}$ (Lewis et al., 1998). The object is strongly lensed (Ledoux et al., 1998; Ibata et al., 1999; Egami et al., 2000), with a magnification factor of 50-100, but even taking this magnification factor into account, the object is still among the most luminous in the Universe. The low resolution discovery spectrum of APM 08279+5255 showed a broad absorption trough (BAL) on the blue side of Ly$_\\alpha$. An excellent high resolution optical spectrum was obtained with HIRES at the Keck telescope (Ellison et al., 1999) and a detailed study of the physical conditions in the broad absorption line flow of the QSO (Srianand \\& Petitjean, 2000) came to the conclusion, that the corresponding gas stream, outflowing with velocities up to 12000 km~s$^{-1}$, is heavily structured and highly ionized.  In this paper we report X--ray observations of APM 08279+5255 with {\\it XMM--Newton} obtained in October 2001 and April 2002, where we detected an ionized Fe K edge in the continuum of the QSO which is very likely related to the UV broad absorption line flow. In section 2 we present the X--ray observations and analysis. In section 3 we discuss the results. Throughout this paper we use a Hubble constant of  50~km~s$^{-1}$~Mpc$^{-1}$ and a deceleration parameter $q_0=0.5$. ",
        "conclusions": "For the first time, we have reported the detection of a highly ionized iron edge in the spectrum of a high-redshift BAL quasar. This indicates a very high column density, ionized absorber in the line of sight. Compared to the warm absorbers typically observed in Seyfert galaxies, which are dominated by OVII -- OVIII edges of optical depth around unity (see Komossa 1999 and references therein), where strong Fe K edges have not been seen so far, the column density of ionized Fe is considerably higher here and the absorption is very likely associated with the highly ionized BAL flow observed in the UV spectrum of the source.  For the quasar 3C 351 it has been shown, that the associated absorption lines observed in the UV spectrum are most likely produced by the same highly ionized material that is also responsible for the moderate X--ray warm absorber (Mathur et al., 1994). In case of BAL systems, however, interpretation is complicated by the uncertainties in optical depth measurements of the lines: the broad, saturated UV absorption lines from different atomic species do not allow an unambiguous column density determination (e.g., Hamann 1998).   \\subsection{Implications for X--ray Broad Absorption Line Models}  Murray \\& Chiang (1995) constructed a model in which the BALs arise in an accretion-disk wind driven by line pressure. They predicted that objects with very broad CIV absorption ($>$ 5000 km~s$^{-1}$) should also show Fe absorption edges in the X-ray regime. This is similar to what we now observe.  The high column density inferred for the X-ray absorber in APM 08279+5255 is also consistent with predictions of the unified model by Elvis (2000). In that model, the UV broad absorption lines occur, when the line of sight to the central object grazes along the funnel--shaped shell of ionized matter thought to constitute the BAL outflow. Using NGC\\,5548 as reference object in constructing his new unified model,  Elvis (2000) predicted  a column density of a few 10$^{23}$ cm$^{-2}$ for an NGC\\,5548-like object viewed along the funnel. We find $\\sim$10$^{23}$ cm$^{-2}$ for APM 08279+5255.  The ionization state of the warm absorber is characterized by the ionization parameter $U$. The number rate of ionizing photons, $Q$,  was estimated from our piecewise powerlaw spectrum with $\\alpha_{\\rm uv-x}$=--1.4 in the EUV and $\\Gamma_{\\rm x}$ as observed, normalized to the observed X--ray flux. We then find $Q = 1.2 \\times 10^{58} k^{-1}$ s$^{-1}$, $k$ being the lensing magnification factor. Given the best estimate of the ionization parameter (log U = 0.4) and assuming a density of the absorber of log $n$ = 9.5 -- this is similar to the value suggested by Elvis (2000) to ensure pressure equilibrium with the BELR clouds -- we expect a lower limit on the location of the absorber of $r \\approx 2 \\times 10^{18} k^{-1/2}$ cm. A more detailed discussion of the ionized gas properties and consequences for BAL models will be given in a forthcoming paper (G. Hasinger et al. 2002, in prep.).   \\subsection{Relation of X--ray and UV Absorbers and Metal Abundances}  Comparison with UV data immediately implies that the X--ray absorber must be dust-free, else the UV continuum would be heavily extincted. The high degree of ionization of the X--ray absorber is consistent with the UV spectrum which shows strong high-ionization BAL lines including OVI, but relatively weak neutral H absorption.  The UV absorption line data show that the BAL flow is highly structured with saturated absorption bands and thus, contrary to the X--ray absorption edge, cannot determine the total column density in the flow. Based on our best-fit {\\em{Cloudy}} model, we expect to see UV absorption in e.g.,  HI, CIV, NV, OVI and NeVIII with column densities on the order of $N_{\\rm HI} \\approx 4\\,\\,10^{16}$cm$^{-2}$, $N_{\\rm CIV} \\approx 10^{16}$cm$^{-2}$, $N_{\\rm NV} \\approx 7\\,\\,10^{17}$cm$^{-2}$, $N_{\\rm OVI} \\approx 6\\,\\,10^{17}$cm$^{-2}$, and $N_{\\rm NeVIII} \\approx 7\\,\\,10^{17}$cm$^{-2}$. The higher metal column densities compared to HI can be traced back to the high degree of ionization of the absorber and do therefore not necessarily indicate chemical overabundances. For a detailed comparison with the UV data, {\\em simultaneous} multi-wavelength observations are required.  No diagnostically valuable, strong Fe lines are present in the UV spectra of BALs. Thus, iron abundance determinations for BALs are still scarce. Fe-K edge X-ray observations therefore nicely supplement UV measurements of other metal species, and allow tests of different scenarios for the origin of the BAL gas.  We find an overabundance of iron of 2--5 times the solar abundance (see Fig. 3). Due to the fact that the measured X--ray absorption is produced mainly by O, Ne, Mg, Si and Fe, this iron overabundance is actually an estimate of $Fe/O$.  The outflowing material must therefore have already been processed in a starburst environment. Detailed chemical evolutionary scenarios of the emission-line gas in quasars (Hamann \\& Ferland 1993) predict the iron enrichment that depends mostly on the lifetime of supernova Type Ia precursors, leading to an expected delay of $\\sim$1 Gyr until $Fe/O$ reaches solar values. Fe measurements in high-$z$ objects, like APM 08279+5255, therefore (1) are of profound relevance for understanding the early star formation history of the universe and (2) provide important constraints on cosmological models. Assuming an Fe abundance of APM 08279+5255 of at least solar, we can place severe constraints on some of the enrichment models of Hamann \\& +Ferland (1993; their Fig. 1). Furthermore, at the redshift of $z\\approx 4$ the age of the universe is a little less than 1 Gyr (for $q_{\\rm o}$=0.5), which compares to the timescale of $\\sim$1 Gyr necessary to enrich $Fe/O$ up to the solar value. Given that we find strong indications of a supersolar $Fe/O$ abundance, we are beginning to constrain cosmological models, favoring those which predict larger galaxy ages at a given $z$.  In the near future, Fe abundance measurements of larger samples of high-$z$ quasars may provide another valuable path to measure cosmological parameters (Hamann \\& Ferland 1993)."
    },
    "0207/astro-ph0207143_arXiv.txt": {
        "abstract": "We show that the second knee in the cosmic ray spectrum (i.e. the steepening occurring at $E\\simeq 4\\times 10^{17}$~eV) could be related to drift effects affecting the heaviest nuclear component, the iron group nuclei, in a scenario in which the knee at $3\\times  10^{15}$~eV indicates the onset of drift effects in the lighter proton component. We also study the anisotropies resulting from diffusion and drift currents in the Galaxy, showing their potential relevance to account for the AGASA observations at $E\\sim 10^{18}$~eV, before the extragalactic component becomes dominant. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207469_arXiv.txt": {
        "abstract": "{ Rotation curves for four spiral galaxies with recently determined Cepheid-based distances are reconsidered in terms of modified Newtonian dynamics (MOND).  For two of the objects, NGC 2403 and NGC 7331, the rotation curves predicted by MOND are compatible with the observed curves when these galaxies are taken to be at the Cepheid distance. For NGC 3198, the largest distance for which reasonable agreement is obtained is 10\\% smaller than the Cepheid-based distance; i.e., MOND clearly prefers a smaller distance.  This conclusion is unaltered when new near-infrared photometry of NGC 3198 is taken as the tracer of the stellar mass distribution.  For the large Sc spiral, NGC 2841, MOND requires a distance which is at least 20\\% larger than the Cepheid-based distance.  However, the discrepancy of the Tully-Fisher and SNIa distances with the Cepheid determination casts some doubt upon the Cepheid method in this case. ",
        "introduction": "It is well established that, in the context of Newtonian dynamics, the observable mass in spiral galaxies cannot account for the observed flat rotation curves in the outer regions of galaxies (Bosma 1978; Begeman 1987, van Albada et al. 1985).  The standard explanation for this discrepancy is the proposal that galaxies are embedded in an extended dark halo which dominates the gravitational field in the outer regions (Trimble 1987).  An alternative explanation for the discrepancy is the possibility that dynamics becomes non-Newtonian in the limit of low accelerations. The most successful such proposal is Milgrom's (1983) modified Newtonian dynamics or MOND. Here the idea is that below a certain acceleration threshold ($a_o$) the effective gravitational acceleration approaches $\\sqrt{a_og_n}$ where $g_n$ is the usual Newtonian acceleration.  This modification yields asymptotically flat rotation curves of spiral galaxies and a luminosity -- rotation velocity relationship of the observed form, $L \\propto v^4$, the Tully-Fisher relation (Tully \\& Fisher 1977). But apart from these general aspects the prescription also successfully predicts the observed form of galaxy rotation curves from the observed distribution of stars and gas with reasonable values for the mass-to-light ratio of the stellar component (Begeman et al. 1991; Sanders 1996; Sanders \\& Verheijen 1998, McGaugh \\& de Blok 1998). A crucial element of a very specific prescription like MOND is that the precise form of the rotation curve is predicted by the observed mass distribution given the value of a single universal parameter; in this case, the critical acceleration $a_0$. Consequently MOND can in principle be falsified as soon as there is one galaxy for which the predicted rotation curve disagrees significantly with the observed curve; although, in practice, the usual uncertainties inherent in astronomical data render a definitive falsification problematic in any individual case.  In Begeman et al. (1991, hereafter BBS) MOND is applied to a sample of galaxies for which high quality \\hi rotation curves are available. For a value of $a_0$ equal to 1.21~10$^{-8}$ cm~s$^{-2}$ the rotation curves of the sample could be reasonably reproduced, the free parameter in each case being the mass-to-light ratio of the visible disc.  Because MOND is an acceleration dependent modification, this derived value of $a_o$ depends upon assumed distance scale ($H_o = 75$ \\kmss Mpc$^{-1}$ in this case).  Moreover, the quality of an individual fit depends upon the adopted distance to the galaxy, and, since the relative distances to these nearby galaxies have not been known to within an accuracy, typically, of 25\\%, this has provided some freedom to adjust the distance in order to improve the MOND fit; i.e., distance, within certain limits, can be considered as an additional second parameter in the fitting procedure.  For most of the galaxies in the sample of BBS, the distance did not have to be adjusted significantly ($<10\\%$) to improve the MOND fits, and the improvement was not significant.  However, one object, NGC 2841, required a large readjustment:  The Hubble law distance to this galaxy is about 9 Mpc, but MOND clearly prefers a distance which is twice as large.  Using ground-based and Hubble Space Telescope observations Cepheid distances to 21 inclined galaxies have now been determined as part of the HST key program on the extragalactic distance scale (e.g. Sakai et al. 2000). Three of the galaxies in this Cepheid sample are also in the sample with high quality rotation curves considered by BBS. These are NGC 2403 (Freedman \\& Madore 1988), NGC 3198 (Kelson et al. 1999) and NGC 7331 (Hughes et al. 1998).  For these three galaxies the MOND prescription can now be considered in the context of the Cepheid distance that is generally considered to be the most precise indicator.  NGC 2841 has been discussed as a critical case for MOND by Sanders (1996). For this galaxy, there is also a large discrepancy between the Tully-Fisher distance and the Hubble law distance (for plausible values of the Hubble constant). Moreover, the galaxy was the site of a recent SNIa (1999by). For these reasons this galaxy has been included, subsequently, in the HST program (Macri et al. 2001).  Here we demonstrate that for two galaxies in the BBS sample the rotation curve predicted by MOND is consistent with the observed curve when the galaxies are placed at the Cepheid distance.  However, for NGC 3198 at the Cepheid distance of 13.8 $\\pm$ 0.6 Mpc, the shape of the rotation curve predicted by MOND systematically deviates (by up to 10 \\kms) from the observed curve, both in the inner and outer regions.  The largest distance which can be compatible with MOND is about 10\\% lower than the Cepheid-based distance. This is not particularly problematic because of likely uncertainties in the Cepheid method and in the determination of a rotation curve from the observed two-dimensional velocity field. NGC 2841, however, remains a difficult case for MOND.  The minimum distance which is consistent with MOND is about 17 Mpc whereas the Cepheid-based distance is 14.1 $\\pm$ 1.5 Mpc.  We discuss the implications and seriousness of this discrepancy for MOND, or, alternatively, for the Cepheid method. ",
        "conclusions": "The main conclusions of this paper can be summarized as follows:  \\begin{enumerate} \\item For the galaxies NGC 2403 and NGC 7331, MOND rotation curves agree acceptably with the observed curves when these galaxies are taken at the Cepheid-based distance. \\item For NGC 3198 at the Cepheid-based distance of 13.8 Mpc, the MOND curve shows small ($<10$ \\kms) but significant systematic deviations from the observed curve. \\item If the distance to NGC 3198 is taken to be 12.5 Mpc, or 10\\% less than the Cepheid-based distance, the MOND curve is an acceptable representation of the observed curve.  This lower distance is probably within the uncertainties of the Cepheid method. \\item These conclusions are unaltered by utilizing recent near-infrared photometry of NGC 3198 which does show evidence for a small central bulge and bar component. \\item For NGC 2841, the rotation curve predicted by MOND when the galaxy is taken to be at the upper limit on the Cepheid-based distance (15.6 Mpc) remains inconsistent with the observed curve, with systematic deviations of more than 30 \\kms. \\item The smallest distance for which the MOND curve is compatible with the observed curve (given the uncertainties involved in the tilted ring technique for modelling warps), is 17 Mpc or 20\\% larger than the Cepheid-based distance.  The preferred MOND distance is 23 Mpc. \\item The TF distance to NGC 2841, based upon the Cepheid-re-calibrated TF relation is 24 Mpc.  If the distance to this galaxy is really 14.1 Mpc, then it would deviate from the mean I-band TF relation by more than 1 magnitude. \\item NGC 2841 has been the host of a type Ia supernova, 1999by.  If this galaxy is at the Cepheid-based distance, this would be one of the least intrinsically luminous supernovae ever observed.  Calibrating the peak luminosity by the Phillips relation, the SN-based distance is 23.5 Mpc. \\end{enumerate}  It is clear that NGC 2841 remains a critical case for MOND. The discrepancy between Cepheid-based distance and both the TF and SNIa based distances to NGC 2841 suggests that there may be a problem with the derived Cepheid-based distance.  In general, it is evident that accurate distance determinations to nearby galaxies are extremely relevant to the question of the viability of MOND. MOND, as a modification of Newtonian dynamics attached to an acceleration scale, is far more fragile than the dark matter hypothesis in this regard.  It would be useful to obtain more Cepheid-based distances to the sample of galaxies with well-observed rotation curves.  Particularly useful would be a Cepheid distance estimate to the Ursa Major cluster as many of these galaxies have well-measured rotation curves and near-infrared photometry."
    },
    "0207/astro-ph0207157_arXiv.txt": {
        "abstract": "We report the $Beppo$SAX observations of 6 Flat Spectrum Radio Quasars. Three of them have a clear detection up to 100 keV with the PDS instrument. For 4 objects the X-ray spectrum is satisfactorily fitted by a power-law continuum with Galactic absorption. 2251+158 show the presence of absorption higher than the galactic value, while the spectrum of the source 0208-512 shows a complex structure, with evidence of absorption at low energy. We construct the Spectral Energy Distributions adding historical data to the broad band X-ray spectra obtained with $Beppo$SAX and reproduce them with a one-zone Synchrotron-Inverse Compton model (including both SSC and External Compton). The implications are briefly discussed. ",
        "introduction": "Blazars are the best laboratory to study the physics of relativistic jets. The non-thermal continuum amplified by relativistic beaming (e.g., Urry \\& Padovani 1995), is a unique tool to probe the physical processes acting in the jet.  In the last decade the interest for Blazars has been renewed by the EGRET discovery of 66 Blazars (Hartman et al. 1999) as strong $\\gamma $-ray emitters. In several cases the $\\gamma $-ray emission dominates the apparent bolometric luminosity. Moreover the short variability timescales, together with the absence of absorption of high energy photons (through the $\\gamma \\gamma$-pair production), directly imply that the $\\gamma $-ray emission is also relativistically beamed (e.g. Dondi \\& Ghisellini 1995).  Despite the large variety of classifications, there is growing evidence that Blazars form a single population. Their different spectral properties can be unified within a spectral sequence in which the leading parameter is the total luminosity (Fossati et al. 1998). At the high-luminosity extreme of the sequence we find the Flat Spectrum Radio Quasars, the most luminous blazars, characterized by the presence of bright emission lines in the optical-UV spectrum and, in some cases, strong {\\it UV bumps}, signatures of the underlying accretion process (e.g. Pian et al. 1999 and references therein). Most of the Blazars detected by EGRET (Hartman et al. 1999) belong to this sub-class but not all FSRQs with comparable SEDs have been detected in the $\\gamma $-ray domain.  In the study of FSRQs, observations in the X-ray band play a fundamental role. The X-ray emission from FSRQ is dominated by the continuum originating from Inverse Compton scattering between relativistic electrons in the jet and the soft photons produced by the disk and/or in the Broad Line Region (Dermer \\& Schlickeiser 1993; Sikora, Begelman \\& Rees 1994). The possibility to measure the IC component yelds important constraints on the emission models. Moreover the study of the soft X-ray spectrum (below 1 keV) provides insight on the presence of intrinsic absorption, on the possible contribution of the high energy end of the synchrotron emission or on the minimum energy of emitting particles. The latter point is particularly interesting because the determination of the minimum energy of particles is a fundamental step in the estimate of the total power transported by the jet (e.g. Celotti, Padovani \\& Ghisellini 1997).  The {\\it Beppo}SAX satellite has the unique capability of covering the wide energy range $0.1- 200$ keV. For this reason it is ideal to study FSRQs. In particular the knowledge of the spectrum in this band, together with the $\\gamma $-ray spectrum measured by EGRET, provide the most complete spectral information on the high energy component (see e.g., Tavecchio et al. 2000, hereafter Paper I).  For these reasons we started an observational program with {\\it Beppo}SAX of a sample of FSRQs, (containing a total of 50 sources) extracted from the 2-Jy sample by Padovani \\& Urry (1992). Using the flux threshold $F_{\\rm 1\\, keV}>0.5 \\mu$Jy the number of sources reduces to 19. In Paper I we discussed the case of three FSRQ detected by EGRET, namely 0836+710, 1510-089 and 2230+114. Here we report the analysis of $Beppo$SAX data of 6 other sources, 3 detected by EGRET and three without an EGRET detection. Adding these 6 sources to the 3 FSRQs analyzed in Paper I and other 3 sources (3C279, 0528+134, PKS 0537-441) discussed elsewhere (Maraschi et al. 1998; Ghisellini et al. 1999, Pian et al., in prep) a total of 12 FSRQs (out of 19 with X-ray flux $F_{\\rm 1\\, keV}>0.5 \\mu$Jy) from the 2 Jy sample were observed by $Beppo$SAX. The sources discussed here are listed in Table 1. Partial and preliminary results were reported in Maraschi \\& Tavecchio (2001)  The paper is organized as follows: in section 2 we report the analysis of {\\it Beppo}SAX observations, in section 3 we discuss the emission models for the SEDs and finally in Sect. 4 we discuss our results. ",
        "conclusions": "We have presented the results of the analysis of 6 FSRQs observed with $Beppo$SAX. The smooth power-law ($\\Gamma =1.6-1.7$) X-ray continuum is produced through Inverse Compton scattering on photon external to the jet, except for the ``atypical'' source 0521-365. In the case of two sources, 0208-512 and 2251+158, there is evidence of more complex X-ray spectra. In particular 0208-512 shows an absorption feature at $\\sim 0.6$ keV, that we suggest could be explained allowing the presence of a warm absorption intercepting the emission produced in the jet.  The soft X-ray spectrum of 2251+158 shows a clear deficit of photons, consistent either with an intrinsic absorption (with rest-frame column density $\\sim 5\\times 10^{21}$ cm$^{-2}$) or an intrinsic break in the continuum occurring below $\\sim $1 keV. In the latter case the position of the break could provide important insights into the physics of the jet, in particular regarding the value of the minimum energy of the radiating particles. On the other hand there is growing evidence that the presence of intrinsic absorption is common in high-redshift radio-loud quasars (e.g., Reeves \\& Turner 2000). Indeed $Beppo$SAX observations of two distant ($z\\sim 4$) blazars (Fabian et al. 2001a,b) show the presence of huge absorption, with column density of the order of few $10^{23}$ cm$^{-2}$.  For three sources (0208-512, 1641+399 and 2251+158) we detect hard X-ray emission up to $\\sim 100$ keV. For 1641+399 and 2251+158 the hard X-ray component appears to be the smooth extrapolation of the soft-medium X-ray continuum, while in 0208-512 the PDS spectrum is likely contaminated by another source present in the large FOV of the instrument. We note however that 0208-512 is one of the few blazars detected by COMPTEL (Blom et al. 1995) and because of its exceptionally intense flux in the MeV band is considered as a prototype of the so-called ``MeV Blazars''. Therefore we can not completely exclude the possibility that the spectral component responsible for the MeV flux contributes also in the PDS band, explaining the observed excess.  The modelling of the SEDs shows that FSRQs (except the case 0521-365, see below) are characterized by jets with Doppler beaming factors larger than 10 and magnetic field of the order of $B\\sim 2-3$ G, roughly in equipartition with the emitting electrons. The size of the emitting region is of the order of $10^{16}$ cm, consistent with the absence of short-timescale variability ($\\lta 1$ day) in this class of blazars. These results confirm, with a larger sample, the conclusions of Paper I  An interesting issue for Blazars is the value of the minimum Lorentz factor of the emitting electrons, $\\gamma _{\\rm min}$. In Paper I we showed how, from the absence of a spectral break in the soft X-ray continuum, it is possible to infer that the electron energy distribution extends down to $\\gamma _{\\rm min} \\sim 1$. In fact a break in the EC continuum is expected at the frequency $\\nu _{\\rm br}\\simeq \\Gamma ^2 \\nu _{\\rm ext} \\gamma _{\\rm min} \\sim 10^{17} \\Gamma _{1}\\sim $ (e.g. Sikora et al. 1997), corresponding to an energy $\\sim 1$ keV (in the rest frame of the source). We recall that the lack of soft X-ray photons observed in 2251+158 (as for the case of 0836+710 reported in Paper I) could be interpreted as due to such intrinsic break in the continuum. Observations with larger signal to noise (possible with $XMM-Newton$ or $Chandra$) will allow to disentangle the two possible interpretations, intrinsic curved continuum or absorption.  In our fits $n_1$, the index of the low energy portion of the electron energy distribution, is smaller than 2. Clearly such a flat spectrum can not be produced by the cooling of high energy electrons (in that case ine would expect $n_1=2$). On the other hand this spectrum is even flatter than the standard prediction $n_1=2$ for a population of non-thermal electrons produced through Fermi acceleration by a non-relativistic shock (e.g. Blandford \\& Ostriker 1978). However recent calculations (see the review by Kirk \\& Dendy 2001) show that relativistic shocks, such as those likely present in the jet of blazars, could produce extremely flat ($n_1< 2$) spectra.  The case of 0521-365 is atypical. Previous observational evidence suggested that the jet forms a relatively large angle with the line of sight. Pian et al. (1996), using independent arguments (in particular the presence of an optical jet,e.g. Scarpa et al. 1999), concluded that $\\theta \\sim 30^{\\rm o }$. Assuming that the bulk Lorentz factor of the emitting plasma is similar to that of the other blazars, $\\Gamma _{\\rm b} \\sim 10$, this implies that the emission from 0521-365 is weakly boosted ($\\delta =1-2$). The SED reported in Fig.5 has been calculated assuming $\\Gamma _{\\rm b} =10$ and $\\theta =15^{\\rm o }$, implying $\\delta=3$. In this situation the X-ray/$\\gamma$-ray continuum, contrary to the case of the other sources, is dominated by the SSC emission, while the EC spectrum gives only a small contribution in the $\\gamma $-ray band.  Part of the motivation for this investigation was to study the difference between the EGRET and non-EGRET sources. In the whole FSRQ sample observed with \\sax there are 9 EGRET sources and 3 sources not detected by EGRET. The average X-ray spectral indices ($\\Gamma _{\\rm EGRET}=1.56 \\pm 0.06$ and $\\Gamma _{\\rm non-EGRET}= 1.69\\pm 0.07$ for EGRET and no-EGRET sources, respectively) indicate that, within the statistical uncertainty, the X-ray spectral characteristics are similar for both groups. The modelling of the SEDS confirms the strong similarity between the two groups.  The lack of the EGRET detection could be attributed to the variability behaviour typical of these sources in the $\\gamma $-ray band (e.g. Mukherjee et al. 1997) rather than to an intrinsic difference in the properties of the jet. New $\\gamma $-ray missions (AGILE, GLAST) will help to further investigate these problems."
    },
    "0207/astro-ph0207361_arXiv.txt": {
        "abstract": "The search for gamma radiation in clusters of galaxies represents a precious tool to investigate the history of these large scale structures. Clusters or sources within them accelerate cosmic rays, as demonstrated by the detection of radio halos, hard X-rays and UV emission, and confine them over cosmological time scales. Nonthermal and thermal phenomena may be closely related and observations of gamma rays may tell us about this link. In this paper we review the physics of cosmic ray acceleration and confinement in clusters of galaxies and the related gamma ray signatures. In particular we describe in some detail the role of cluster mergers for the acceleration of nonthermal particles. The perspectives for gamma ray detection with GLAST and with ground based detectors are also discussed. ",
        "introduction": "The presence of nonthermal particles in clusters of galaxies is a well established fact. These particles are responsible for extended synchrotron radio halos in several clusters (see Feretti et al., 2000 for a recent rewiew), as well as for hard X-ray (HXR) and extreme ultraviolet (EUV) excesses (see e.g. Fusco Femiano et al., 1999, 2000; Lieu et al., 1996). Several explanations have been proposed for the origin of this radiation, but at present there is no conclusive evidence in favor or against any of these models.  The simplest explanation for the HXR excess is based on the inverse compton scattering (ICS) of the same relativistic electrons that are responsible for the radio halos. In this case, low values of the intracluster magnetic field are required, that seem in some contradiction with the much larger values of the fields evaluated through faraday rotation measurements (RM) (Eilek 1999; Clarke, Kronberg \\& B\\\"{o}ringer 1999). These measurements are however quite difficult, and the discrepancy needs to be considered critically. The severe energy losses associated with the relativistic electrons make the nonthermal phenomena due to synchrotron and ICS of relativistic electrons transient phenomena, with duration not longer than a few billion years. However, if the acceleration processes occur during some violent phenomenon such as cluster mergers, some level of reacceleration may be expected due to the presence of turbulence in the intracluster medium (Brunetti et al. 2001a, 2001b).  Alternative explanations of the HXR excess have also been proposed, based on acceleration of electrons from the thermal bath and bremsstrahlung radiation from these particles (Ensslin, Lieu \\& Biermann 1999; Blasi 2000; Dogiel 2000; Sarazin \\& Kempner 2000). These models also have their drawbacks, as discussed by Petrosian (these proceedings) and Petrosian (2001).  An important theoretical insight transformed our way of looking at clusters: cosmic rays accelerated in clusters are trapped there for cosmological times (Berezinsky, Blasi, \\& Ptuskin 1997; V\\\"{o}lk, Aharonian, \\& Breitschwerdt 1996). Clusters behave as cosmological storage rooms for cosmic rays. The combination of this argument and the ever-increasing mass of observations of nonthermal phenomena, generated an unprecedented interest in clusters as possible sources of gamma rays. The detection (or not) of gamma radiation by one of the future gamma ray telescopes such as GLAST, or even current ground based telescopes such as STACEE or HEGRA would allow us to weigh the nonthermal content of clusters and achieve a better understanding of the nonthermal history of these large scale structures.  The issue of nonthermal radiation is clearly related to the problem of acceleration of  particles: although the common wisdom is that the acceleration occurs during mergers of subclusters, there are several arguments which complicate this simple picture. We discuss this important issue at length in this review.  Throughout the paper we assume a flat cosmology ($\\Omega_0=1$) with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=1-\\Omega_m=0.7$ and a value for the Hubble constant of $70\\, \\rm{km/s/Mpc}$.  The paper is planned as follows: in \\S 2 we discuss the physics of cosmic ray confinement in clusters of galaxies; in \\S 3 we summarize the gamma ray predictions for gamma rays from clusters of galaxies. \\S 4 is devoted to the investigation of merger shocks as cosmic ray accelerators. The consequences of gamma ray production from clusters onto the diffuse gamma ray background are discussed in \\S 5, while our conclusions are presented in \\S 6. ",
        "conclusions": "The diffuse medium in the intracluster volume is filled with a nonthermal gas of particles that are the relics of all the events occurred within the cluster itself. Gamma ray astronomy is an important tool to study this component and infer information about particle acceleration and confinement and about the specific processes (mergers, active galactic phases, and many others) that inject nonthermal particles in a cluster and possibly contribute to its heating.  We discussed here the expectations for gamma ray fluxes in the presence of both protons and primary electrons in clusters. Even if the amount of hadronic cosmic rays trapped on cosmological scales is, say, $10\\%$ of the equipartition energy, we expect that the gamma ray fluxes may be detectable by GLAST for energies above 100 MeV, and by future ground based gamma ray telescopes such as VERITAS, MAGIC and HESS at higher energies. Even current observations with STACEE and HEGRA could actually provide interesting information about the abundance of cosmic rays in the intracluster gas of nearby clusters, as shown in fig. 1. Unfortunately, we are not aware of any scientific report of such attempt to look for high energy gamma rays with either STACEE or HEGRA. While the higher energy gamma ray emission is likely to be generated by hadronic interactions, the lower energy gamma rays (in the MeV-GeV range) can be generated by several processes related to electrons, and most of the fluxes derived in the literature are in the range of interest for GLAST.  The importance of gamma ray observations can be appreciated particularly well in the context of the growing multifrequency observations, that one piece at a time, are allowing us to understand the processes that occur in the intracluster volume, enriching it with hot gas, nonthermal particles and magnetic fields, in a way that at present is still unclear."
    },
    "0207/astro-ph0207011_arXiv.txt": {
        "abstract": "We observed a group of galaxies, HCG 57, with ASCA. Regardless that their member galaxies are dominated by spiral galaxies, we detected extended thermal X-ray emission that is attributed to hot gas with a temperature of $1.04\\pm0.10$ keV. This is the second clear detection of thermal X-ray emission from a spiral-dominant group of galaxies after HCG 92. The luminosity of the thermal emission is about $5\\times10^{41}$ erg s$^{-1}$ in the 0.5--10 keV band, which is higher than that of HCG 92, but relatively less luminous among groups of galaxies. The X-ray emission is extended over several member galaxies, and is thus associated with the group rather than an individual galaxy. The metal abundance cannot be well constrained with a lower limit of 0.08 solar. The gas-to-stellar mass ratio is $\\sim0.3$. Although this is relatively low among groups, the hot gas is also a significant component even in the spiral-dominant group. We suggest that the X-ray faintness of spiral-dominant groups is due to the low surface brightness and somewhat low gas mass, at least in the case of HCG 57. ",
        "introduction": "Groups of galaxies are the poorest systems of clusters of galaxies. Since the number of member galaxies is quite small, individual galaxies could have great effects on the group properties. Many groups of galaxies have been found to contain a significant amount of X-ray emitting hot gas with a temperature of $\\sim$1 keV (Ponman, Bertram 1993; Mulchaey et al. 1996). ROSAT observations revealed that the X-ray luminosity and mass of hot gas scatter widely with a range of 2 orders of magnitude, although their galaxy mass is similar within a factor of $\\sim$5 (Mulchaey et al. 1996, 2000). Ponman et al. (1996) claimed that this scatter is caused by a steep relation between the temperature and the X-ray luminosity, indicating the action of galactic winds to the hot gas.  The metal abundance is also different with a range of 0.05--1 solar among groups (Fukazawa et al. 1996; Davis et al. 1999), in contrast to a small scatter of 0.2--0.4 solar for rich clusters of galaxies. They speculated that the shallow potential cannot bind metal-rich gas ejected from member galaxies. This scenario is also supported by the low Si-to-Fe abundance ratio of groups, as compared with rich clusters (Fukazawa et al. 1998). However, we have not yet understood the cause of significant scatters in metal abundances and in the hot gas-to-stellar mass ratio. In order to resolve this issue, further X-ray investigations of groups of galaxies are needed. It has been reported that some correlations exist among the physical parameters of groups, such as the galaxy velocity dispersion, optical luminosity, X-ray luminosity, X-ray temperature, and so on (Ponman et al. 1996; Helsdon, Ponman 2000; Helsdon et al. 2001).  In the past observations, the temperature and metal abundance have been measured mainly for groups of galaxies whose member galaxies are dominated by elliptical ones. This is due to the X-ray faintness of ones dominated by spiral galaxies (Mulchaey et al. 1996). Why are spiral-dominant groups of galaxies X-ray faint? How much hot gas is there in such groups? How high is the metal abundance of hot gas in such groups? These are important questions to understand concerning how member galaxies have a relation with a hot intragroup medium. In the past, only one spiral-dominant group of galaxies has shown reasonably bright intergalactic X-ray emission. This is HCG 92, famous as the Stephan's Quintet, which contains a significant amount of hot intragroup medium (Sulentic et al. 1995; Awaki et al. 1997). The X-ray luminosity is $1\\times10^{41}$ erg s$^{-1}$, and the gas-to-stellar mass ratio is $\\sim$0.1. There is another possible evidence of a hot intragroup medium in the spiral-dominant group HCG 16 (Ponman et al. 1996), although its X-ray properties are not well known. Since the sample is still too poor for the purpose described above, we planned to perform ASCA (Tanaka et al. 1994) observations of a spiral-dominant group of galaxies.  HCG 57 is one of the Hickson Compact galaxy groups (Hickson et al. 1988; Hickson 1993) at a redshift of 0.0304. It contains eight original member galaxies in the Hickson Catalog, as tabulated in tables 1 and 2. For a comparison, an X-ray bright elliptical-dominant galaxy group, HCG 51 (Fukazawa et al. 1996), is also shown. It can be seen that these two groups exhibit quite similar properties, such as redshift, galaxy velocity dispersion, galaxy density, optical luminosity, galaxy number, and group size. The most different feature is a morphology of dominant member galaxies. The brightest three galaxies of HCG 57 are spiral, while those of HCG 51 are elliptical or S0. Therefore, a comparison of the X-ray properties between HCG 57 and HCG 51 is valuable when considering the relation between the member galaxy morphology and the hot gas in groups. Ponman et al. (1996) reported, based on a short serendipitous ROSAT pointed observation, that HCG 57 is an X-ray emitter with a luminosity of $10^{41.98}$ erg s$^{-1}$ and a temperature of $0.8\\pm0.2$ keV. Imaging information is, however, not available, possibly due to the poor data quality; therefore, the existence of a hot intragroup medium is not known. We are especially interested in how much X-ray hot gas exists in HCG 57, because HCG 51 was found to contain a large amount ot hot gas whose mass is comparable to the total stellar mass in member galaxies. Note that the galaxy velocity dispersion of HCG 57 is high enough to bind hot gas gravitationally. This paper reports on the ASCA detection of X-ray emitting hot gas from HCG 57. This is the second detailed X-ray analysis of such a spiral-dominant galaxy group after HCG 92. We employed 90\\% confidence limits throughout this work, and used a Hubble constant of 50 km s$^{-1}$ Mpc$^{-1}$. The solar abundances are taken from Anders and Grevesse (1989). ",
        "conclusions": "We observed a spiral-dominant group of galaxies, HCG 57, with ASCA, and detected extended thermal X-ray emission with a temperature of 1 keV around HCG 57, together with the hard X-ray component, whose extent cannot be explained by one point source. Soft thermal emission was not associated with the specific galaxy, but extended up to 5 arcmin ($\\sim$250 kpc). The luminosity and mass of the X-ray emitting thermal hot gas are $4.2\\times10^{41}$ erg s$^{-1}$ (0.5--10 keV) and $4.0\\times10^{11} M_{\\odot}$, respectively, within 5 arcmin of the group center. Compared with the temperature and X-ray luminosity relation of groups of galaxies taken from Xue and Wu (2000) in figure 4, the X-ray luminosity of HCG 57 is relatively low for a temperature of 1 keV, but still within the scatter. The ratios of the gas mass to stellar and total mass are 0.3 and 0.068, respectively, both of which are somewhat smaller than that of rich clusters and X-ray bright groups. This gives us a precious detailed X-ray result for the spiral-dominant groups of galaxies after HCG 92, and the third significant detection of X-ray emitting hot gas from such groups.  According to the X-ray properties of the thermal hot gas such as extent, temperature, mass, and luminosity, we conclude that it is not an interstellar medium in a specific galaxy, but an intragroup medium bound gravitationally by the HCG 57 itself. The low mass and low luminosity of hot gas is consistent with groups whose spiral fraction is large (Mulchaey et al. 1996). Our results indicate that the X-ray low luminosity of spiral-dominant groups is not due to the lack of an intragroup medium for some systems. For HCG 57, the X-ray faintness is mostly attributed to quite a low central gas density of at most $\\sim3\\times10^{-4}$ cm$^{-3}$. This is also the same as in the case of HCG 16 (Ponman et al. 1996). We think that the high central gas density of HCG 92 is unusual for spiral-dominant groups because no other such objects are found. We claim that the amount of hot gas in HCG 57 is not extremely small compared with X-ray bright groups. Therefore, it can be said that the gravitational potential of HCG 57 is deep enough to bind a significant amount of hot gas. This property differs from that for individual galaxies; the amount of hot gas in individual spiral galaxies is much less than that in individual elliptical galaxies by two or three orders of magnitude. Then, why are spiral-dominant groups always X-ray faint even if they contain massive hot gas? A low gas density of $<10^{-3}$ cm$^{-3}$, as in the case of HCG 57, or a low gas temperature of $kT<0.5$ keV is thought to be attributed. It is speculated that the environment of a low gas density makes spiral galaxies free from ram-pressure stripping, while a high gas density of bright groups of galaxies and rich clusters strips off the disk component of spiral galaxies to convert the morphology into S0.  The temperature of 1 keV for hot gas in the HCG 57 is typical for groups, and sufficiently high to bind a large amount of hot gas. Such a deep gravitational potential has been mainly found for elliptical-dominant groups, and HCG 57 also shows a clear evidence of such a potential for spiral-dominant groups as HCG 92. This implies that the gravitational potential or dark matter associated with groups of galaxies does not depend on the type of member galaxies. Apart from the group-scale gravitational potential, a galaxy-scale potential must exist, inferred from the double-$\\beta$ structure of the X-ray surface brightness around the group central galaxy (Ikebe et al. 1996; Mulchaey, Zabludoff 1999). For HCG 57, the measurement of the gravitational potential associated with central spiral galaxies is quite important, because it has not ever been performed with other methods. However, we cannot resolve it because of poor spatial resolution; and observation by Chandra is hopefully awaited.  A significant amount of metals in the hot gas is also found in HCG 57. Assuming that the metal abundance is 0.2 solar, the iron mass becomes $2.1\\times10^8 M_{\\odot}$. Accordingly, the iron-mass-to-light ratio (IMLR) is $8.0\\times10^{-4}$, which is small by a factor of 10 compared with that of rich clusters of galaxies. There are three early-type (E/S0) galaxies in HCG 57, and their total optical luminosity is $7.8\\times10^{10}L_{\\odot}$. Then, an IMLR for only early type galaxies becomes $2.3\\times10^{-3}$. Considering the IMLR of $\\sim10^{-2}$ for rich clusters, it is possible that the total iron in the hot gas was supplied mainly by these early type member galaxies. Nevertheless, we suggested that spiral member galaxies also contribute to the iron enrichment of the hot gas because of their dominance in HCG 57. The metal abundance ratio is an essential key to resolve this issue, and hence XMM/Newton and Astro-E2/XRS/XIS observations with good photon statistics will enable us to measure it.  The hard X-ray component is also detected from HCG 57. We cannot determine whether this emission is point-like or extended. If extended, several possible mechanisms can be considered. One is inverse Compton scattering via cosmic microwave background (CMB) photons by high-energy electrons in the intragroup space, as suggested for several rich clusters (e.g. Fusco-Femiano et al. 1999) and HCG 62 group (Fukazawa et al. 2001; Nakazawa 2001). The similar X-ray luminosity of the hard component with that of HCG62 and HCG51 is also interesting when considering the origin; the luminosity of the hard component does not depend on the luminosity of X-ray hot gas. For example, if the hard component is due to inverse Compton scattering, high-energy electrons should exist. Such electrons might be accelerated commonly in groups of galaxies through galaxy--galaxy or galaxy--plasma interactions (Nakazawa 2001), and luminosity of the hard component is proportional to the amount of electrons not dependent on their distribution. On the other hand, the luminosity of the X-ray hot gas strongly depends on the gas distribution because it is a square measure of the hot gas density. Therefore, the luminosity ratio of the soft and hard component can differ from groups to groups, and spiral-dominant groups could exhibit a clear hard component without being hindered by the soft component. Anyway, it is important to measure the spatial distribution of the hard X-ray component in order to consider its origin. Astro-E2/XIS will enable us such measurements thanks to the large effective area and low background.  The authors thank T.J. Ponman for many helpful comments. The authors are also grateful to the ASCA team for their help in the spacecraft operation and calibration."
    },
    "0207/astro-ph0207227_arXiv.txt": {
        "abstract": "The r-process of nucleosynthesis requires a large neutron-to-seed nucleus abundance ratio.  This does not, however, require that there be an excess of neutrons over protons.  If the expansion of the matter is sufficiently rapid and the entropy per nucleon is sufficiently high, the nucleosynthesis enters a heavy-element synthesis regime heretofore unexplored. In this extreme regime, characterized by a persistent disequilibrium between free nucleons and the abundant alpha particles, heavy r-process nuclei can form even in matter with more protons than neutrons.  This observation bears on the issue of the site of the r-process, on the variability of abundance yields from r-process events, and on constraints on neutrino physics derived from nucleosynthesis.  It also clarifies the difference between nucleosynthesis in the early universe and that in less extreme stellar explosive environments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/hep-ph0207238_arXiv.txt": {
        "abstract": "\\noindent {\\bf Abstract} \\vspace{2.5mm}  We investigate the possibility to use the neutrinos coming from a future galactic supernova explosion to perform neutrino oscillation tomography of the Earth's core. We propose to use existing or planned detectors, resulting in an additional payoff. Provided that all of the discussed uncertainties can be reduced as expected, we find that the average matter densities of the Earth's inner and outer cores could be measured with a precision competitive with geophysics. However, since seismic waves are more sensitive to matter density jumps than average matter densities, neutrino physics would give partly complementary information.  \\vspace*{0.2cm} \\noindent {\\it PACS:} 14.60.Lm; 13.15.+g; 91.35.-x; 97.60.Bw\\\\ \\noindent {\\it Key words:} Neutrino oscillations; Supernova neutrinos; Neutrino tomography; Geophysics ",
        "introduction": "In order to obtain more information about the interior of the Earth, neutrino tomography has been considered as an alternative method to geophysics. There exist, in principle, two different such techniques, neutrino absorption tomography~\\cite{Volkova74,DeRujula83,Wilson84,Askar84,Borisov87,Nicolaidis91,Crawford95,Kuo95,Jain:1999kp} and neutrino oscillation tomography~\\cite{Ermilova:1988pw,Chechin:1991,Ohlsson:2001ck,Ohlsson:2001fy,Ioannisian:2002yj}. Neutrino absorption tomography, based on the absorption of neutrinos in matter, is in some sense similar to X-ray tomography and unfortunately faces several problems including the need of extremely high energetic neutrino sources, huge detectors, and the prerequisite of many baselines. Neutrino oscillation tomography uses the fact that neutrino oscillations are influenced by the presence of matter~\\cite{mikh85,mikh86,wolf78}. Neutrino oscillation tomography would, in principle, be possible with a single baseline, since interference effects provide additional information on the matter density profile. However, it requires quite precise knowledge about the neutrino oscillation parameters and stringent bounds on the contribution of non-oscillation physics, such as neutrino decay, CPT violation, non-standard interactions, sterile neutrinos, \\etc{} Supernovae as neutrino sources are especially interesting, since the neutrinos come in large numbers from a short burst, which could be used to obtain a snapshot of the Earth's interior. In addition, compared to solar neutrinos, their energy spectrum has a high-energy tail, which is more sensitive to Earth matter effects. The influence of Earth matter on supernova neutrinos has, for example, been studied in \\Refs~\\cite{Lunardini:2001pb,Takahashi:2001dc,Takahashi:2000it}.  We assume that technologically feasible detectors exist, such as Super-Kamiokande, SNO, Hyper-Kamiokande, and UNO, which are originally built for different purposes, but also capable to detect supernova neutrinos. We discuss the possibility to use the neutrinos coming from a future galactic supernova explosion to determine with the assumed detectors some of the measurable quantities describing the structure of the Earth's interior. We especially focus on the outer and inner core of the Earth, since they are much harder to access with conventional geophysical methods than the mantle of the Earth. ",
        "conclusions": "\\label{sec:S&C}  We have discussed the possibility to use the neutrinos from a future galactic supernova explosion to obtain additional information on the Earth's core. First, we have summarized geophysical aspects and unknowns of the Earth's core. Then, we have investigated core-collapse supernovae as potential neutrino sources for a snapshot of the Earth's interior. Next, we have discussed the neutrino propagation from the production to the detection in detail, where we have especially focused on Earth matter effects on the neutrino oscillations of the supernova neutrinos. We have also put these effects into the context of the supernova parameters, \\ie, temperatures and deviation from energy equipartition. Furthermore, we have stressed the importance of supernova neutrinos arriving at the surface of the Earth as mass eigenstates for this technique, which we have also supported by a discussion of decoherence of neutrino oscillations. We have shown that we need one detector on the surface of the Earth on the side towards the supernova, and another one in the shadow of the Earth's core. For the most likely scenario of not crossing the Earth's inner core, we have shown that the Earth's average core matter density could be determined up to 9~\\% with a Super-Kamiokande-like and 1.3~\\% with a Hyper-Kamiokande-like detector (all errors at the $2 \\sigma$ confidence level). In addition, for a less likely two-parameter measurement of the outer and inner core matter densities, Hyper-Kamiokande could verify the existence of the inner core at the $3 \\sigma$ confidence level and measure the outer core matter density with a precision of about 3.1~\\%, as well as the density jump between outer and inner core matter densities with a precision of about 59~\\%. The latter error is comparable to seismic wave geophysics, where, however, not the difference between the average matter densities, but the matter density jump at the outer-inner core boundary is measured. Thus, neutrino physics could provide complementary information to geophysics. However, the quoted numbers for the precisions depend on the supernova parameter values, indirectly determined by the on-the-surface measurement, and are in some cases better, in others worse. Thus, the actual precisions will not be known before the supernova goes off. In general, we find that the more muon and tau neutrinos are produced, the larger the temperature difference between the different flavors is, and the more different the degeneracy parameters for the different flavors are, the better the application works. Finally, we have discussed several uncertainties to these measurements and we have found that especially the determination of the total muon antineutrino energy of the supernova causes problems to our method. However, measuring not only electron antineutrinos, but also the other two flavors could completely eliminate the dependence on the supernova parameters. Furthermore, the leading solar neutrino parameters have to be known with sufficient precision, which is about 0.2~\\% for Hyper-Kamiokande-like measurements. In summary, supernova neutrino tomography could be a nice additional payoff of existing or planned detectors if all of the prerequisites can be met at the time when the next supernova explodes."
    },
    "0207/astro-ph0207541_arXiv.txt": {
        "abstract": "$\\eta$ Carinae is a very luminous and unstable evolved star. Outflowing material ejected during the star's giant eruption in 1843 surrounds it as a nebula which consists of an inner bipolar region (the {\\it Homunculus\\/}) and the {\\it Outer Ejecta\\/}. The outer ejecta is very filamentary and shaped irregularly. Kinematic analysis, however, shows a regular bi-directional expansion despite of the complex morphology. Radial velocities in the outer ejecta reach up to 2000 kms$^{-1}$ and give rise to X-ray emission first detected by ROSAT. We will present a detailed study of the outer ejecta based on HST images, high-resolution echelle spectra for kinematic studies, images from CHANDRA/ACIS and HST-STIS spectra. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207407_arXiv.txt": {
        "abstract": "{ The initial descriptions of the FITS format provided a simplified method for describing the physical coordinate values of the image pixels, but deliberately did not specify any of the detailed conventions required to convey the complexities of actual image coordinates.  Building on conventions in wide use within astronomy, this paper proposes general extensions to the original methods for describing the world coordinates of FITS data.  In subsequent papers, we apply these general conventions to the methods by which spherical coordinates may be projected onto a two-dimensional plane and to frequency/wavelength/velocity coordinates. ",
        "introduction": "\\label{s:intro}  The Flexible Image Transport System, or FITS format, was first described by Wells et al.~(\\cite{kn:WGH}).  This format is characterized by a fixed logical record length of 2880 bytes, and the use of an unlimited number of character-format ``header'' records with an 80-byte, keyword-equals-value substructure. The header is followed by the header-specified number of binary data records, which are optionally followed by extension records of the specified length, but, at that time, of unspecified format.  Since then, a number of authors have suggested various types of extensions (e.g.~Greisen \\&\\ Harten \\cite{kn:GH}; Grosb\\o l et al.~\\cite{kn:GHGW}; Harten et al.~\\cite{kn:HGGW}; Cotton et al.~\\cite{kn:CTP}).  Because of its great flexibility, the FITS format has been, and continues to be, very widely used in astronomy. In fact, the FITS tape format was recommended (resolution C1) for use by all observatories by Commission 5 at the 1982 meeting of the IAU at Patras (\\cite{kn:IAU}) and the General Assembly of the IAU adopted (resolution R11) the recommendations of its commissions, including the FITS resolution. A committee of the NASA/Science Office of Standards and Technology has codified the current state of FITS into a document which has been accepted as the official definition of the standard (Hanisch et al.~\\cite{kn:NOST}).  Wells et al.~(\\cite{kn:WGH}) anticipated the need to specify the physical, or world, coordinates to be attached to each pixel of an {\\it N\\/}-dimensional image.  By {\\em world coordinates}, we mean coordinates that serve to locate a measurement in some multi-dimensional parameter space.  They include, for example, a measurable quantity such as the frequency or wavelength associated with each point in a spectrum or, more abstractly, the longitude and latitude in a conventional spherical coordinate system which define a direction in space.  World coordinates may also include enumerations such as ``Stokes parameters'', which do not form an image axis in the normal sense since interpolation along such axes is not meaningful.  Wells et al.~(\\cite{kn:WGH}) viewed each axis of the image as having a coordinate type and a reference point for which the pixel coordinate, a coordinate value, and an increment were given.  Note that this reference point was not required to occur at integer pixel locations nor even to occur within the image.  An undefined ``rotation'' parameter was also provided for each axis. Since there are, in general, more coordinates to be attached to a pixel than there are ``real'' axes in the {\\it N\\/}-dimensional image, the convention of declaring axes with a single pixel was also established in both examples given by Wells et al.  The keywords defined were \\begin{center}\\begin{tabular}{ll} \\CRVAL{n} & coordinate value at reference point \\\\ \\CRPIX{n} & array location of the reference point in pixels\\\\ \\CDELT{n} & coordinate increment at reference point \\\\ \\CTYPE{n} & axis type (8 characters) \\\\ \\CROTA{n} & rotation from stated coordinate type \\end{tabular}\\end{center} \\noindent A list of suggested values for \\CTYPE{n}\\ was provided with few of the details actually required to specify coordinates.  The units were specified to be The International System of Units ``SI'' (meters, kilograms, seconds) with the addition of degrees for angles.  The simplicity of this initial description was deliberate.  It was felt that a detailed specification of coordinate types was a lengthy and complicated business, well beyond the scope intended for the initial paper.  In addition, the authors felt that a detailed specification would probably be somewhat controversial and thus likely to compromise the possibility of wide-spread agreement on, and use of, the basic structures of the format.  Hindsight also suggests that we were rather naive at the time concerning coordinates and it is fortunate that the detailed specification was postponed until greater experience could be obtained.  The descriptions of coordinates in the initial FITS paper are simply inadequate.  They provide no description of the meaning of the world coordinates and suggest a rather incomplete list of coordinate types. The use of a single rotation per axis cannot describe any general rotation of more than two axes.  While participating in the development of the AIPS software package of the National Radio Astronomy Observatory, Greisen (\\cite{kn:G1}, \\cite{kn:G2}) found it necessary to supply additional details to the coordinate definitions for both spectral and celestial coordinates. Since the latter have been widely used, a NASA-sponsored conference held in January 1988 recommended that they form the basis for a more general coordinate standard (Hanisch \\&\\ Wells~\\cite{kn:HW}); such a standard is described below.  The present work generalizes the set of world coordinate system (WCS) FITS keywords with a view to describing non-linear coordinate systems and any parameters that they may require.  Alternate keywords which should be supported are described.  It also addresses the questions of units, multiple coordinate descriptions, uncertainties in the coordinate values, and various other coordinate related matters.  Conventions for attaching coordinate information to tabular data are also described.  Paper II (Calabretta \\&\\ Greisen \\cite{kn:CG}) and Paper III (Greisen et al.~\\cite{kn:GVCA}) extend these concepts to the ideal, but non-linear angular and spectral coordinates used in astronomy.  Paper IV (Calabretta et al.~\\cite{kn:CVGTDAW}) then provides methods to describe the distortions inherent in the image coordinate systems of real astronomical data.  The complex questions related to time systems and to other kinds of coordinates will be deferred. ",
        "conclusions": "\\label{s:summary}  The changes to FITS-header keywords are summarized in Table~\\ref{ta:keyword}.  As described in Paper II, for one purpose, \\CROTA{i}\\ may in some cases be used instead of the new keywords so that the coordinate information may be understood by software systems which have yet to be converted to these new conventions."
    },
    "0207/astro-ph0207631_arXiv.txt": {
        "abstract": "The gravitational lensing properties of cosmological halos depend upon the mass distribution within each halo. The description of halos as nonsingular, truncated isothermal spheres, a particular solution of the isothermal Lane-Emden equation (suitably modified for $\\Lambda\\neq0$), has proved to be a useful approximation for the halos which form from realistic initial conditions in a CDM universe. We derive here the basic lensing properties of such halos, including the image separation, magnification, shear, and time-delay. We also provide analytical expressions for the critical curves and caustics. We show how the scale-free results we derive yield scale-dependent lensing properties which depend upon the cosmological background universe and the mass and collapse redshift of the lensing halos, according to the truncated isothermal sphere (TIS) model of CDM halos derived elsewhere. We briefly describe the application of these results to the currently-favored $\\Lambda$CDM universe. ",
        "introduction": "The gravitational lensing of distant sources has in recent years become one of the most powerful tools in observational cosmology (see, for example, \\citealt{soucail01} and references therein). Since the effects of gravitational lensing depend upon the redshift of the source, the cosmological background, and the distribution of matter in the universe, they can be used to constrain the cosmological parameters and the primordial power spectrum of density fluctuations from which structure originates. In addition, many of the effects produced by gravitational lenses, such as image multiplicity, separations, and time delay, depend strongly upon the matter distribution inside the lenses. Hence, measurements of these effects can provide a unique tool for probing the matter distribution inside collapsed objects like galaxies and clusters, providing the only direct measurement of their dark matter content, and constraining the theory of their formation and evolution.  Until recently, the internal structure of halos adopted in lensing studies was generally some gravitational equilibrium distribution, either singular or nonsingular (e.g., King model, singular isothermal sphere, pseudo-isothermal sphere), not necessarily motivated directly by the theory of cosmological halo formation (see, e.g., \\citealt{young80}; \\citealt{tog84}; \\citealt{hk87}; \\citealt{nw88}; \\citealt{bsbv91}; \\citealt{jaros91,jaros92}; \\citealt{kochanek95}; \\citealt{pmm98}; \\citealt{premadi01}; \\citealt{rm01}). As the theory of halo formation in the CDM model has advanced in recent years, however, the halo mass profiles adopted for lensing models have been refined to reflect this theory. Numerical simulations of large-scale structure formation in Cold Dark Matter (CDM) universes predict that galaxies and clusters have a singular density profile which approaches a power law $\\rho\\propto r^{-n}$ at the center, with the exponent $n$ ranging from 1 to 1.5 (\\citealt{cl96}; \\citealt{nfw96,nfw97}; \\citealt{tbw97}; \\citealt{fm97,fm01a,fm01b}; \\citealt{moore98,moore99}; \\citealt{hjs99}; \\citealt{ghigna00}; \\citealt{js00}; \\citealt{klypin00}; \\citealt{power02}). These results are in apparent conflict with observations of rotation curves of dark-matter-dominated dwarf galaxies and low surface brightness galaxies, which favor a flat-density core (cf. \\citealt{primack99}; \\citealt{bs99}; \\citealt{moore99}; \\citealt{moore01}). On the scale of clusters of galaxies, observations of strong gravitational lensing of background galaxies by foreground clusters also favor the presence of a finite-density core in the centers of clusters (see, e.g., \\citealt{tkda98}).  Several possible explanations have been suggested in order to explain this discrepancy. The rotation curve data might lack sufficient spatial resolution near the center to distinguish unambiguously between a density profile with a flat-density core and one with a singular profile (e.g. \\citealt{vdbs01}). Attempts have also been made to improve the numerical resolving power of the simulations to obtain a more accurate determination of the slope of the predicted density profiles at small radii (e.g. \\citealt{moore99}; \\citealt{power02}). However, if the flat-core interpretation of the observations and the singular cusps predicted by the numerical simulations are both correct, then the simulation algorithms may be ignoring some physical process which would, if included, serve to flatten the halo density profiles at small radii relative to the results for purely gravitational, N-body dynamics of cold, collisionless dark matter, while retaining the more successful aspects of the CDM model. For example, gasdynamical processes (see, e.g. \\citealt{esh01}) and a modification of the microscopic properties of CDM, such as the proposal of self-interacting dark matter \\citep{ss00}, both have the potential to lower the central density of halos and possibly reconcile simulations with observations.  Lensing by the two kinds of halo mass profiles, singular versus flat-core, will be different. This has led to attempts to predict the differences expected if the halos have the singular cusp of the NFW or Moore profiles or else a profile with a flat core (e.g. \\citealt{kochanek95}; \\citealt{km01}; \\citealt{rm01}; \\citealt{wts01}; \\citealt{tc01}; \\citealt{lo02}). Singular profiles like that of NFW are physically motivated by the N-body simulations, and the latter have been used to place these halo profiles empirically in a proper cosmological context which permits statistical predictions for the CDM model. The nonsingular profiles which have been adopted to contrast with these singular ones, however, are generally no more than parameterized, mathematical fitting formulae, with no particular physical model to motivate them or put them in a proper cosmological context.  We have developed an analytical model for the postcollapse equilibrium structure of virialized objects that condense out of a cosmological background universe, either matter-dominated or flat with a cosmological constant (\\citealt{sir99}, hereafter Paper~I; \\citealt{is01a}, hereafter Paper~II). This {\\it Truncated Isothermal Sphere\\/}, or TIS, model assumes that cosmological halos form from the collapse and virialization of ``top-hat'' density perturbations and are spherical, isotropic, and isothermal. This leads to a unique, nonsingular TIS, a particular solution of the Lane-Emden equation (suitably modified when $\\Lambda\\neq0$). The size $r_t$ and velocity dispersion $\\sigma_V$ are unique functions of the mass $M$ and formation redshift $z_{\\rm coll}$ of the object for a given background universe. The TIS density profile flattens to a constant central value, $\\rho_0$, which is roughly proportional to the critical density of the universe at the epoch of collapse, with a small core radius $r_0\\approx r_t/30$ (where $\\sigma_V^2=4\\pi G\\rho_0r_0^2$ and $r_0\\equiv r_{\\rm King}/3$, for the ``King radius'' $r_{\\rm King}$, defined by \\citealt{bt87}, p. 228).  Even though the TIS model does not produce the central cusp in the density profile of halos predicted by numerical CDM simulations at very small radii, it does reproduces many of the average properties of these halos quite well, suggesting that it is a useful approximation for the halos which result from more realistic initial conditions (Papers I, II; \\citealt{is01b} and references therein). In particular, the TIS mass profile agrees well with the fit by NFW to N-body simulations (i.e. fractional deviation of $\\sim20\\%$ or less) at all radii outside of a few TIS core radii (i.e. outside a King radius or so). It also predicts the internal structure of X-ray clusters found by N-body and gasdynamical simulations of cluster formation in the CDM model. For example, the TIS model reproduces to great accuracy the mass-temperature and radius-temperature virial relations and integrated mass profiles derived empirically from the simulations of cluster formation \\citep{emn96}. The TIS model also successfully reproduces to high precision the mass-velocity dispersion relation for clusters in CDM simulations of the Hubble volume by the Virgo Consortium \\citep{evrard02}, including its dependence on redshift for different background cosmologies. The TIS model also correctly predicts the average value of the virial ratio in N-body simulations of halo formation in CDM.  The TIS profile matches the observed mass profiles of dark-matter-dominated dwarf galaxies. The observed rotation curves of dwarf galaxies are generally well fit by a density profile with a finite density core suggested by \\citet{burkert95}, given by \\begin{equation} \\rho(r)=\\frac {\\rho_{0,B}}{(r/r_c+1)(r^2/r_c^2+1)}\\,. \\end{equation}  \\noindent The TIS model gives a nearly perfect fit to this profile, with best fit parameters $\\rho_{0,B}/\\rho_{0,{\\rm TIS}}=1.216$, $r_{c}/r_{0,{\\rm TIS}}=3.134$, correctly predicting the maximum rotation velocity $v_{\\rm max}$ and the radius $r_{\\rm max}$ at which it occurs. The TIS  model can also explain the mass profile with a flat density core measured by \\citet{tkda98} for cluster CL 0024+1654 at $z=0.39$, using the strong gravitational lensing of background galaxies by the cluster to infer the cluster mass distribution \\citep{si01}. The TIS model not only provides a good fit to the projected surface mass density distribution of this cluster within the arcs, but also predicts the overall mass, and a cluster velocity dispersion in close agreement with the value $\\sigma_v=1150$ km/s measured by \\citet{dressler99}.  Several authors have studied the effect of lensing by halos with a flat-density core (\\citealt{jaros91,jaros92}; \\citealt{kochanek95}; \\citealt{pmm98,premadi01}) or by NFW or Moore profiles that have been generalized, so that the inner slope of the density profile is arbitrary (\\citealt{km01}; \\citealt{rm01}; \\citealt{wts01}; \\citealt{lo02}). These particular density profiles are essentially mathematical conveniences without physical motivation. There is no underlying theoretical model in these cases that was used to predict the value of the core radius or the departure of the inner slope of the density profile from the value found by N-body simulations of CDM. By contrast, the TIS model is based on a set of physical assumptions concerning the origin, evolution, and equilibrium structure of halos in CDM universes. Observations of gravitational lenses have the potential to distinguish between the TIS profile and singular ones like the NFW profile, as several observable properties of gravitational lenses will be strongly affected by the presence, or absence of a central cusp in the density profile. One example of an important observable that can distinguish between various density profiles is the parity of the number of images. Lenses with nonsingular density profiles, such as the TIS, obey the {\\it odd number theorem}. The number of images of a given source is always odd, unless the source is extended and saddles a caustic (see \\citealt{sef92}, hereafter SEF, p.~172). Lenses with singular profiles, like the singular isothermal sphere, the NFW profile, or the Moore profile, need not obey this theorem, even for point sources. Most observed multiple-image gravitational lenses have either 2 or 4 images, and this may argue against profiles with a central core \\citep{rm01}. There are, however, other possible explanations for the absence of a third or fifth image. That image tends to be very close to the optical axis, and might be hidden behind the lens itself. Also, it is usually highly demagnified, and might be too faint to be seen.  We can use the TIS solution to model observed gravitational lenses individually. Alternatively, we can use the observations collectively to constrain the distribution of halo properties as characterized by the TIS solution. These properties, core radius, velocity dispersion, central density, and so on, depend upon the mass of the lensing halos and the redshift at which they form. Observational constraints on the statistical distribution of these properties will, in turn, impose constrains on the cosmological parameters and the primordial power spectrum of density fluctuations.  The problem of studying gravitational lensing of distant sources in an inhomogeneous universe can be divided into two parts. The first part consists of determining the intrinsic properties of the lenses. In particular, we need to determine the relationship between the observables (image multiplicity, magnification, brightness ratio, sheer, image separation, time delay, $\\ldots$) and the lens parameters. The second part consists of tying the lens properties to the cosmology. This involves using cosmological models of structure formation and evolution to determine the statistical distribution of the lens parameters, the clustering properties of the lenses, and the nature of their environments. In this paper, we focus on the first part, determining the intrinsic properties of individual lenses, which is an essential building block. The second part will be the subject of forthcoming papers.  The remainder of this paper is organized as follows. In \\S2, we derive the lens equation. In \\S3, we compute the critical curves and caustics. In \\S4, we study the properties of multiple images: separation, magnification, brightness ratios, and time delay. In \\S5, we place the scale-free description of these properties in the cosmological context of the CDM model and explain how the dimensionless parameters of \\S2--\\S4 are related by the TIS model to the properties of lensing halos in physical units. Summary and conclusion are presented in \\S6. ",
        "conclusions": "We have derived the lensing properties of cosmological halos described by the Truncated Isothermal Sphere model. The solutions depend on the background cosmological model through the critical surface density $\\sigma_{\\rm crit}$, which is a function of the cosmological parameters and the source and lens redshifts, and the TIS parameters $\\rho_0$ and $r_0$, which are functions of the mass and collapse redshift of the halo, and the cosmological parameters. By expressing the surface density of the halo in units of $\\sigma_{\\rm crit}$ and the distances in units of $r_0$, all explicit dependences on the cosmological model disappear, and the solutions are entirely expressible in terms of two dimensionless parameters, the central convergence $\\kappa_c$ and the scaled position $y$ of the source. We have computed solutions, and we provide either analytical expressions or numerical fits, for the critical curves and caustics, the image separations, the magnification and brightness ratios, the shear, and the time delay. Lensing of a point source by a TIS produces either one or three images, depending on whether the source is located outside or inside the radial caustic. When three images are produced, the central one is usually very faint, being highly demagnified. Degenerate image configurations occur when an extended source overlaps a caustic. Two images are produced when the source overlaps the radial caustic, while an Einstein ring with a central spot is produced when the source overlaps the tangential caustic, which is a single point located at $y_t=0$. These degenerate cases correspond to maxima of the total magnification, which diverges as the source size goes to zero. When three images are produced, the angular separation between the two outermost images depends strongly on $\\kappa_c$, but only weakly on the source location.  The lens properties are often qualitatively different at small and large $\\kappa_c$. For instance, at small $\\kappa_c$, $\\kappa_c\\ga1$, the image separation $\\Delta x$ decreases as source position $y$ increases (i.e. as the projected separation between the source position and the lens center increases) (Fig.~7), the brightness ratio $\\mu_1/\\mu_3$ increases with $y$ (Fig.~8), the shear $\\gamma_3$ of the outermost image decreases with $y$, and the time delay $\\tau_{23}$ decreases with $y$ for $y\\la y_r$ (Fig.~10), while these quantities behave differently for large $\\kappa_c$. This is easily understood. Multiple imaging can only occur when the central surface density $\\sigma(0)$ exceeds the critical density $\\sigma_{\\rm crit}$ (or equivalently $\\kappa_c>1$). Since $\\sigma(x)$ is a decreasing function of $x$, there is a natural scale in the system: the position $x_{\\rm crit}$ on the lens plane where $\\sigma(x_{\\rm crit})=\\sigma_{\\rm crit}$. If the lens profile was scale-free, as in the cases of a Schwarzschild lens or a singular isothermal sphere, $x_{\\rm crit}$ would be the only length scale in the problem, and the properties of the lens would be self-similar. But the TIS has a characteristic length scale, the core radius $r_0$, and the existence of this second length scale prevents the solutions from being self-similar.  This paper focused on the intrinsic properties of individual lenses described by the TIS model. It provides all the necessary formula one needs to study gravitational lensing by TIS halos in specific cosmological models. We will present such studies in forthcoming papers. As an illustration here, we applied the TIS model to the currently-favored $\\Lambda$CDM universe, to calculate the central convergence $\\kappa_c$ expected for TIS halos of different masses and collapse epochs. We found that high-redshift sources (e.g. $z_S\\approx3$) will be strongly lensed by TIS halos (i.e. $\\kappa_c>1$) only for cluster-mass halos. We also calculated the characteristic angular scale of image features produced by strong lensing by TIS halos relative to the Einstein radius $\\theta_E$ of a lens with the same total mass, for comparison with the results for other lensing halo mass profiles and with observed lensing systems."
    },
    "0207/astro-ph0207389_arXiv.txt": {
        "abstract": "The true costs of high performance computing are currently dominated by software.  Addressing these costs requires shifting to high productivity languages such as Matlab.  MatlabMPI is a Matlab implementation of the Message Passing Interface (MPI) standard and allows any Matlab program to exploit multiple processors. MatlabMPI currently implements the basic six functions that are the core of the MPI point-to-point communications standard. The key technical innovation of MatlabMPI is that it implements the widely used MPI ``look and feel'' on top of standard Matlab file I/O, resulting in an extremely compact ($\\sim$250 lines of code) and ``pure'' implementation which runs anywhere Matlab runs, and on any heterogeneous combination of computers.  The performance has been tested on both shared and distributed memory parallel computers (e.g. Sun, SGI, HP, IBM and Linux).  MatlabMPI can match the bandwidth of C based MPI at large message sizes. A test image filtering application using MatlabMPI achieved a speedup of $\\sim$300 using 304 CPUs and $\\sim$15\\% of the theoretical peak (450 Gigaflops) on an IBM SP2 at the Maui High Performance Computing Center.  In addition, this entire parallel benchmark application was implemented in 70 software-lines-of-code (SLOC) yielding 0.85 Gigaflop/SLOC or 4.4 CPUs/SLOC, which are the highest values of these software price performance metrics ever achieved for any application.  The MatlabMPI software will be made available for download. ",
        "introduction": "Matlab \\cite{Matlab} is the dominant programming language for implementing numerical computations and is widely used for algorithm development, simulation, data reduction, testing and system evaluation.  The popularity of Matlab is driven by the high productivity that is achieved by users because one line of Matlab code can typically replace ten lines of C or Fortran code.  Many Matlab programs can benefit from faster execution on a parallel computer, but achieving this goal has been a significant challenge. There have been many previous attempts to provide an efficient mechanism for running Matlab programs on parallel computers \\cite{MATABP,Morrow98,ParAl,RTExpress,Tseng99,MultiMATLAB, ParaMat,Fabozzi99,Matpar,MPITB,Quinn,CMTM}. These efforts of have faced numerous challenges and none have received widespread acceptance.  In the world of parallel computing the Message Passing Interface (MPI) \\cite{MPI} is the de facto standard for implementing programs on multiple processors. MPI defines C and Fortran language functions for doing point-to-point communication in a parallel program.  MPI has proven to be an effective model for implementing parallel programs and is used by many of the worlds' most demanding applications (weather modeling, weapons simulation, aircraft design, and signal processing simulation).  MatlabMPI consists of a set of Matlab scripts that implements a subset of MPI and allows any Matlab program to be run on a parallel computer. The key innovation of MatlabMPI is that it implements the widely used MPI ``look and feel'' on top of standard Matlab file I/O, resulting in a ``pure'' Matlab implementation that is exceedingly small ($\\sim$250 lines of code). Thus, MatlabMPI will run on any combination of computers that Matlab supports.  The next section describes the implementation and functionality provided by MatlabMPI.  Section three presents results on the bandwidth performance of the library from several parallel computers.  Section four uses an image processing application to show the scaling performance and the very high software price performance that can be achieved using MatlabMPI. Section five presents our conclusions and plans for future work. ",
        "conclusions": "The use of file I/O as a parallel communication mechanism is not new and is now increasingly feasible with the availability of low cost high speed disks.  The extreme example of this approach are the now popular Storage Area Networks (SAN), which combine high speed routers and disks to provide server solutions. Although using file I/O increases the latency of messages it normally will not effect the bandwidth. Furthermore, the use of file I/O has several additional functional advantages which make it easy to implement large buffer sizes, recordable messages, multi-casting, and  one-sided messaging.  Finally, the MatlabMPI approach is readily applied to any language  (e.g. IDL, Python, and Perl).  MatlabMPI demonstrates that the standard approach to writing parallel programs in C and Fortran (i.e. using MPI) is also valid in Matlab.  In addition, by using Matlab file I/O, it was possible to implement MatlabMPI entirely within the Matlab environment, making it instantly portable to all computers that Matlab runs on.  Most potential parallel Matlab applications are trivially parallel and don't require high performance.  Never-the-less, MatlabMPI can match C MPI performance on large messages.  The simplicity and performance of MatlabMPI makes it a very reasonable choice for programmers that want to speed up their Matlab code on a parallel computer.  MatlabMPI provides the highest productivity parallel computing environment available.  However, because it is a point-to-point messaging library, a significant amount code of must be added to any application in order to do basic parallel operations.  In the test application presented here, the number of lines of Matlab code increased from 35 to 70.  While a 70 line parallel program is extremely small, it represents a significant increase over the single processor case.  Our future work will aim at creating higher level objects (e.g. distributed matrices) that will eliminate this parallel coding overhead (see Figure~\\ref{fig:LayeredArch}).  The resulting ``Parallel Matlab Toolbox'' will be built on top of the MatlabMPI communication layer, and will allow a user to achieve good parallel performance without increasing the number lines of code."
    },
    "0207/astro-ph0207340_arXiv.txt": {
        "abstract": "The formation and evolution of  active regions is an inherently complex phenomenon. Magnetic fields generated at the base of the convection zone follow a chaotic evolution  before reaching the solar surface. In this article, we use a 2-D probabilistic Cellular Automaton (CA) to model the statistical properties of the magnetic patterns formed on the solar surface and to estimate the magnetic energy released in the interaction of opposite polarities. We assume that newly emerged magnetic flux tubes stimulate the emergence of new magnetic flux in their neighborhood. The flux-tubes move randomly on the surface of the sun, and they cancel and release their magnetic energy when they collide with magnetic flux of opposite polarity, or diffuse into the ``empty'' photosphere. We assume that cancellation of magnetic flux in collisions causes ``flares\" and determine the released energy as the difference in the square of the magnetic field flux ($E\\sim B^2$). The statistics of the simulated ``flares\" follow a power-law distribution in energy, $f(E)\\sim E^{-a},$ where $a=2.2\\pm 0.1$. The size distribution function of the simulated active regions exhibits a power law behavior with index $k\\approx 1.93 \\pm 0.08$, and the fractal dimension of the magnetized areas on the simulated solar surface is close to $D_F\\sim 1.42\\pm 0.12.$ Both quantities, $D_F$ and $k$, are inside the range of the observed values. ",
        "introduction": "Many solar phenomena, such as coronal heating and solar flares, are closely related to the evolution of  active regions. Active regions are interpreted in this article as domains of strong magnetic field on the solar surface.  The appearance of active regions on the solar surface is the result of the complex interplay between the buoyant forces and the turbulent convection zone working in the solar interior \\citep{Par79}. Convection zone dynamics have been studied numerically \\citep{Nor96}, but the magneto-convection still remains in its infancy and many questions remain unanswered (see review by \\citet{Wei97}).  Many models have been suggested for the formation and evolution of active regions, such as the rise of a kink-unstable magnetic flux tube \\citep{Mor92}, or the statistical description of the dynamical evolution of large scale, two dimensional, fibril magnetic fields \\citep{Bog85}. \\citet{Sch97} assumed that the thin flux tubes that constitute the active region move and interact only in the intergranular lanes, and form in this way a fractal pattern. Several models have been developed using the anomalous diffusion of magnetic flux in the solar photosphere in order to explain the fractal geometry of the active regions \\citep{Law91, Sch92, Lsr93, Mze93}. Last, a percolation model was used to simulate the formation and evolution of active regions \\citep{Wen92, Sei96}, which models the evolution of active regions by reducing all the complicated solar MHD and turbulence to three dimensionless parameters. This percolation model explains the observed size distribution of active regions and their fractal characteristics \\citep{Meun99}.    Numerous observational studies have investigated the statistical properties of active regions, using full-disc magnetograms and Ca II plage regions observations from the Mount Wilson Observatory and from the National Solar  Observatory (see review by \\citet{How96}). These studies have examined among other parameters the size distribution of  active regions, and their fractal dimension: {\\em The size distribution function} of the newly formed active regions exhibits a well defined power law with index $\\approx -1.94$, and active regions cover only a small fraction of the solar surface (around $\\sim 8\\%$) \\citep{Hzw93}. The \\textit{fractal dimension} of the active regions has been studied using high-resolution magnetograms by \\citet{Bak93}, and more recently by \\citet{Meun99}. These authors found, using not always the same method, a fractal dimension $D_F$ in the range $1.3<D_F<1.8$.    In this article, we model the emergence and evolution of magnetic flux on the solar surface  using a 2-D cellular automaton (CA), following techniques developed initially by \\citet{Sei96}. The dynamics of this automaton is probabilistic and is based on the competition between two ``fighting\" tendencies: \\textbf{stimulated} or \\textbf{spontaneous} emergence of new magnetic flux, and the disappearance of flux due to \\textbf{diffusion} (i.e.\\ dilution below observable limits), together with random \\textbf{motion} of the flux tubes on the solar surface. The basic new element we add to the \\citet{Sei96} model is that we keep track of the \\textbf{energy release} through flux cancellation (reconnection) if flux tubes of opposite polarities collide. We concentrate our analysis only on the newly formed active regions, since the old active regions undergo more complicated behavior, the dipoles are participating in large scale flows associated with differential rotation and meridional flow \\citep{Lei64, Wan94}. ",
        "conclusions": "We have introduced a probabilistic cellular automaton to explain the formation and evolution of active regions and the associated energy release. The rules used implement the competition between the stimulated emergence of new magnetic flux and the gradual diffusion, as well as the motion of flux across the photosphere. We have compressed all the complicated physics involved in the formation of active regions into three dimensionless parameters, the probabilities $P$,  $D$, and $D_m$ of which actually only $P$ and $D_m$ play a decisive, major role on the results. The fourth parameter $\\mathcal{E}$ plays no significant role in the evolution of the model (see Sec.\\ 2).  Our main results are: (i) The model yields a power-law distribution of the active region sizes, and the active regions form a fractal set. (ii)  The collision and cancellation of oppositely polarized magnetic   flux causes ``flares\". The energy distribution of the ``flares\" follows a power-law distribution with exponent $a=2.2 \\pm 0.1$. Varying the parameters, the power-law is conserved, just the power-law index change. (iii) The ``flares'' are distributed along the lines where regions of opposite flux meet (see Fig.~ \\ref{fig2}a). (iv) Using different values for the stimulation probability $P$, we find a close correlation between the size distribution exponent $k$ and the energy release index $a$ (see Fig.~\\ref{fig2}b).  The results derived from our model concerning size-distribution and fractal dimension are in quantitative and qualitative agreement with the existing observations and with the results of the percolation models proposed previously \\citep{Sei96,Meun99}. The basically new result of our model is the connection of photospheric dynamics to energy release. This primary energy release through magnetic reconnection may well be the basic physical cause for the various types of observed low atmosphere emission events, such as coronal bright points \\citep{Par00,Kru98}, X-ray networks flares \\citep{Krucker1997}, transition region impulsive EUV emission \\citep{Benz2002}, and H$\\alpha$ flares (Ellerman bombs \\citep{Georgoulis2002}). All these emissions may play a role in coronal heating (this point is still under strong debate), and in particular it was recently claimed that impulsive, chromospheric EUV emissions are important signatures of not just the coronal heating process but also of mass supply to the corona \\citep{Brown2000}.  Interesting to note is that all the mentioned observed explosive phenomena exhibit power-law distributions in energy, with slopes roughly in the range between $1.5$ and $2.5$, coinciding with what we find here for the primary photospheric energy release. We believe  that the observed discrepancies in the indices (especially on $a$) are directly related to the dynamical evolution of the active regions. The predicted correlation between the power-law indices of the distributions of the released energy and of the sizes of active regions can be checked on existing observations. These results may stimulate a new explanation for the statistical properties of impulsive solar energy release events in the low atmosphere.  We believe that CA models, when used in conjunction with MHD numerical codes, can become a valuable tool for the study of the statistical properties of active regions."
    },
    "0207/astro-ph0207176_arXiv.txt": {
        "abstract": "{ This paper presents the results of a high resolution radio imaging survey of all known (96) low-luminosity active galactic nuclei (LLAGNs) at \\d19. We first report new 2~cm (150~mas resolution using the VLA) and 6~cm (2~mas resolution using the VLBA) radio observations of the previously unobserved nuclei in our samples and then present results on the complete survey. We find that almost half of all LINERs and low-luminosity Seyferts have flat-spectrum radio cores when observed at 150~mas resolution. Higher (2~mas) resolution observations of a flux-limited subsample have provided a 100\\% (16 of 16) detection rate of pc-scale radio cores, with implied brightness temperatures $\\gtrsim\\,10^8\\,$K. The five LLAGNs with the highest core radio fluxes also have pc-scale `jets.' Compact radio cores are almost exclusively found in massive ellipticals and in type~1 nuclei (i.e. nuclei with broad \\ha\\ emission). Only a few `transition' nuclei have compact radio cores; those detected in the radio have optical emission-line diagnostic ratios close to those of LINERs/Seyferts. This indicates that some transition nuclei are truly composite Seyfert/LINER+\\hii\\ region nuclei, with the radio core power depending on the strength of the former component. The core radio power is correlated with the nuclear optical `broad' \\ha\\ luminosity, the nuclear optical `narrow' emission-line luminosity and width, and with the galaxy luminosity. In these correlations LLAGNs fall close to the low-luminosity extrapolations of more powerful AGNs. The scalings suggest that many of the radio-non-detected LLAGNs are simply lower power versions of the radio-detected LLAGNs. The ratio of core radio power to nuclear optical emission-line luminosity increases with increasing bulge luminosity for all LLAGNs. Also, there is evidence that the luminosity of the disk component of the galaxy is correlated with the nuclear emission-line luminosity (but not the core radio power). About half of all LLAGNs with multiple epoch data show significant inter-year radio variability. Investigation of a sample of $\\sim\\,$150 nearby bright galaxies, most of them LLAGNs, shows that the nuclear ($\\leq$150~mas size) radio power is strongly correlated with both the black hole mass and the galaxy bulge luminosity; linear regression fits to all $\\sim\\,$150 galaxies give: log P$_{\\rm 2cm}$ = 1.31($\\pm$0.16)~log~\\mmdo\\              $+\\,8.77$ and log P$_{\\rm 2cm}\\,=$ 1.89($\\pm$0.21)~log~L$_{\\rm B}$(bulge) $-\\,0.17$. Low accretion rates ($\\leq10^{-2}-10^{-3}$ of the Eddington rate) are implied in both advection- and jet-type models. In brief, all evidence points towards the presence of accreting massive black holes in a large fraction, perhaps all, of LLAGNs, with the nuclear radio emission originating in either the accretion inflow onto the massive black hole or from jets launched by this black hole - accretion disk system. ",
        "introduction": "The debate on the power source of low-luminosity active galactic nuclei (LLAGNs, i.e. low-luminosity Seyferts, LINERs, and ``transition'' nuclei [nuclei with spectra intermediate between those of LINERs and \\hii\\ regions]) is a continuing one. Their low emission-line luminosities \\citep[L$_{\\rm H\\alpha}~\\leq$ 10$^{40}$ erg s$^{-1}$ by definition;][hereafter H97a]{hoet97a} can be modeled in terms of photoionization by hot, young stars \\citep{termel85,filter92,shi92}, by collisional ionization in shocks \\citep{kosost76,foset78,hec80,dopsut95} or by aging starbursts \\citep{aloet99}.  On the other hand, evidence has been accumulating that at least some fraction of LLAGNs share characteristics in common with their more powerful counterparts - radio galaxies and powerful Seyfert galaxies. These similarities include the presence of compact nuclear radio cores \\citep{hec80}, water vapor megamasers \\citep{braet97}, nuclear point-like UV sources \\citep{maoet95,baret98}, broad \\ha\\ lines \\citep[][hereafter H97c]{hoet97c}, and broader \\ha\\ lines in polarized emission than in total emission \\citep{baret99}. If LLAGNs are truly scaled down AGNs then the challenge is to explain their much lower accretion luminosities. This requires either very low accretion rates ($\\sim\\,10^{-8}$ of the Eddington accretion rate, L$_{\\rm Edd}$) or radiative efficiencies (the ratio of radiated energy to accreted mass) much lower than the typical value of $\\sim\\,$10\\% \\citep[e.g.~Chapter~7.8 of][]{fraet95} assumed for powerful AGNs.  One well-known property of some powerful AGNs is a compact (sub-parsec), flat-spectrum (usually defined as $\\alpha~\\geq$ $-$0.5, S$_\\nu~\\propto \\nu^\\alpha$) nuclear radio source, usually interpreted as the synchrotron self-absorbed base of the jet which fuels larger-scale radio emission. Astrophysical jets are known to be produced in systems undergoing accretion onto a compact object \\citep[e.g.][]{pri93,bla93} so such compact radio sources in galactic nuclei may reasonably be considered a signature of an AGN. Much theoretical work \\citep[e.g.][]{beget84,lovrom96,falbie99} has been devoted to this disk-jet relationship in the case of galactic nuclei and it has been suggested that scaled-down versions of AGN jets can produce flat-spectrum radio cores in LLAGNs \\citep{falbie99}. Compact nuclear radio emission with a flat to inverted spectrum is also expected from the accretion inflow in advection-dominated \\citep[ADAF;][]{naret98} or convection-dominated \\citep[CDAF;][]{naret00} accretion flows, possible forms of accretion onto a black hole at low accretion rates \\citep{reeet82}. Flat-spectrum radio sources can also result through thermal emission from ionized gas in normal \\hii\\ regions or through free-free absorption of non-thermal radio emission, a process which probably occurs in compact nuclear starbursts \\citep{conet91}. The brightness temperature, T$_{\\rm b}$, in such starbursts is limited to log~[T$_{\\rm b}~$(K)]~$\\lesssim$ 5 \\citep{conet91}. Thus it is necessary to show that T$_{\\rm b}$ exceeds this limit before accretion onto a black hole can be claimed as the power source.  From the observational perspective, \\citet{hec80} showed that LINER nuclei tend to be associated with a compact radio source, and compact, flat-spectrum radio cores are known to be present in many `normal' E/S0 galaxies \\citep{sadet89,wrohee91,sleet94}. Flat-spectrum radio cores are, however, uncommon in normal spirals or Seyfert galaxies \\citep{ulvwil89,vilet90,sadet95}, though \\citet{ulvho01a} have pointed out that low-luminosity Seyferts have a higher incidence of flat-spectrum cores than more powerful Seyferts.   How does one distinguish accretion-powered LLAGNs from LLAGNs powered by hot stars or supernova shocks? Broad \\ha\\ lines and bright unresolved optical or UV sources are ambiguous indicators because they can all be produced in starburst models \\citep[see][]{teret92}. Searches for broader polarized \\ha\\ emission are currently feasible in only a few of the brightest LLAGNs \\citep{baret99}. X-ray emission from the nucleus can be confused with sources in the bulge or disk of the galaxy, even at \\textit{Chandra}'s sub arcsecond resolution. Further, all these indicators may be affected by viewing geometry, obscuration, and the signal-to-noise of the observations, the last problem being exacerbated by the low optical, UV, and X-ray luminosities of LLAGNs and the need to subtract the starlight. The radio regime, however, offers several advantages. Gigahertz (cm wave) radiation does not suffer the obscuration that affects the UV to infra-red. Also, at tens of gigahertz the problems of free-free absorption can be avoided in most cases. Finally, high resolution, high sensitivity radio maps can be routinely made with an investment of less than an hour per source at the Very Large Array{\\footnote{The VLA and VLBA are operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}} \\citep[VLA;][]{thoet80} and the Very Long Baseline Array \\citep[VLBA;][]{napet94}. At their resolutions of $\\sim\\,$100~milli-arcsec (mas) and $\\sim\\,$1~mas, respectively, it is easy to pick out the AGN, since any other radio emission from the galaxy is usually resolved out.  Closely related to these theoretical and observational issues are the increasing number of accurate mass determinations for ``massive dark objects'' (MDOs) in nearby galactic nuclei. Indeed, dynamic signatures of nuclear black holes are being found in almost every nearby galaxy studied. Mass estimates from high-resolution stellar-, maser-, or gas-dynamics are probably accurate to a factor of two \\citep{ricet98,gebet00}. Slightly less reliable (factor $\\sim\\,$2-5) estimates of M$_{\\rm MDO}$ for large samples are possible using the tight correlation between central stellar velocity dispersion of the galaxy bulge ($\\sigma_c$) and M$_{\\rm MDO}$ \\citep{merfer01,gebet00}. In this paper we assume that the few~$\\times$ 10$^{5-9}\\,$\\Msun\\ MDOs traced by dynamic methods are black holes, and interchangeably use the terms MDO and black hole. With the new MDO results the degeneracy between black hole mass and accretion rate when accounting for radiated luminosity is removed. Given AGN spectral energy densities and black hole masses in large samples of nearby AGNs, one can attempt to isolate factors which determine accretion rates and luminosities.  A weak correlation between the nuclear centimeter radio luminosity and the mass of the nuclear black hole has been found by several authors, e.g. \\citet{fraet98}, \\citet{yibou99} and \\citet{dimet01}. This result has been used as evidence for the existence of an ADAF-like inflow since, in such a flow, the radio luminosity is predicted to scale with the black hole mass and the accretion rate \\citep[e.g.][]{mah97}. Two significant weaknesses here are the low resolutions of the radio data, which allows significant contamination from other processes in the bulge, and the small sample sizes. Our 150~mas resolution radio observations of a large number of nearby LLAGNs considerably increase the number of LLAGNs with reliable black hole mass estimates \\textit{and} high resolution radio observations, allowing a better test of the relationship between the two quantities.  In the following sections we first define the samples used in the paper (Sect.~\\ref{secsample}), and report on new VLA and VLBA observations which complete our survey of the distance-limited (\\d19) sample of LLAGNs (Sects.~\\ref{secnewobs} and \\ref{secnewres}). Observational results on radio-variability in LLAGNs are also presented in Sect.~\\ref{secnewres}. In Sects.~\\ref{secdetrate} to ~\\ref{secgalbulge} we combine the new observations with those reported in the earlier papers in this series, and  present comprehensive results on the radio properties of LLAGNs in the distance-limited (\\d19) sample, and their implications for the central engines in LLAGNs. In Sect.~\\ref{secoverallmdo} we study the relationship between nuclear radio power and black hole mass, both for the distance-limited LLAGN sample and for the larger sample of nearby bright galaxies (most of them LLAGNs) defined in Sect.~\\ref{secsample}. The results are followed by a brief discussion (Sect.~\\ref{secdiscussion}) on the nature of the central engines in LLAGNs. ",
        "conclusions": "The completion of our radio survey of the 96 nearest (\\d19) LLAGNs from the Palomar sample of nearby bright galaxies has yielded the following main results: \\newline a)~All (16 of 16) of the LLAGNs investigated at mas resolution with the VLBA have pc-scale cores with brightness temperatures $\\gtrsim\\,10^8$~K. The five of these nuclei with the highest core flux all have pc-scale jets. The luminosity, brightness temperature, morphology, and variability of the radio emission all argue against an origin in star-formation related processes or as thermal emission. Thus, the core radio emission probably originates either in an accretion inflow onto a supermassive black hole or from jets launched by this black hole-accretion disk system. The latter explanation is supported by the morphologies and spectra of the radio cores; \\newline b)~there is no reason to believe that the remaining LLAGNs with compact radio cores (investigated at 150~mas resolution) are different from the 16 LLAGNs investigated at mas resolution. Thus, at least half of all LINERs and low-luminosity Seyferts probably contain accreting black holes. The incidence for transition nuclei is much lower; \\newline c)~compact radio cores are preferentially found in massive ellipticals and in type~1 nuclei, i.e.  nuclei in which broad \\ha\\ emission is present. For these nuclei, the core radio power is proportional to the broad \\ha\\ luminosity. The preferential detection of type~1 nuclei could result from: 1) observational selection effects, in which broad \\ha\\ and radio cores have been found only in the more powerful LLAGNs, 2) only the type~1 LLAGNs are bona-fide AGNs, or 3) within the unified scheme, type~1 nuclei are beamed and/or type~2 nuclei are free-free absorbed in the radio; \\newline d)~the radio power of the compact core is correlated with both the galaxy luminosity and the luminosity and width of the nuclear emission lines. These trends suggest that we have detected only the brighter LLAGNs, i.e. the true incidence of accreting black holes in LLAGNs is higher than indicated by our radio detections alone. \\newline e)~The core radio and nuclear emission-line properties of LLAGNs fall close to the low-luminosity extrapolations of more powerful AGNs, providing further support for a common central engine;  \\newline f)~low-luminosity Seyferts and LINERs share many of the same characteristics in the radio. The transition nuclei detected are those which are the closest, in terms of emission-line diagnostic ratios, to Seyferts and LINERs. Thus at least some transition nuclei are really composite Seyfert/LINER + \\hii\\ region nuclei, with the core radio power dependent on the Seyfert/LINER component; \\newline g)~with increasing bulge luminosity, the radio power increases more rapidly than the emission-line luminosity. Also, there is evidence that the disk luminosity of the galaxy is correlated with the nuclear emission-line luminosity (but not the core radio power). Both these factors contribute to a radio loud/quiet division between ellipticals and non-ellipticals at high bulge luminosities; \\newline h)~investigation of a sample of $\\sim\\,$150 nearby bright galaxies, most of them LLAGNs, shows that the nuclear ($\\leq$150~mas) radio power is strongly correlated with both the estimated MDO mass and the galaxy bulge luminosity. Partial correlation analysis on the two correlations yields the result that each correlation is meaningful even after removing the effect of the other correlation. The nature of these important correlations is discussed. Low accretion rates ($\\leq\\,$10$^{-2}$--10$^{-3}$ of the Eddington rate) are implied in both ADAF- and jet-type models; \\newline i)~about half of all LLAGNs investigated show significant inter-year variability at 2~cm and 3.6~cm.  \\newline In short, all evidence points towards the presence of accreting black holes in a large fraction, perhaps all, of LLAGNs."
    },
    "0207/astro-ph0207206_arXiv.txt": {
        "abstract": "{ We have mapped the large-scale structure of the Serpens cloud core using moderately optically thick ($^{13}$CO(1--0) and CS(2--1)) and optically thin tracers (C$^{18}$O(1--0), C$^{34}$S(2--1), and N$_2$H$^+$(1--0)), using the 16-element focal plane array operating at a wavelength of 3mm at the Five College Radio Astronomy Observatory. Our main goal was to study the large-scale distribution of the molecular gas in the Serpens region and to understand its relation with the denser gas in the cloud cores, previously studied at high angular resolution. All our molecular tracers show two main gas condensations, or sub-clumps, roughly corresponding to the North-West and South-East clusters of submillimeter continuum sources. We also carried out a kinematical study of the Serpens cloud. The  \\COI\\ and \\COII(1--0) maps of the centroid velocity show an increasing, smooth gradient in velocity from East to West, which we think may be caused by a global rotation of the Serpens molecular cloud whose rotation axis is roughly aligned in the SN direction. Although it appears that the cloud angular momentum is not sufficient for being dynamically important in the global evolution of the cluster, the fact that the observed molecular outflows are roughly aligned with it may suggest a link between the large-scale angular momentum and the circumstellar disks around individual protostars in the cluster. We also used the normalized centroid velocity difference as an infall indicator. We find two large regions of the map, approximately coincident with the SE and NW sub-clumps, which are undergoing an infalling motion. Although our evidence is not conclusive, our data appear to be in qualitative agreement with the expectation of a slow contraction followed by a rapid and highly efficient star formation phase in localized high density regions. ",
        "introduction": "\\label{sec:intro}  It is now well established that a substantial fraction of the young pre-main sequence stars in molecular clouds are found in compact groups (clusters) rather than being uniformly distributed throughout the clouds (e.g.  Gomez et al.~\\cite{Gea93}; Testi et al.~\\cite{TPN99}; Carpenter~\\cite{C00}). It is thus clear that the formation and early evolution of stellar clusters hold the key for the understanding of the galactic disk stellar population, such as the distribution of stellar masses at birth, or initial mass function (IMF), as well as the kinematical and multiplicity properties (Clarke et al.~\\cite{CBH00}; Kroupa~\\cite{K95}; Kroupa et al.~\\cite{Kea01}). While there is a general consensus that the latter two properties are linked to the clusters dynamical evolution, the nature of the IMF is still debated. Observations show that the emergent mass distribution of young stellar clusters is consistent with the field stars IMF (Palla \\& Stahler~\\cite{PS99}; Hillenbrand \\& Carpenter~\\cite{HC00}; Meyer et al.~\\cite{Mea00}), implying that the stellar mass distribution must be linked with the cluster formation process itself. Three main classes of models have been considered: those which link the final stellar masses to (i) the structure or fragmentation of  the parent molecular cloud (e.g. Elmegreen~\\cite{E97}; Myers~\\cite{M00}; Padoan \\& Nordlund~\\cite{PN01}); or (ii) the feedback of the protostellar accretion process (Adams \\& Fatuzzo~\\cite{AF96}); or, finally, (iii) competitive accretion between already formed (proto--)stars in clusters (Bonnell et al.~\\cite{Bea97};~\\cite{Bea01}).  The progressive improvement in millimeter wavelength techniques, in particular the development of bolometer arrays for large single dish telescopes and the upgrade of millimeter wave arrays,  have provided a new wealth of observational data to constrain theories (Andr\\'e \\& Motte~\\cite{AM00}; Testi \\& Sargent~\\cite{TS00}). Millimeter continuum surveys of the cluster forming clouds in Serpens (Testi \\& Sargent~\\cite{TS98}, hereafter TS98), $\\rho$--Oph (Motte et al.~\\cite{MAN98}; Johstone et al.~\\cite{Jea00}), NGC~1333 (Sandell \\& Knee~\\cite{SK01}), and Orion~B (Johnstone et al.~\\cite{Jea01}; Motte et al.~\\cite{Mea01}) have shown that the mass distribution of the prestellar cores, progenitors of  individual stellar systems, is consistent with the field stars IMF and thus significantly steeper that the mass function of the larger gaseous clumps in molecular clouds (e.g. Williams et al.~\\cite{WBM00}).  These observations suggest that the self-similarity of molecular clouds breaks down at the cores level, and the fragmentation or the small-scale structure of the clouds is the likely mechanism responsible for the stellar IMF.  In a first attempt to link the small-scale structure of the prestellar cores and young stellar objects in clusters with the kinematics and large-scale distribution of the molecular gas, Testi et al.~(\\cite{tsoo00}; hereafter TSOO) combined the high resolution 3~mm continuum survey of TS98 with new interferometric and single dish molecular line surveys of the Serpens core. The high resolution CS(2--1) observations showed high velocity molecular gas associated with outflows emanating from several of the protocluster members, while the single dish N$_2$H$^+$(1--0) revealed the presence of two main gaseous clumps, well separated in velocity, in which the TS98 sources are embedded. The combination of the new high and low resolution millimeter-wave data with previously published observations showed that the protocluster is structured in sub-clusters, each internally coherent, well separated both spatially and kinematically. The star formation efficiency and proto-stellar density within each core is a factor of $\\sim$10 higher than the average values. The fact that the molecular gas within each sub-cluster is kinematically coherent suggests that clusters are assembled through fragmentation of the molecular clump in coherent sub-clumps which produce independent sub-clusters (TSOO). It is not clear, however, whether the sub-clusters in Serpens will eventually merge into a final cluster, such as other young clusters (e.g. the Orion Nebula Cluster) which show little evidence for sub-structures (Clarke et al.~\\cite{CBH00}).  Similar results on the spatial segregation of pre-stellar cores have also been found in the $\\rho$-Oph star forming region (Johstone et al.~\\cite{Jea00}), but unfortunately the essential kinematical information is missing in this region. In both regions, the prestellar cores and protostars densities are too low for dynamical interactions to be efficient over the timescale in which the bulk of the mass of the final stars is assembled. Dynamical interaction and competitive accretion can possibly play a role for the final tuning of stellar masses but cannot play a major role in determining the global shape of the stellar IMF in these regions. Note however that both these regions are at least an order of magnitude less dense than the typical clusters associated with massive stars (such as the Orion Nebula Cluster).  Given for granted that, at least in Serpens and $\\rho$-Oph, the stellar IMF appears to be linked to the individual small-scale cloud fragments, each progenitor of a single stellar system, the question remains on how the fragmentation and collapse processes occur. There are currently two opposite views of these processes: a ``slow'' and a ``fast'' mode. In the first mode molecular clouds may contract as a whole in a slow fashion evolving through a series of quasi-equilibrium conditions; star formation is triggered in localized regions of the cloud when the density exceeds the threshold for gravitational collapse. Star formation in the cloud as a whole may be regulated by external agents, such as in the photoionization-regulated view of McKee~(\\cite{MK89}), or it may progressively accelerate until the feedback of newly formed stars disperse the molecular material, halting star formation (Palla \\& Stahler~\\cite{PS00}). In the ``fast'' mode, since turbulence is expected to decay on very short timescales within molecular clouds (Mac~Low~\\cite{McL99}), there is no force to balance the clouds against gravitational collapse and star formation occurs globally on a few free-fall times (Elmegreen~\\cite{E00}; Hartmann et al.~\\cite{Hea01}). From the observational point of view, the problem can be approached by studying the large-scale kinematical and physical conditions in star forming molecular clouds and compare them with position, mass and age distributions of forming stars (Palla \\& Stahler~\\cite{PS00}).  The Serpens cloud core appears to be an ideal place to test the predictions of the various models. At a distance of $\\sim$310~pc (de Lara et al.~\\cite{dL91}), the Serpens cloud core and protocluster have been the subject of numerous detailed observational studies in recent years, at infrared and (sub-)millimeter wavelengths. Among the most recent are: Davis et al.~\\cite{dav99}; Giovannetti et al.~\\cite{Gea98}; Hodapp~\\cite{H99}; Hogerheijde et al.~\\cite{Hea99}; Kaas~\\cite{K99}; Kaas \\& Bontemps~\\cite{KB01}; McMullin et al.~\\cite{mcm94}; \\cite{mcm00}; Williams \\& Myers~\\cite{WM99}; \\cite{WM00}; in addition to TS98 and TSOO. The reasons for these numerous observational efforts lie in its compactness, richness and relative proximity to the Sun, which make it possible to obtain and combine high resolution observations of the relatively large number of (proto-)stellar sources within the (proto-)cluster and large-scale observations in various tracers of the entire cloud core. Regions further away from the Sun cannot be adequately resolved with single dish telescopes and cannot be observed at the required sensitivity with current millimeter arrays, whereas nearby regions require a substantially larger observing time to be fully mapped.  Most of the previous single dish millimeter studies of the large-scale structure of the molecular clump, within which the (proto-)cluster is embedded, focused either on the study of molecular outflows by means of optically thick CO(2--1) observations (Davis et al.~\\cite{dav99}), or on the chemistry (McMullin et al.~\\cite{mcm94}; \\cite{mcm00}), or presented partial maps of the cloud core (White, Casali \\& Eiroa~\\cite{Wea95}). In this paper we complement these and our previous high resolution studies (TS98 and TSOO) with observations of the molecular clump large-scale structure using complete maps in moderately optically thick ($^{13}$CO(1--0) and CS(2--1)) and optically thin tracers (C$^{18}$O(1--0), C$^{34}$S(2--1), and N$_2$H$^+$(1--0)), taking advantage of the 16-element focal plane array operating at a wavelength of 3mm at the  Five College Radio Astronomy Observatory\\footnote{The Five College Radio Astronomy Observatory is operated with support from the National Science Foundation and with permission of the Metropolitan District Commission} (FCRAO). These new observations are used to derive refined values of the physical parameters of the cloud, and are combined with previous observations to study the kinematical and fragmentation properties of the cloud core.  \\begin{figure} \\vspace*{4mm} \\centerline{\\psfig{figure=fig1a.eps,width=6.8cm,angle=270}} \\centerline{\\psfig{figure=fig1b.eps,width=6.8cm,angle=270}} \\centerline{\\psfig{figure=fig1c.eps,width=6.8cm,angle=270}} \\caption[ ]{ Maps of the \\COI, \\COII\\ and \\N2H\\ integrated intensity in the velocity interval 4.0 to 12.0~\\kms. In the \\COI\\ map the value of the first contour is 1.5~\\Kkms and the contour interval is 1.0~\\Kkms. In the \\COII\\ and \\N2H\\ maps the value of the first contour and the contour interval are both 0.4~\\Kkms. The cross marks the reference position (see  Table~\\ref{tab:lines}) and the filled triangles represent the positions of the submillimeter continuum sources. The letters mark the positions of the main \\N2H\\ cores as defined in TSOO. The ``spur'' is also visible, an extended, lower surface brightness narrow region, stretching from the NW to SE clumps and continuing beyond these (see also Fig.~\\ref{fig:SCUBA850}). } \\label{fig:comap} \\end{figure}   The paper is organized as follows: in Sect.~\\ref{sec:due} and~\\ref{sec:mor} we describe the new observations, the morphology of the Serpens cloud in the observed tracers and derive the physical conditions of the molecular gas; in Sect.~\\ref{sec:disc} we discuss the cloud kinematical and fragmentation properties; finally, in Sect.~\\ref{sec:concl} we summarize the main results of this study, and we discuss the implications for cluster formation and evolution theories. ",
        "conclusions": "\\label{sec:concl}  We performed wide-field imaging (up to about $16'\\times 16'$) of the emission towards the Serpens cloud core using moderately optically thick ($^{13}$CO(1--0) and CS(2--1)) and optically thin tracers (C$^{18}$O(1--0), C$^{34}$S(2--1), and N$_2$H$^+$(1--0)). Our main goal was to study the large-scale distribution of the molecular gas in the Serpens region and to understand its relation with the denser gas in the cloud cores, previously studied at high angular resolution.  The \\COI\\ and \\COII(1--0) emission traces the distribution of the column density in the Serpens cloud. The \\COI\\ emission is extended and encompasses both the NW and SE clusters of submillimeter continuum sources, whereas the \\COII\\ emission shows the clear demarcation between the NW and SE sub-clumps, which are also separated in velocity (see also TSOO). Maps of the CS and \\CSI(2--1) emission show several condensations in the two main sub-clumps, NW and SE, but present otherwise a different appearance from \\N2H. The latter follows the spatial distribution of the submillimeter sources in both the NW and SE sub-clumps much more closely than the other molecular tracers. In both the NW and SE sub-clumps the distribution of the column density follows quite closely the map of the \\COII\\ and \\CSI\\ integrated emission, suggesting that the emission peaks are column density maxima and are not due to peaks in the excitation of the molecular gas. We also find that the SE and NW sub-clumps are virialised with masses typically of order $60\\, M_\\odot$ and peak number densities of order $\\nmh\\simeq 7\\times 10^5$~\\cmc.  The  \\COI\\ and \\COII(1--0) maps of the centroid velocity show an increasing but smooth gradient in velocity from E to W. Using most of the data in a map at once, by least-squares fitting maps of line-center velocity for the direction and magnitude of the best-fit velocity gradient, we find that \\COII\\ has a 1.25~\\kms~pc$^{-1}$ velocity gradient which is directed  from E to W. The velocity gradient is extended to the whole Serpens cloud and is unlikely that the \\COI\\ {\\it and} \\COII\\ centroid velocities are seriously affected by the outflows.  We can also reasonably exclude that the observed velocity gradient is due to separate clumps with different systemic velocities, or that it is due to a global, non-symmetric cloud contraction. We thus think that the observed velocity gradient may indeed be caused by a global rotation of the Serpens molecular cloud whose rotation axis is very close to the SN direction, similar to the average outflow axis found by TSOO. Following Goodman et al. (\\cite{good93}) we also calculated the parameter $\\beta$, defined as the ratio of rotational kinetic energy to gravitational energy. We find $\\beta \\simeq 0.07$ which, according to Goodman et al. (\\cite{good93}), is low enough to inhibit fragmentation due to rotationally driven instabilities. The cloud angular momentum is not sufficient for being dynamically important in the global evolution of the cluster. Nevertheless, the fact that the outflows are roughly aligned with it may suggest a link between the large-scale angular momentum and the circumstellar disks around individual protostars in the cluster. However, our data does not allow to prove this connection.  The normalized centroid velocity difference, $\\delta V_{\\rm c}$, is a typical indicator of underlying infall. Using CS and \\CSI\\ we find two large regions of the map where the negative values of $\\delta V_{\\rm c}$ are more concentrated, which are approximately coincident with the SE and NW sub0clumps. % We find that the SE and NW sub-clumps have greater non-thermal than thermal motions, and we calculate infall velocities of about 0.1 to 0.4~\\kms. % Together with the results of Williams \\& Myers~(\\cite{WM00}), our observations suggest that infalling conditions are met in a relatively large region where the column density of the gas exceeds a certain threshold; when a protostellar object is formed, its outflow locally acts against the infall.  Whithin the high column density regions, where the infalling signature is detected, the {\\it local} star formation efficiency is much larger than in the entire cloud. This is consistent with the picture that where the conditions for star formation are met, then the process is highly efficient. Given that outflows from young protostars appear to oppose locally to the infall, such high star formation efficiency can be achieved only if most of the objects in the cluster are formed over a very short timescale. This short formation timescale, which does not allow much time for interaction between protostars given the relatively low density of objects observed in Serpens, and the indication that the prestellar cores mass function is consistent with the field stars IMF (TS98), suggest that dynamical interactions and competitive accretion may not be the dominating processes in these types of regions.   The observational connection of the large-scale rotation and the individual stellar systems flows, and the extended slow inflow motion limited to the high column density regions, appear to be consistent with a slow global evolution of the cloud that locally reachs the conditions for gravitational collapse and star formation. The very short timescale for the formation of the bulk of the cluster is advocated by both the accelerated star formation view of Palla \\& Stahler~(\\cite{PS00}) or the dynamical star formation view of Elmegreen~(\\cite{E00}).  It is difficult, from our observations, to give a definitive answer to whether star formation occurs in a slow quasi static fashion up to the point of instability or whether the entire process is driven by the fast dissipation of turbulence in clouds. Our data appear to be in qualitative agreement with the expectation of a slow contraction followed by a rapid and highly efficient star formation phase in localized high density regions. Our evidence is not conclusive, but we provide a series of observational constraints with which numerical simulation have to be tested against quantitatively. From the observational point of view, on the one hand the possible connection between the large-scale rotation and the individual circumstellar regions should be made more precise by means of high angular observations, on the other hand comparisons between wide area studies of the physical and kinematical conditions within giant molecular clouds and the distribution and properties of the young (forming) stars within them should be extended to other complexes."
    },
    "0207/astro-ph0207030_arXiv.txt": {
        "abstract": "We present \\hst\\//STIS (1160--1700\\thinspace\\AA) echelle spectra of the cataclysmic variable (CV) star, QU Car, at three epochs. In catalogues this binary is classified as a nova-like variable. QU Car was observed three times in time-tag mode for 2300\\thinspace sec, 2600\\thinspace sec and 2600\\thinspace sec, allowing us to study the spectral time evolution on timescales down to $\\sim$10\\thinspace sec. We find evidence of a high-state non-magnetic CV at low inclination, with unusually high ionisation.  We observed narrow absorption lines ($\\sim$ few hundred \\kms wide) in N\\Vl1240, O\\Vl1371 and Si\\IVl1398, as well as broader (HWZI $\\sim1000\\kms$) emission in C\\IIIl1176, C\\IVl1549 and He\\IIl1640, all with a superposed absorption component. High ionisation is indicated by the He\\II\\ emission, which is unusually strong in comparison with C\\IV, and the relative strength of the O\\V\\ absorption line. The dereddened UV continuum spectral index of, on average, $-2.3$ suggests that disc accretion dominates the spectral energy distribution. In two observations velocity shifting is noted in the absorption lines on a timescale long enough not to repeat within the $\\sim2600$-sec exposures. The absorption superposed on the C\\IV\\ emission line moves coherently with the N\\V\\ and Si\\IV\\ absorption, suggesting the same origin for all absorption lines -- most likely to be in the accretion disc atmosphere.  Weak blueshifted absorption in N\\V\\ and C\\IV\\ provides evidence of an outflow component and we estimate a maximum outflow velocity of $\\sim2000\\kms$. This may be linked to a wind launched from further out in the disc than is typically seen in those high-state non-magnetic CV whose wind speeds are observed to reach to $\\ga4000\\kms$. Unusually, three ionisation stages of carbon -- C\\II, C\\III\\ and C\\IV\\ -- are present in emission, with line width increasing with higher ionisation. The presence of C\\II\\ in emission and the positive line-width/ionisation correlation is most easily reconciled with an origin in a disc chromosphere, beyond the influence of the EUV-emitting inner disc. ",
        "introduction": "Cataclysmic Variable (CV) stars are interacting binary systems in which a late-type secondary star overflows its Roche lobe and accretes matter onto a compact primary. Amongst these, the nova-like variables are defined as those that persist in a high mass-accretion rate. A result of such a heterogeneous label is that, amongst the nova-like variables, there are many objects that, displaying similarities to members of other CV classes, sit uneasily with their classification. In this paper we present high time-resolution UV spectroscopy of one such example -- QU Car.  At $m_{\\rm v}=11.4$, QU Car is one of the brightest known CV, so it is rather surprising that until now it has been mostly ignored in observational studies. From among the small number of hitherto published studies, Gilliland \\& Phillips \\shortcite{82gilliland} examined optical spectra (4300--4800\\thinspace\\AA) at time-resolutions down to 3 minutes. They derived a long period (10\\fh9) from radial velocity variations in the optical emission lines. They also made note of the relative weakness of the Balmer line emission, which they deemed due to a high temperature and high mass accretion rate in the disc and suggested a classification of old nova or nova-like variable -- a view in keeping with Schild's (1969) comment that the object displayed some characteristics indicative of an old nova. Since this time, QU Car has been classified in catalogues (Ritter \\& Kolb 1998, Downes et al. 2001) as a nova-like variable.  {\\it International Ultraviolet Explorer} (\\iue\\/) low-dispersion spectra reveal broad blueshifted UV absorption lines in N\\V\\ and C\\IV, with an emission component in C\\IV\\ and He\\II\\ (Knigge, Woods \\& Drew 1994), indicative of an optically thick disc accompanied by mass outflow. Knigge et al. (1994) noted variation in the C\\IV\\ line caused by a periodic increase in the strength of the emission component, but were unable to decide whether this variablity was linked to the orbital phase. They also commented on the strength of the O\\Vl1371 absorption line and the He\\IIl1640 emission -- both unusually strong for high-state non-magnetic CV (HnMCV).  The data that are analysed in this paper were obtained as part of an observing campaign to detect spectral signatures of disc-wind variability of short timescales. These signatures are predicted to show up as highly time-variable fine structure components in wind-formed UV lines (Proga, Stone \\& Drew 1998). To achieve the necessary time and wavelenght resolution, the Space Telescope Imaging Spectrograph (STIS) on the {\\it Hubble Space Telescope} (\\hst\\/) was employed with an echelle grating in photon-counting (TIME-TAG) mode. This allows a time-resolution of down to $\\sim10\\,$s and velocity resolution of down to $\\sim10$\\kms, with the trade-off that higher time-resolution will result in a reduced useable spectral resolution and vice-versa.  On examining the results of the \\hst\\/ observations we were struck by the ways in which QU Car distinguished itself from the other targets in the programme (the nova-like variables IX Vel and V3885 Sgr, presented in Hartley et al. 2002). In this paper we seek to present the current \\hst\\/ observations and then re-examine archive data, with a view to establishing a physical origin for the unusual features noted. In section \\ref{s:obs} we describe the observations and data extraction processes. The results are then presented in two ways: in section \\ref{s:mean} we describe the time-averaged spectrum obtained at each epoch of observation, then in section \\ref{s:timevar} we re-present the data as spectral time-series at 30-sec time-resolution and as continuum light curves. Finally, in section \\ref{s:discuss} we complare QU~Car with other non-magnetic HnMCV and consider the origin of QU~Car's UV line spectrum. ",
        "conclusions": "\\label{s:discuss} In this paper we have presented high time- and high spectral-resolution UV spectra of QU Car obtained at three epochs.  These data have pointed toward a higher than typical degree of ionization, and reveal a mix of fairly narrow, emission and absorption line features along with only modest evidence of mass loss.  Before considering the physical implications of the new spectroscopy, we review the constraints on the inclination of QU~Car.  Gilliland \\& Phillips proposed $i\\la60\\degr$ from the lack of orbital modulation in the binary's light curve and the relatively narrow optical emission features.  The UV emission lines are not particularly broad either, generally fitting within approximately $\\pm1000$\\kms.  The simple fact that a number of UV spectral lines are firmly in absorption, and that broadened blueshifted absorption is detected from time to time, also reinforces the case for low inclination in that line emission is accepted as the general pattern for high-state high inclination systems.  We can venture a little further than this and note that the observed mix of emission and absorption lines argues persuasively against re-classifying QU~Car as a magnetic system (see Warner 1995) -- a step that might otherwise be contemplated on the grounds of the strong He\\IIl1640 emission typifying this object.  \\begin{figure} \\vspace{9cm} \\special{psfile=fig8.eps hscale=45 vscale=45 voffset=-65 hoffset=-23} \\caption{STIS spectrum of QU Car (dataset o5bi07010) compared to a GHRS spectrum of the dwarf nova SS~Cyg in outburst (datasets Z3DV0104T and Z3DV0105T) and a STIS spectrum of the novalike variable IX Vel (dataset o5bi01010). All specta were downloaded via the MAST facility at the STScI.} \\label{fig:comp} \\end{figure}  We now compare the 1150--1660~\\AA\\ spectrum of QU Car with those of more typical high-state non-magnetic CV in order to make the case that QU~Car is extreme in terms of the apparent degree of ionization.  We choose as our comparison objects SS~Cyg near peak of outburst and IX~Vel, for the reason that similarly high quality \\hst\\ UV spectra have been obtained for them.  IX~Vel and SS~Cyg are also non-eclipsing systems, like QU~Car.  The SS~Cyg observations used (Z3DV0104T and Z3DV0105T) were obtained by one of us (KSL) and, to ensure the most up-to-date calibration, were retrieved from the MAST facility at STScI.  The IX~Vel observations were presented in Hartley et al (2002).  The spectra of these three sources are shown together in figure~\\ref{fig:comp}.  The three objects are placed in a sequence of increasing degree of ionization, with IX~Vel the least-ionized, QU~Car the most ionized and SS~Cyg somewhere in between.  This assessment is based largely on how the He\\IIl1640 and O\\Vl1371 line profiles compare with C\\IVl1549 -- these signatures of high excitation are indeed relatively strongest in QU Car.  Anther point of contrast between QU~Car and the other two systems concerns the broad absorption dips at $\\sim$1300~\\AA\\ and $\\sim$1340~\\AA . These are presumably due to disc atmospheric blanketing.  Wade \\& Hubeny (1998) show in their disc atmosphere models that these features would be expected to fade as the maximum disc temperature rises (see their figure 9).  That this happens to at least the $\\sim$1300~\\AA\\ feature is consistent with the probable identification of much of the blend with modest-to-low ionization Si~{\\sc iii}/O~{\\sc i} transitions.  The pattern in figure~\\ref{fig:comp} is that IX~Vel presents with the highest equivalent width in these dips, SS~Cyg is intermediate, whilst they are completely absent from QU~Car. This trend operates in the same sense as the rise in relative prominence of the O\\V\\ and He\\II\\ lines.  Two final points of interest with regard to figure~\\ref{fig:comp} are the larger number of interstellar lines in the spectrum of QU~Car compared to either IX~Vel or SS~Cyg, and the appearance of both the C\\IIIl1176 and C\\IIl1335 lines in emission in QU~Car (albeit with a superposed absorption component in the former).  The relatively rich interstellar line spectrum has to do with the greater distance to QU~Car. This is estimated to be at least $500\\,$pc (Gilliland \\& Phillips 1982), compared to $96\\,$pc and $166\\,$pc, from astrometric parallaxes, for IX Vel ({\\it HIPPARCOS}) and SS~Cyg (Harrison et al. 2001), respectively.  The seeming predeliction for carbon line emission may be due to a carbon over-abundance in the emission line gas.  Both the question of the distance to QU~Car and the abundances issue are focused on in a separate study (Drew et al. 2002).  Although wind signatures are scarcely apparent in the QU~Car spectrum shown in figure~\\ref{fig:comp}, the comparison with archive \\iue\\ data (section~\\ref{s:mag}) showed that they have been more apparent at other epochs.  However there is a clear point of contrast with objects like SS~Cyg and IX~Vel with regard to the maximum outflow velocities observed.  In QU~Car they are usually about 2000~\\kms, whilst in both SS~Cyg and IX~Vel this is in the region of 4000~\\kms.  This higher figure is indeed typical of HnMCV (see e.g. Prinja \\& Rosen 1995).  Very crudely, if it is assumed these speeds scale as the rotation speeds in the accretion disc where presumably the wind is launched, this factor of 2 difference may be viewed as implying a substantially larger launch radius in QU~Car -- perhaps up to a factor of 4 larger.  Quantitatively this conclusion is somewhat dependent on both the binary inclination and accretor mass -- qualitatively it will be hard to avoid.  Its implication, like the impression of relatively high ionization, is a higher than typical mass accretion rate.  In the above we have assumed that the UV continuum emission in QU~Car is dominated by disc accretion.  This seems uncontroversial given the similarity between the slope of its dereddened energy distribution and that of indisputably accretion-dominated HnMCV (see section~\\ref{s:mean}).  It is likely that all absorption features, including those seen superposed on the C\\III, C\\IV\\ and He\\II\\ emission profiles, form in the disc atmosphere.  This view is supported by the way in which the pattern of variability seen in N\\V\\ and Si\\IV\\ absorption during the Q3 observation is also apparent in the C\\IV\\ absorption component. The less variable He\\II\\ absorption component can be viewed as mimicking the less variable O\\V\\ profile.  \\begin{figure} \\vspace{9cm} \\special{psfile=fig9.eps hscale=45 vscale=45 voffset=-70 hoffset=-23} \\caption{Three ionisation stages of carbon present in the UV spectrum: C\\IIIl1176, C\\IIl1335 and C\\IVl1549, plotted on the same velocity scale. The C\\II\\ and C\\III\\ lines are taken from a mean of the three STIS spectra to improve the S/N ration, whereas the C\\IV\\ line is taken from I3. The narrow absorptions superposed on the C\\II\\ line are interstellar lines formed in the same transition.} \\label{fig:c} \\end{figure}  We now consider the likely origin of the line emission. Figure~\\ref{fig:c} shows the emission, from the three ionization stages of carbon present in the spectrum, plotted on a common velocity scale.  The C\\II\\ line is clearly the narrowest with a HWZI of order 600\\kms.  The HWZIs in the C\\III\\ and C\\IV\\ lines are not so easy to tell apart.  On fitting a single gaussian through the line wings we find that the FWHM, respectively for these two lines, is $\\sim$950\\kms and $\\sim$1100\\kms.  These figures translate to HWZIs of $\\sim$1200 and $\\sim$1400\\kms.  We do not correct for intrinsic multiplet splitting because the magnitude of the correction depends on the unknown flux ratios among the multiplet components.  Fortunately these splittings are not so large that correction for them would alter the basic finding -- namely, that broader emission is associated with higher ionization. The near equality of the C\\III\\ and C\\IV\\ emission widths probably indicates that the excited C\\III\\ line contains a substantial contribution from recombination.  To see all three carbon ion stages simultaneously, with this pattern of velocity widths probably points to an origin in a disc `chromosphere'.  This might be created by irradiation from the EUV-bright inner disc and accreting object.  Emission produced by an outflow from the accretor or the innermost disc is rendered less likely by the modest ionization of the C\\II\\ and C\\III\\ lines (with the emissive flux in the C\\III\\ line only a little less than that in C\\IV ).  Since the emission component in each of the C\\III, C\\IV\\ and He\\II\\ lines is broader than the superposed absorption, we would have to refer to a specific kinematic model for the system to position the emission line region relative to the absorption line region. In the absence of this, we cannot distinguish whether the line emission is optically thin and overlies the source of line absorption, or whether it comes from a different part of the disc.  The emergent picture of QU~Car is then that it is relatively highly-ionized compared to the more typical properties of HnMCV.  The narrowness of the UV spectral lines observed may have to do with their production at larger disc radii than usual.  Of course we cannot yet rule out the possibility that QU~Car might be oriented nearly pole on.  We have to leave as an unsolved mystery the origin of the short-term variability within the spectral line profiles (section~\\ref{s:timevar}) that has been revealed thanks to the HST/STIS TIME-TAG mode of data-taking.  Since this variability is not present at all epochs and indeed presents with different properties on the two occasions we have noticed it, it is unlikely to be connected to anything as fundamental as the binary orbit.  Our third set of data hints at a repetition timescale of around an hour if, that is, the detected absorption profile disturbances do in fact repeat.  To progress any further, another more intensive campaign of spectroscopic observation would be required.  Lastly, we have not yet been able to begin to find an explanation for the remarkable inverse P~Cygni character of, most notably, the N\\Vl1240 and Si\\IVl1397 line profiles detected in our third dataset. It is already clear that QU~Car is a remarkable binary that has been overlooked for too long.  In a separate study focusing on exploitation of the UV interstellar line data and new optical spectroscopy (Drew et al. 2002), we show that this binary may in fact be closely related to the supersoft sources and then present evidence of carbon over-abundance.  It is to be hoped that this re-evaluation of the status of QU~Car and the presentation here, of its striking UV spectroscopic properties, will significantly broaden our perception of the ways the Universe finds to configure such high-state compact binaries."
    },
    "0207/astro-ph0207083_arXiv.txt": {
        "abstract": "{We present an analysis of new and old high frequency VLBI data collected during the last 10 years at 5--22\\,GHz. For the jet components in the mas-VLBI jet, two component identifications are possible. One of them with quasi-stationary components oscillating about their mean positions. Another identification scheme, which formally gives the better expansion fit, yields motion with $\\sim 9$\\,$c$ for $H_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $q_0=0.5$. This model would be in better agreement with the observed rapid IDV and the expected high Lorentz-factor, deduced from IDV. } ",
        "introduction": "The S5 blazar 0716+714 is one of the most active BL Lac objects. It is extremely variable on time scales from hours to months. 0716+714 has yet no known redshift. Optical imaging of the underlying galaxy however reveals an estimate of $z \\geq 0.3$ (Wagner et al. 1996). In the radio bands 0716+71 is an intraday variable (IDV). It exhibits a very flat radio spectrum, extending up to at least 300\\,GHz. The variability appears to be correlated over wide ranges of the electromagnetic spectrum (Quirrenbach et al. 1991). The simultaneous variations between X-ray, optical and radio strongly suggest an intrinsic origin of variability. From IDV a high intrinsic Doppler factor could be expected. VLBI studies covering more than 20\\,years at 5\\,GHz show a core-dominated evolving jet extending to the north. The VLBI jet is oriented at 90$^\\circ$ with respect to the VLA jet. At present it is not clear if 0716+714 is a superluminal source and how fast the motion of the VLBI components is. There are several possible models proposed with motions between $0.05-0.5$\\,mas\\,yr$^{-1}$ but some of them are based on relative unfrequently sampled data (Eckart et al. 1986, 1987; Witzel et al. 1988; Schalinski et al. 1992; Gabuzda et al. 1998; Tian et al. 2000).  We present and discuss our preliminary results of the reanalysis of 10 years (25 epochs) of VLBI data on 0716+714 at frequencies between 5\\,GHz and 22\\,GHz. Our analysis is based on four epochs from the CJF-Survey at 5\\,GHz between 1992 and 1999 (\\cite{britzen}), five epochs at 8.4\\,GHz from astrometric observations (\\cite{ros}) and our own observations between 1994 and 1999, five epochs at 15\\,GHz from the 2cm-Survey from 1994 to 2001 (\\cite{Kellermann}) and 11 epochs at 22\\,GHz from Jorstad et al. (2001) and our own data from 1992 to 1997. ",
        "conclusions": "At present we prefer Model\\,2 which would be in better agreement with the observed rapid IDV. Due to the observed radio optical correlation, the rapid IDV in 0716+714 cannot be explained by refractive interstellar scattering. Therefore the source has to be very small ($T_{\\rm B}\\geq10^{17}$\\,K) and the Doppler factor has to be $\\geq50$. The large oscillation of the components in Model\\,1 should make one suspicious about the identification. Further analysis with a comprehensive error calculation will show which model will prevail."
    },
    "0207/astro-ph0207610_arXiv.txt": {
        "abstract": "{\\small A linear dependence of the amplitude of broadband noise variability on flux for GBHC and AGN has been recently shown by Uttley \\& McHardy \\cite{uttley:01}. We present the long term evolution of this rms-flux-relation for Cyg~X-1 as monitored from 1998--2002 with RXTE. We confirm the linear relationship in the hard state and analyze the evolution of the correlation for the period of 1996--2002. In the intermediate and the soft state, we find considerable deviations from the otherwise linear relationship. A possible explanation for the rms-flux-relation is a superposition of local mass accretion rate variations.} ",
        "introduction": "The Cyg~X-1 data presented here were obtained by the PCA onboard RXTE, mostly gained through our monitoring campaign 1998--2002 (for a description of the data see \\cite{pottschmidt:02a}). We split each light curve into segments of 1\\,s length and determine the mean flux of each segment. The segments are binned into 41 equally segmented flux bins, for each of which we calculate the power spectral density (PSD) of all contained lightcurve segments via the DFT (see, e.g., \\cite{{nowak:98b},{nowak:00a}}).  From all periodograms of each flux bin the mean PSD is determined using standard methods. For better statistics we chose to take into account only flux bins containing at least 20 periodograms.  Integrating the Poisson noise corrected PSD over the range $\\nu=1$--$32$\\,Hz, we arrive at the squared fractional rms variability. In a next step we multiply the fractional rms variability by the mean flux of the bin, to obtain the absolute rms variability $\\sigma$. ",
        "conclusions": "An explanation of the rms-flux-relation seems to be the modification of a theory which explains the PSD of GBHC by the superposition of accretion disk (AD) instabilities occuring at several radii in the AD that are propagated towards the innermost region of the disk \\cite{{lyubarskii:97a},{churazov:01a}}. Local $\\dot{M}$ variations, which are caused by short term changes of the disk viscosity and whose time scale is dependent on radius, are superimposed on long term $\\dot{M}$ modulations. If the fractional amplitude of the $\\dot{M}$ variations is independent of the long term $\\dot{M}$, a linear flux-rms relation will be observed \\cite{uttley:01}. The rms-flux-relation will break down if the corona is severely disturbed. This seems to be the case during the intermediate state of Cyg~X-1, in agreement with earlier results for the behavior of X-ray time lags \\cite{pottschmidt:00a}. Here, the observed X-ray time lag is much larger than during the normal hard state, which could indicate changes in the geometry of the AD corona, possibly related to the observed radio emitting outflow."
    },
    "0207/astro-ph0207426_arXiv.txt": {
        "abstract": "Using extensive \\hst\\ imaging, combined with \\cha\\ X-ray and Parkes radio data, we have detected the optical binary companion to a second millisecond pulsar (MSP) in the globular cluster 47 Tucanae. This faint ($V=22.3$) blue ($V-I=0.7$) star shows a large amplitude (60--70\\%) sinusoidal variation in both $V$ and $I$. The period (3.19066 hr) and phase of the variation match those of the MSP 47 Tuc W (which does not have an accurate radio timing position) to within 0.5 seconds and 1.2 minutes respectively, well within the 1~$\\sigma$ errors. The phase dependence of the intensity and color implies that heating of a tidally locked companion causes the observed variations. The eclipsing nature of this MSP in the radio, combined with the relatively large companion mass ($>$ 0.13 \\mdot) and the companion's position in the color-magnitude diagram, suggest that the companion is a main sequence star, a rare circumstance for an MSP companion. This system is likely to have had a complex evolution and represents an interesting case study in MSP irradiation of a close companion.  We present evidence for another optical variable with similar properties to the companion of 47 Tuc W. This variable may also be an MSP companion, although no radio counterpart has yet been detected. ",
        "introduction": "More than 50 millisecond pulsars (MSPs), half of all known, have been found in radio surveys of globular clusters, where the high stellar densities and interaction rates cause large numbers of neutron stars (NSs) to be spun up to millisecond periods (via binary production and subsequent accretion onto the compact object). Studies of these objects therefore offer insight into the formation and evolution of NS binaries (see, e.g., Rasio, Pfahl, \\& Rappaport 2000) and the frequency of stellar interactions in the dense cores of clusters, sampling different evolutionary channels compared to the disk population.  Recent improvements in the sensitivity of the Parkes radio telescope have resulted in a dramatic increase in the number of cluster MSPs detected. In particular, the 11 known MSPs in the cluster 47 Tuc (Manchester et al. 1991; Robinson et al. 1995) have recently been increased to 20 (Camilo et al. 2000, hereafter CLF00), with 16 of these currently having accurate (milli-arcsecond) timing positions (Freire et al. 2001a; Freire 2001). These enable studies of the cluster's gravitational potential well (Freire et al. 2001a) and of the intra-cluster medium (Freire et al. 2001b).  Another significant observational advance has been the ability of \\cha\\ to detect cluster MSPs in large numbers (Grindlay et al. 2001a, 2002; D'Amico et al. 2002).  Accurate (0\\farcs1--0\\farcs2) X-ray positions for the MSPs combined with the detection of large numbers of cataclysmic variables (CVs) and active binaries with both \\cha\\ and \\hst\\ (Grindlay et al. 2001a and Edmonds et al. 2002a, in preparation) has allowed the radio, X-ray and optical data to be placed on a common astrometric frame, good to $\\lesssim$0\\farcs1.  These advances, and the superb spatial resolution of \\cha\\ and \\hst, have allowed the recent detection of two optical counterparts to cluster MSPs, PSR~J0024$-$7203U, hereafter 47 Tuc U, in 47 Tuc (Edmonds et al. 2001a, hereafter EGH01), and PSR~J1740$-$5340, hereafter NGC 6397 A, in NGC 6397 (Ferraro et al. 2001). The MSP 47 Tuc U has a $\\sim$0.15 \\mdot\\ He white dwarf (WD) companion, a 10.3 hr binary period and small amplitude (0.004 mag in $V$) orbital variations caused by heating of one side of the companion (EGH01).  The $\\sim$0.45 \\mdot\\ companion in NGC 6397 A is either a subgiant or a heated main sequence (MS) star (Ferraro et al. 2001). The 32.5 hr, 0.1 mag orbital variations are caused by tidal distortion of the secondary by the NS, and are roughly sinusoidal with a (32.5/2) hr period. Another notable feature of NGC 6397 A is that the spectrum of the X-ray counterpart is relatively hard, and is suggestive of non-thermal emission. By contrast, 8 out of 9 of the 47 Tuc MSPs which are bright enough to have useful spectral information are soft (Grindlay et al. 2002).  Currently, NGC 6397 A appears to be unique among known MSPs in having a likely non-degenerate companion. The MSPs in 47 Tuc, for example, are either single, have likely He WD companions (e.g. 47 Tuc U), or have secondaries with masses of $\\sim$0.02--0.03 \\mdot\\ (CLF00) which cannot be normal MS stars.  As argued by Burderi et al. (2002) and Ferraro et al. (2001), the subgiant companion to NGC 6397 A may have been the star that recently spun up the MSP, consistent with the relatively small (characteristic) age for the system (D'Amico et al. 2001). The position of the star in the color-magnitude diagram (CMD) is roughly consistent with that of a mass-losing subgiant like that in the CV AKO~9 in 47 Tuc (Albrow et al. 2001).  The equally interesting possibility presented by Ferraro et al. (2001) is that an MS star has been captured in an exchange interaction with an NS already spun-up to millisecond periods. Further work is needed to test whether the position of the star in the CMD is consistent with estimates of heating of the companion by the MSP (Ferraro et al. 2001). The relatively large amplitude of the ellipsoidal variations in NGC 6397 A suggests a high inclination orbit (Ferraro et al. 2001; Orosz \\& van Kerkwijk 2002), consistent with the detection of eclipses in the radio (D'Amico et al. 2001).  Five out of eight eclipsing MSPs known have very low companion masses of $\\sim$0.02--0.03 \\mdot, and short orbital periods (these include the 47 Tuc MSPs J, O, and R). The three exceptions are NGC 6397 A; PSR~B1744$-$24A in the cluster Terzan~5, with an orbital period of $\\sim$1.8 hr, companion mass $\\sim$0.10 \\mdot, and displaying irregular eclipses (Lyne et al. 1990); and 47 Tuc W, an eclipsing MSP in 47 Tuc with a $\\sim$3.2 hr period but with a companion mass of $\\sim$0.15 \\mdot\\ (CLF00; see \\S~\\ref{sec.radio}).  We present here extensive \\hst\\ evidence for an optical counterpart to 47 Tuc W based on the detection of a faint, large amplitude, optical variable with an orbital period and phase which perfectly match those of 47 Tuc W within the uncertainties. The X-ray counterpart detected with \\cha\\ has properties which are similar to those of NGC 6397 A. This detection represents the best evidence yet obtained for an MSP with an MS companion, given uncertainties about the evolutionary status of NGC 6397 A.  Following a brief summary of the radio data available on 47 Tuc W, we present the astrometry, absolute photometry and time series photometry for the 47 Tuc W optical companion in \\S~2, where we also present evidence for a 2nd possible MSP companion with similar optical and X-ray properties. In \\S~3 we discuss the implications of these results. ",
        "conclusions": "\\label{sec.disc}  \\subsection{\\wopt}  As shown above, there is excellent period and phase agreement between \\wopt\\ and 47 Tuc W and we regard the association between these two objects as secure.  This marks the first time a radio pulsar has been localized with sub-arcsecond precision by means other than radio pulse timing or synthesis observations.  We now discuss the nature of \\wopt\\ and its variability in the context of all that is known about the binary system, including radio (\\S~\\ref{sec.radio}) and X-ray data.  As noted in \\S~\\ref{sec.radio}, the MSP 47 Tuc W was eclipsed for $\\sim$30\\% of its orbit on the only day in which it was detected at radio wavelengths.  The two eclipsing MSPs known in the field (PSRs B1957+20 and J2051$-$0827, with companion masses of 0.02--0.03 \\mdot) show, to first order, persistent eclipses at similar orbital phases (other eclipsing MSPs with similar binary parameters, in 47 Tuc, are far more poorly studied and relatively little is known in detail about their eclipsing properties).  By contrast, the MSPs in NGC~6397 (NGC 6397 A; D'Amico et al. 2001) and Terzan~5 (PSR~B1744$-$24A; Nice \\& Thorsett 1992) show eclipses that are variable in time and orbital phase.  These two systems, while significantly different from each other, are in addition distinct from the other eclipsing MSPs in having larger companion masses, $\\sim$0.45 \\mdot\\ and $\\sim$0.10 \\mdot, respectively.  Naturally we do not know if the eclipses in 47 Tuc W are persistent, but the minimum companion mass of 0.13 \\mdot\\ suggests it is a fundamentally different star from the $\\sim$0.02 \\mdot\\ objects.  The companion mass is possibly similar to that of 47 Tuc U's companion and other presumed He WD companions to MSPs, but the observed eclipse could not be caused by a He WD companion, and in any case as seen in Figure \\ref{fig.cmds} the photometry for \\wopt\\ appears inconsistent with a He WD.  The companion mass {\\em is} above the hydrogen burning limit.  It therefore appears plausible that \\wopt\\ is an MS star, the first such known example as a companion to an MSP (if the companion in NGC 6397 A is an evolved subgiant rather than an MS star).  The large amplitude, nearly sinusoidal light curve of \\wopt\\ is reminiscent of the optical light curves of PSRs B1957+20 (Callanan et al. 1995) and J2051$-$0827 (Stappers et al. 2001). This similarity suggests that the optical variations of \\wopt\\ are caused by rotational modulation, as we see alternately the heated and (relatively?) unheated sides of the tidally locked companion once per orbit (as expected, \\wopt\\ appears to be hottest when it is brightest; Fig. \\ref{fig.w29col}). A similar, large amplitude light curve may also be present in the companion to PSR~B1744$-$24A in Terzan~5, but the formidable reddening and crowding in this cluster (Cohn et al. 2002) have thwarted a recent attempt to detect the companion with \\hst\\ (Edmonds et al. 2001b).  Since the blue color of \\wopt\\ probably results from the effects of heating, we have estimated the effects and intensity of the `radiation bath' \\wopt\\ experiences, based on the size of the binary and the expected spin-down luminosity (\\edot) of 47 Tuc W. Unfortunately, without a timing solution, no direct estimate of \\edot\\ is possible. Using the observed X-ray luminosity of 7.8$\\times 10^{30}$ erg s$^{-1}$ and the \\lx\\ vs \\edot\\ relationship of Grindlay et al. (2002), we estimate that \\edot\\ for 47 Tuc W is $\\sim 3\\times 10^{35}$ erg s$^{-1}$ (the 1~$\\sigma$ range is $\\sim 2-8\\times 10^{35}$ erg s$^{-1}$). A much smaller value of \\edot\\ (7.8$\\times 10^{33}$ erg s$^{-1}$) is derived using the linear relation of Becker \\& Tr\\\"umper (1999), \\lx$\\sim$10$^{-3}$\\edot, derived largely for field MSPs.  The stellar evolution models of Bergbusch \\& Vandenberg (1992) give an approximate guide to the relationship between luminosity and mass for a MS star.  We adopt the lowest mass model given by Bergbusch \\& Vandenberg (1992), a 0.15 \\mdot\\ model with \\teff\\ $=3300$~K and $M_V=13.2$, and test whether significant heating of such a star is expected.  We use the measured orbital period, an assumed mass for the NS of 1.4 \\mdot, and the above secondary mass to estimate the binary separation from Kepler's Third Law.  Then, using the estimated secondary radius from the stellar models, and assuming that the MSP radiates its wind isotropically, the `high' \\edot\\ estimate given above implies that the MSP energy intercepted by \\wopt\\ should be $\\sim 1 \\times10^{33}$erg s$^{-1}$, a factor of $\\sim 120 \\times$ the luminosity of the assumed companion (by comparison, the corresponding numbers for the companion to 47 Tuc U are $\\sim7 \\times10^{30}$erg s$^{-1}$ or 0.013$\\times$ the luminosity of the companion; EGH01).  Assuming that all of the spin-down energy is reradiated as a blackbody, we derive a temperature of 15,500~K. This would imply that heating effects are important and the MS star should be brighter and hotter than in its unperturbed state, possibly explaining the blue colors and relative brightness of \\wopt.  The above temperature is an upper limit (unless there is beamed emission), since a significant fraction of \\edot\\ could be powering other processes such as mass loss. For example, for PSR 1957+20, Callanan et al. (1995) find that the fraction of the MSP flux intercepted by the secondary that is converted into optical emission (defined as $\\eta$) is either in the range 0.07--0.2 for a Roche lobe filling model, or $\\eta \\gtrsim$ 3 (significant beaming) for a secondary considerably underfilling its Roche lobe. Stappers et al. (2001) find $\\eta \\sim$0.3 for J2051$-$0827 and EGH01 find $\\eta \\sim 1$ for 47 Tuc U.  Guided by these numbers we assume a conservatively small value for $\\eta$ of $\\sim$0.1 to give a reasonable lower limit on the heating flux, and derive a blackbody temperature of $\\sim$8,700~K.  Therefore heating effects should still be important.  For this model, the ratio between the heating flux and the undisturbed luminosity of the assumed stellar companion (120) is only a factor of 2.4 larger than the ratio (caused by heating) between the $V$-band flux of \\wopt\\ and the $V$-band flux of the stellar model (implying $\\eta \\sim$0.4).  However, if this model is approximately correct, the relatively small difference in intensity between the heated and unheated sides of \\wopt\\ (a factor of 6 compared to the above factor of 120) has to be explained.  One possibility is that heating of the companion has occurred fairly evenly over its surface. Although the secondary is probably tidally locked, significant heating of the darker side of the star might be expected because of energy transport within the star from the combined effects of convection (this low mass MS star should be fully convective), and the effects of relativistic $e^{+}e^{-}$ pairs and $\\gamma$-ray photons (D'Antona 1996). It is also possible that a crescent of the heated hemisphere is visible at all orbital phases if the orbital inclination of the binary is significantly less than 60\\degr\\ (Fruchter et al. 1995). This requires a higher mass companion for consistency with the radio results and the brighter star also gives a correspondingly smaller amplitude. In the case of NGC 6397 A, where the inclination is known to be high, the companion also has a much smaller variability amplitude than expected based on heating estimates (Orosz \\& van Kerkwijk 2002). This may be a feature common to all MSP systems containing non-degenerate companions.  Approximate upper limits to the mass and radius of the 47 Tuc W companion are set by the requirement that mass transfer be avoided.  Using the Roche-lobe formula from \\citet{pac71} ($r/a=\\mathrm{0.462[(M_{W}/(M_{NS}+M_W)]^{1/3}}$, where $r$ is the Roche-lobe radius, $a$ is the binary separation, $\\mathrm{M_{W}}$ is the mass of the secondary and $\\mathrm{M_{NS}}$ is the mass of the NS), we derive that, in its unperturbed state, a 0.29 \\mdot\\ model with \\teff $=3680$~K and $M_V=10.8$ would underfill its Roche lobe by $\\sim$1\\% (the corresponding upper limit on the orbital inclination is 27\\degr).  Our 0.15 \\mdot\\ model would (in its unperturbed state) significantly underfill its Roche lobe (the ratio of the stellar radius to $r$ is 0.6), explaining why the NS is not accreting.  However, under the influence of the extreme radiation bath from the MSP, \\wopt\\ may expand to fill or almost fill its Roche lobe (D'Antona \\& Ergma 1993), with accretion being inhibited by radiation pressure from the MSP (as may be occuring in NGC 6397 A; Burderi et al. 2002).  Using the `low' \\edot\\ estimate given above, the MSP energy intercepted by \\wopt\\ is only $\\sim4 \\times10^{31}$erg s$^{-1}$ (with a reradiated blackbody temperature of 6,300~K), a factor of $\\sim$3.3 times the luminosity of the assumed companion. This luminosity ratio is smaller than the observed variation and is a factor of 15 {\\em smaller} than the ratio between the $V$-band flux of \\wopt\\ and the $V$-band flux of the stellar model. So, unless significant beaming is present, only a fainter MS star than assumed above can give a larger ratio between the heating flux and the stellar luminosity. However, the disparity between the $V$-band fluxes then increases by the same amount.  A smaller ratio between the $V$-band fluxes is given by a brighter, more massive secondary (requiring an inclination $<$ 60\\degr), but then the heating flux would be too small to cause the observed variations.  These inconsistencies suggest that the higher \\edot\\ value given earlier is closer to the true value, providing indirect support for the \\lx\\ vs \\edot\\ relationship of Grindlay et al. (2002).  The similarities between the X-ray properties of W29 and those of the X-ray counterparts of NGC 6397 A (see \\S~\\ref{sec.ast}) and 47 Tuc J offer valuable insight.  These similarities at the very least allow for a similar physical origin of the X-ray emission.  The extended nature of the hard X-ray source in NGC 6397 A suggests it is due to shocked gas lifted from the binary companion (Grindlay et al. 2002), and similar behavior may be occurring in W29. This match in X-ray properties between W29 and two other eclipsing MSPs suggests evidence for a trend, since apart from 47 Tuc J, all 9 of the X-ray-detected MSPs in 47 Tuc that are bright enough to have useful hardness ratios determined have soft spectra (caused by likely thermal emission from their polar caps; Grindlay et al. 2002).   As we have already mentioned, 47 Tuc W is an extremely unusual (and so far possibly unique) MSP in having a likely MS companion. This likely represents evidence for stellar interactions, not surprising for a system so close to the cluster center. One possible formation scenario for this system is that an NS was spun up by accretion to form an MSP, which eventually evolved into an MSP/very low mass companion binary. Later, a collision between the MSP binary and the currently observed MS star caused the very low mass star (the lightest object amongst the 3 stars) to be ejected from the system. Alternatively, a single MSP could have exchanged into a double MS star binary, although in this case tight constraints on the mass of the ejected star may be necessary unless \\wopt\\ was once considerably more massive than it is today. Significant mass loss could be taking place if \\wopt\\ is close to filling its Roche lobe and material is being expelled from the system under the propeller mechanism (Burderi et al. 2002). A third possibility is that the observed binary was formed as a result of direct 2-body tidal capture (Mardling 1995) between a single MSP and an MS star, although this process may be much less important than binary interactions (Rasio et al. 2000).  In testing these various exchange scenarios, we note that 47 Tuc W is located very close to the cluster center (only 3\\farcs8 or 0.16$r_c$ away), consistent with the centrally concentrated spatial distribution of the other MSPs. If the MSP was initially ejected from the cluster core in an exchange encounter it may have had time to sink back into the cluster core, since the MS star could have been captured several Gyr ago, and the central relaxation time for 47 Tuc is 0.1 Gyr (increasing to 3 Gyr at the half-mass radius of 2\\farcm8; Harris 1996). The NGC 6397 MSP, by contrast, is found $\\sim$30$''$ from the cluster center (Ferraro et al. 2001) and is therefore well outside the cluster core ($r_c = 3 $\\farcs0; Harris 1996), possibly showing evidence for ejection from the core (D'Amico et al. 2001; Grindlay et al. 2002). This ejection may have occurred quite recently, since the central relaxation time for the core-collapsed NGC 6397 is only 0.08 Myr, increasing to 0.3 Gyr at the half-mass radius of 2\\farcm3 (Harris 1996). If these numbers are accurate, a system as heavy as NGC 6397 A should sink rapidly back into the core (unless it is in a highly elliptical orbit). One possible explanation is that NGC 6397 A was formed very recently by an NS capturing a subgiant, and was ejected from the core in this interaction. This scenario is broadly consistent with the relatively small (characteristic) age for the system of 0.35 Gyr (D'Amico et al. 2001). It may also have been ejected from the cluster core in a subsequent interaction. A similar ejection scenario may apply to PSR~B1744$-$24A, which is located $\\sim40''$ from the center of Terzan 5 (Nice \\& Thorsett 1992), equivalent to $\\sim5 r_c$ using the recent $r_c$ determination by Cohn et al. (2002).  It may be possible that \\wopt\\ will eventually fill its Roche lobe and cause mass accretion onto the NS to recommence, as may occur for NGC 6397 A (Ferraro et al. 2001 and Burderi et al. 2002). In this scenario, as \\wopt\\ keeps losing mass, it eventually may end up like the very low mass (0.02--0.03 \\mdot) systems such as 47 Tuc J (though these may be ablated He WDs rather than ablated MS stars; Rasio et al. 2000). Then, after continued ablation, it could possibly become an isolated MSP.  \\subsection{\\other}  Unlike with \\wopt, there is no period or phase match for \\other\\ with a known MSP (its period is similar to, but statistically different from, the 95.39 $\\pm$ 0.1 min period of 47 Tuc R, an eclipsing MSP with a very low mass companion; CLF00), but a significant fraction of the 47 Tuc MSPs have not yet been detected in the radio (CLF00; Grindlay et al. 2002).  Given the short period of this system, the heating effects (if it is an MSP) may be even more dramatic than for \\wopt, since the ratio between the heating flux and the stellar luminosity would be higher, assuming a similar luminosity companion star. We note that the color difference between the brighter and darker sides of the companion ($\\sim$30\\%) compared to the amplitude of the variation is larger than the similar ratio for \\wopt. The smaller amplitude of \\other\\ compared to \\wopt\\ may simply be an inclination effect or it may reflect a significant difference in the level of heating of the dark side of the companion, or show evidence for a different type of star, such as a very low mass degenerate star instead of an MS star.  We also note that in this system the low mass model assumed above for \\wopt\\ would overfill its Roche lobe by a few percent.  There are broad similarities in the appearance of the color vs phase plots for \\wopt\\ (Fig. \\ref{fig.w29col}) and \\other\\ (Fig. \\ref{fig.w34col}), particularly in the positions of the peak near \\phb\\ = 0.0, the minimum at \\phb\\ $\\sim$0.4, and the possible `hump' of relatively blue light at \\phb\\ = 0.4--0.7. The relatively blue color at optical maximum is secure, but the effects near \\phb\\ = 0.5 can only be considered marginal, given the faintness of the stars at these times and the possibility that artifacts of the \\hst\\ orbit may have leaked into the time series (half the period of \\wopt\\ is 95.72 min).  These comparisons do not constitute proof that \\other\\ is an MSP companion, and therefore we briefly consider alternative explanations for its behavior.  The near-sinusoidal light curve of \\other, with its short period and large amplitude, is similar to that of a W~UMa variable (a contact, or near contact binary consisting of two MS stars). However, this interpretation is clearly inconsistent with the blue colors of this object. A second possible explanation is that \\other\\ is a CV, since most of the CVs in 47 Tuc (and also most of the CVs in NGC 6397 and the field) are optically variable blue stars with moderately hard X-ray counterparts (Grindlay et al. 2001a). However, the likely CVs in 47 Tuc show flickering (random fluctuations on timescales of minutes with amplitudes of $\\sim$0.05--0.1 mag) which hides the presence of periodic variations in most cases. Of the few CVs in 47 Tuc which do show periodic variations, the light curves are very different from those of \\other\\ (and \\wopt), and show either low amplitude ($\\lesssim$ 0.1 mag) periodic variations from ellipsoidal modulation, or a combination of ellipsoidal modulation and eclipses (Edmonds et al. 2002a, in preparation), with distinctly non-sinusoidal light curves (there is significant signal in the power at both the orbital period and half the orbital period). Given these differences, then unless \\other\\ is a different type of CV from the $\\sim$20 good candidates found in 47 Tuc, we consider it more likely to be an MSP.  As shown in Figure \\ref{fig.w34tseries}, there are several instances of large deviations from the sinusoidal model for \\other. We examined the standard deviation of the residuals (data$-$sinusoid) for each cycle of the \\other\\ light curve, after iteratively removing 3~$\\sigma$ deviations from the mean, and compared these results with those found for \\wopt. The mean standard deviation for \\other\\ (0.15) was similar to the value found for \\wopt\\ (0.21), but there were some differences in the distributions. No 3~$\\sigma$ deviations were found for \\wopt, but two were found for \\other\\ (at HJD $\\sim$51363.6 and HJD $\\sim$51367.8), and two consecutive cycles of $\\sim 2.9 \\sigma$ deviations (in HJD = 51370.35--51370.5) were also found. These variations could be evidence for variable mass loss from the companion or they could be episodes of unusually large flickering, if the CV model is correct."
    },
    "0207/astro-ph0207432_arXiv.txt": {
        "abstract": "{ This letter presents incontrovertible evidence that NGC~5506 is a Narrow Line Seyfert 1 (NLSy1). Our new 0.9--1.4\\micron\\ spectrum of its nucleus clearly shows the permitted \\fullpiroi\\ line (with full width at half maximum $<\\,$2000~\\kms) and the `1 micron \\pfeii\\ lines'. These lines can only originate in the optically-thick broad line region (BLR) and, among Seyfert nuclei the latter series of lines are seen only in NLSy1s. The obscuration to the BLR, derived from a rough estimate of the \\fullpiroi/\\fullpopoi\\ ratio and from the reddening of the near-IR Paschen lines, is A$_{\\rm V}\\,>\\,5$. Together, these results make NGC~5506 the first identified case of an optically-obscured NLSy1. This new classification helps explain its radio to X-ray properties, which until now were considered highly anomalous. However, interesting new concerns are raised: e.g., NGC~5506 is unusual in hosting both a `type 1' AGN and a nuclear water vapor megamaser. As the brightest known NLSy1, NGC~5506 is highly suitable for study at wavebands less affected by obscuration. ",
        "introduction": "The Seyfert nucleus of NGC~5506 has resisted a clear type classification within Seyfert galaxies, and there is a long standing debate on whether it is an intermediate type~1 (broad \\ha\\ directly visible) or type~2 (broad \\ha\\ not directly visible) Seyfert. The presence of `broad' \\pb\\ has been reported by \\citet{blaet90}, \\citet{rixet90}, and \\citet{ruiet94}, but \\citet{gooet94} found that the `narrow' line emission profiles become broader at longer wavelengths and suggested that the `broad' \\pb\\ was the strong, highly reddened wings of this profile. Based on data available at that time \\citet{gooet94} interpreted the broadening of emission lines with wavelength as due to obscuration of the inner parts of the narrow line region. \\citet{morwar85} reported a marginal detection of \\fullpopoi, characteristic of Seyfert 1s, and suggested the presence of high-density optically thick gas. At odds with other type 2 objects, the nucleus of NGC~5506 is dominated by a bright compact core at all near-IR wavelengths and 60\\% of the J-band (1.25\\micron) flux in its central few arcsec is non-stellar in origin \\citep{oliet99,aloet01}. In the hard X-ray it is one of the most luminous and brightest Seyferts in the local universe \\citep[\\L2-10$\\,\\sim\\,10^{43}$;][]{mus82} and its obscuring column \\citep[\\nh$\\,= 3.4\\,\\times\\,10^{22}$ cm$^{-2}$;][]{baset99} is intermediate between typical values for Seyfert 1s and 2s. Nuclear water vapor masers, a property highly correlated with a type 2 spectral classification \\citep{braet96}, have been detected towards its nucleus \\citep{braet94}.  The host galaxy causes additional complications. The galaxy disk is close to edge on ($i\\,=\\,70{\\degr}$), and dust in the galaxy disk is responsible for some or all of the nuclear reddening \\citep{veiet97,ima00}. NGC~5506 is therefore variously treated as a type 1.9 or type 2 Seyfert in the literature and in either case is usually an outlier among the members of its class.  In this letter we report on near-IR spectroscopy of NGC~5506, which unequivocally identifies it as a Narrow Line Seyfert~1 (NLSy1). In a NLSy1 the broad line region (BLR) is directly visible with the BLR emission lines having widths typically $\\leq\\,2000\\,$\\kms, significantly narrower than those in classical Seyfert 1s. NLSy1s show several anomalous properties, most notably in the X-ray \\citep[for a nice overview of these see][]{veret01} and explanations for these include accretion rates close to the Eddington rate (implying lower black hole masses than other Seyferts) or a view to the AGN along its axis. ",
        "conclusions": "The currently used classification for NLSy1s \\citep[from][]{pog00} is: 1.~narrow permitted lines only slightly broader than forbidden lines; 2.~FWHM(\\hb) $<$ 2000~\\kms; 3.~\\oiii/\\hb$\\,<\\,$3, but exceptions allowed if there is also strong [\\ion{Fe}{VII}] and [\\ion{Fe}{X}] present, unlike what is seen in Seyfert 2s.  \\newline We have shown that the BLR emission is detected in the near-IR and that the \\poi\\ and \\pb\\ line profiles likely sample the bulk of the BLR. Thus, with \\poi\\ (from the BLR only) and \\pb\\ (from BLR and NLR) line widths $<$ 2000~\\kms, NGC~5506 directly satisfies the first two conditions. The observed \\oiii/\\hb\\ ratio is 7.5 at the nucleus and this ratio remains high over most of the extended emission-line region \\citep{wilet85}. If the BLR is highly extincted as our results suggest then the BLR contribution to the \\hb\\ flux would likely change the unextincted nuclear \\oiii/\\hb\\ ratio to $<\\,3$, bringing NGC~5506 into agreement with the third condition for classification as a NLSy1. A high extinction to the BLR would also explain the lack of strong optical \\pfeii\\ lines as usually seen in NLSy1s. NGC~5506 shares other properties unique to NLSy1s including the presence of the `1 micron \\pfeii\\ lines' as shown here, a high X-ray luminosity, steep X-ray slope, and fast X-ray variability \\citep{lamet00}. NGC~5506 is now the brightest known NLSy1 and therefore most suited for studies in wavebands not affected by obscuration. An important issue raised is whether several other X-ray bright and highly variable `type 2' Seyferts are, like NGC~5506, partially obscured NLSy1s.  Several properties of NGC~5506 still remain, or now become, anomalous. \\citet{matet01} find evidence that NLSy1s have preferentially lower black hole masses and are accreting at high values of L/L$_{\\rm Eddington}$.  However, the high central velocity dispersion in NGC~5506 \\citep[180~\\kms;][]{oliet99} though somewhat uncertain, implies a relatively high black hole mass among Seyferts, if the scaling between velocity dispersion and black hole mass is valid among Seyferts \\citep{wan02}. NGC~5506 is also unusual in being a type~1 AGN with a nuclear megamaser. Both the X-ray column \\citep{riset02} and narrow maser lines \\citep{braet96} are variable, and it may be that the latter are produced during periods when the column to nucleus is temporarily higher."
    },
    "0207/astro-ph0207118_arXiv.txt": {
        "abstract": "We consider the spin evolution of highly magnetized neutron stars in a hypercritical inflow just after their birth in supernovae. Presence of a strong magnetic field could deform the star and if the symmetry axis of the field is misaligned with that of stellar rotation, the star will be an emitter of gravitational wave. Here we investigate the possibility of gravitational radiation from such a star when there is a hypercritical inflow onto it. For doing this we adopt a simplified model of the system in which the star is approximated as a Newtonian spherical polytrope with index $N=1$. The stellar configuration is slightly deformed away from spherical by the intense magnetic field; the rotational angular frequency of the star is determined by the balance between the accretion torque and the magnetic dipole radiation. We take into account the 'propeller' process in which a rotating stellar magnetic field flings away in-falling matter; the inflow is assumed to be a self-similar advection dominated flow. An estimation of the characteristic amplitude of the gravitational radiation from such systems is given. The computation of the signal-to-noise ratio suggests that for the case of an initially rapidly rotating and highly magnetized star (surface field $10^{15}$G) in the Virgo Cluster, its ellipticity would need to be larger than $10^{-5}$ in order for the gravitational waves to be observed. ",
        "introduction": "Recently Watts \\& Andersson (2002) have considered the possibility of detecting gravitational radiation from a newly born neutron star in a supernova remnant. In their scenario, debris of a supernova explosion (typically $\\sim 0.1M_\\odot$) falls back hypercritically ($\\dot{M}>10^{-4}M_\\odot \\mbox{yr}^{-1}$) onto a neutron star with a normal-strength magnetic field ($B<10^{13}$G) and transfers angular momentum to the star. On the other hand, the r-mode instability\\footnote{ See Andersson \\& Kokkotas (2001) for a review of this instability and its astrophysical implications.} which is driven by gravitational radiation, can extract angular momentum from the system quite efficiently. The accretion torque and the torque of gravitational radiation balance to give the star a characteristic rotational period of a few milliseconds. As a result, this system can be an efficient emitter of gravitational waves of nearly constant frequency for several years. \\footnote{Yoshida \\& Eriguchi (2000) have considered the r-mode instability of a hypercritically accreting neutron star in a common envelope with a normal star companion.} Their work indicates that for stars with a normal-strength magnetic field ($B<10^{13}$G), the system can be a promising source for gravitational wave.  On the other hand, for moderately magnetized stars (i.e., $B\\sim 10^{10}-10^{12}$G) Rezzolla et al. have argued that the coupling of r-modes with the magnetic field can prevent the instability from developing. Thus the results by Watts \\& Andersson may not be relevant to the stars with rather strong magnetic fields ($B>10^{13}$G)  As a complementary study to theirs, we study here a similar system but with the neutron star having a much larger magnetic field ($B\\sim 10^{14}-10^{15}$G). Magnetic fields of this strength are associated with the so-called {\\it magnetars} (Duncan \\& Thompson 1996).  When the magnetic field reaches such large intensities, the equilibrium configuration would be distorted by the magnetic tension considerably. In the case of an axisymmetric magnetic field, the star will be deformed in an axisymmetric way aligned with the axis of the magnetic field (Bonazzola and Gourgoulhon 1996). If this symmetry axis is not aligned with the rotational one, this will produce a time varying quadrupole moment and lead to the emission of gravitational waves.  In this respect, our system is quite different from the one considered by Watts \\& Andersson (2002), since here gravitational radiation is not produced by the r-mode instability, and has only a passive role in the evolution of the system. The evolution of the system, i.e., of the rotational angular frequency and of the mass, is regulated by the evolution of the accretion torque.  In particular, the so-called 'propeller' mechanism (Pringle \\& Rees 1972, Illarionov \\& Sunyaev 1975) becomes active rather early. (This mechanism consists of the in-falling matter being flung away from the stellar magnetosphere when the rotational frequency is larger than the local orbital frequency, because the centrifugal force on the matter exceeds the gravitational force.) Also the torque from magnetic dipole radiation may be important in the later phases for highly magnetized stars.  The plan of the paper is the following. In section 2, the formulation of our model is outlined. In the following section, the results, (i.e., the evolution of the stellar rotational frequency and the mass, the characteristic amplitudes of gravitational waves from the system, and the signal-to-noise ratio of gravitational wave for laser-interferometric detectors), are presented for several parameter sets. The final section contains the summary and some comments about our model. ",
        "conclusions": "Using a simple model, we have computed the evolution of highly magnetized neutron stars with hypercritically in-falling matter just after their birth. The presence of a strong magnetic field would introduce a deformation of the star. If the axis of deformation were misaligned with the rotation axis, the object would emit gravitational waves with frequencies equal to the rotational one and its overtone. We have computed the characteristic amplitude of gravitational waves produced in this way. The results obtained indicate that if a strong magnetic field ($B\\sim 10^{15}$G) efficiently deforms the star ($\\beta\\sim 10$), we could observe the resulting signal from a source in the Virgo Cluster, when matched-filtering techniques are used in up-coming detectors.   Our model is built on several assumptions. Two of them in particular should be commented on.  Firstly, we assume that the accretion flow can be well-described by a simple advection dominated disk. In the present context, the accretion rate is such that an accretion shock is developed far outside the stellar surface (Armitage \\& Livio 2000; see eq.(17) of Brown et al. 2000). Internal to the shock, the accretion flow can be again hypercritical. However, it is possible that the system might have a jet or an explosive outflow to remove the matter in-flowing onto the central objects (Armitage \\& Livio 2000). In that case, the accretion rate onto the object would be significantly reduced. At present the mechanism of formation of an outflow from an accretion disk is unknown. For instance, the standard magnetically driven outflow (Blandford 1976; Lovelace 1976) may not work because the ram pressure of the inflow here is comparable with or exceeding that of the magnetic field of the central star.  Secondly we assume that the magnetic field inside and at the surface of the star does not change in the course of the evolution of the system. The decay of neutron star magnetic fields has been an important but unsolved issue which is related to the observational data for radio pulsars and low mass X-ray binaries (Possenti 1999). To explain the different strengths of the magnetic field in 'younger' systems (ordinary radio pulsars) and 'older' ones (low mass X-ray binaries or millisecond pulsars), some kind of mechanism for field decay seem to be needed. The 'accretion driven decay' model and the 'spin driven decay' model (Possenti 1999) might both be relevant for our model system here although their applicability is not clear since they have much stronger field. Also notice that the time scale of the decay in these scenarios should be much longer than the time scale of the evolution of the system (see also Zanotti \\& Rezzolla 2001)."
    },
    "0207/astro-ph0207604_arXiv.txt": {
        "abstract": "We analyze the data of the gravitational microlensing survey carried out by  by the MOA group during 2000 towards the Galactic Bulge (GB). Our observations are designed to detect efficiently high magnification events  with faint source stars and short timescale events, by increasing the the sampling rate up to $\\sim 6$ times per night and using Difference Image Analysis (DIA). We detect $28$ microlensing candidates in $12$ GB fields corresponding to $16$ deg$^2$. We use Monte Carlo simulations to estimate our microlensing event detection efficiency, where we construct the $I$-band extinction map of our GB fields in order to find dereddened magnitudes. We find a systematic bias and large uncertainty in the measured value of the timescale $t_{\\rm Eout}$ in our simulations. They are associated with blending and unresolved sources, and are allowed for in our measurements.  We compute an optical depth $\\tau = 2.59_{-0.64}^{+0.84} \\times 10^{-6}$ towards the GB for events with timescales $0.3<t_{\\rm E}<200$ days.  We consider disk-disk lensing, and obtain an optical depth $\\tau_{bulge} =  3.36_{-0.81}^{+1.11}\\times 10^{-6}[0.77/(1-f_{\\rm disk})]$ for the bulge component assuming a $23\\%$ stellar contribution from disk stars. These observed optical depths are consistent with previous measurements by the MACHO and OGLE groups, and still higher than those predicted by existing Galactic models. We present the timescale distribution of the observed events, and find there are no significant short events of a few days, in spite of our high detection efficiency for short timescale events down to $t_{\\rm E} \\sim 0.3$ days. We find that half of all our detected events have high magnification ($>10$). These events are useful for studies of extra-solar planets. ",
        "introduction": "\\label{sec:intro} Following the suggestion of \\cite{pac91} and \\cite{gri91}, several groups have carried out microlensing surveys towards the Galactic Bulge (GB), as seen in Baade's window. It is now well understood that these observations are useful for studying the structure, dynamics and kinematics of the Galaxy and the stellar mass function as the event rate and timescale distributions are related to the masses and velocities of lens objects.  The amplification of a microlensing event is described by (\\citealt{pac86}) \\begin{equation} \\label{eq:amp-u} A(u)= \\frac{u^2+2}{u\\sqrt{u^2+4}}, \\end{equation} where $u$ is the projected separation of the source and lens in units of the Einstein radius $R_{\\rm E}$ which is given by \\begin{equation} R_{\\rm E}(M,x) = \\sqrt{\\frac{4GM}{c^2}D_{\\rm s}x(1-x)}, \\label{eq:re} \\end{equation} where $M$ is the lens mass, $x=D_{\\rm l}/D_{\\rm s}$ is the normalized lens distance and $D_{\\rm l}$ and $D_{\\rm s}$ are the observer-lens and the observer-source star distances. The time variation of $u=u(t)$ is \\begin{equation} \\label{eq:u} u(t)=\\sqrt{\\beta^{2} + \\left( \\frac{t-t_{0}}{t_{\\rm E}} \\right)^2}, \\end{equation} where $\\beta$, $t_{0}$, $t_{\\rm E}= R_{\\rm E}/v_{\\rm t}$ and $v_{\\rm t}$ are the minimum impact parameter in units of $R_{\\rm E}$, the time of maximum magnification, the event time scale and the transverse velocity of the lens relative to the line of sight towards the source star, respectively. From a light curve, one can determine the values of $\\beta$, $t_{0}$ and $t_{\\rm E}$, but not the values of $M$, $x$ or $v_{\\rm t}$.  Our MOA (Microlensing Observations in Astrophysics) group started observations towards the GB in 1999. From 2000 we introduced the Difference Image Analysis (DIA) (\\citealt{cro92}; \\citealt{phi95}; \\citealt{ala98}; \\citealt{ala00}; \\citealt{alc99,alc00}; \\citealt{woz00}; \\citealt{bon01}) which is able to perform better photometry than the traditional DoPHOT (\\citealt{sch93}) type analysis in crowded fields at any place even where no star was identified.   To date, hundreds of microlensing events have been detected towards the GB by the OGLE (\\citealt{uda94,uda00}; \\citealt{woz01}) and MACHO collaborations (\\citealt{alc97a,alc00}). They estimate the microlensing optical depth towards the GB to be $3.3 \\pm1.2 \\times 10^{-6}$ from $9$ events by DoPHOT analysis, $3.9 ^{+1.8}_{-1.2} \\times 10^{-6}$ from $13$ events in a clump giant subsample from DoPHOT, and  $3.23^{+0.52}_{-0.50} \\times 10^{-6}$ from $99$ events by DIA respectively. \\cite{pop00} and  \\cite{pop02} estimate values of $2.0 ^{+0.4}_{-0.4} \\times 10^{-6}$ and $2.23 ^{+0.38}_{-0.35} \\times 10^{-6}$ respectively from MACHO data. These values are all more than twice those expected from existing Galactic models, which are somewhere around $0.5 \\sim  1.0 \\times 10^{-6}$ (\\citealt{pac91}; \\citealt{gri91}; \\citealt{kir94}).  This suggests that the standard models of the Galaxy need to be revised. To explain the high optical depth, a number of authors have suggested the presence of a bar oriented along our line of sight to the GB (\\citealt{pac94}; \\citealt{zha95}), and have adopted various values of the bar orientation and mass (\\citealt{pac94}; \\citealt{pea98}; \\citealt{zha96}). Microlensing observations towards the GB therefore appear useful for characterizing the mass and inclination of the bar.   \\cite{pop00} and  \\cite{pop02} raised the possibility of a systematic bias in the optical depth due to the difficulties of measuring $t_{\\rm E}$ associated with blending and unresolved sources.  When the actual source base-line flux is unknown, $t_{\\rm E}$ and $\\beta$ are degenerate in relatively low signal-to-noise ratio (S/N) events (c.f. \\citealt{han99}; \\citealt{bon01};\\citealt{gou02}). The optical depth may be estimated by using red clump giant stars to avoid the bias, or by other methods (e.g. \\citealt{gon99}; \\citealt{ker01}). In this paper, we quantify the bias by using Monte Carlo simulations and take itinto account to estimate the optical depth.    Our observations are designed to detect efficiently high magnification events with faint source stars for the study of extra-solar planets and surface-transit events, by increasing the sampling rate up to $\\sim 6$ times per night (note that the sampling rate of other projects is typically once per night). Our observations are consequently fairly sensitive to short duration events, i.e., events caused by smaller mass lenses. This could lead to a different optical depth estimate from previous studies if there is a significant contribution of low mass objects such as brown dwarfs to the microlensing optical depth. Thus the MOA observations can constrain the contribution of the low-mass population and also the structure of the GB. For this purpose here we present the results of MOA observations.   In this paper, the results of the DIA analysis of data towards the GB taken by MOA in 2000 with DIA are presented. The analysis is aimed at finding how efficiently we can detect high magnification events and short timescale events, and at estimating the optical depth towards the GB. In \\S\\,\\ref{sec:observation} we describe our observations. \\S\\,\\ref{sec:dataanalysis} is devoted to the analysis method. In \\S\\,\\ref{sec:results} we describe the event selection process and results. In \\S\\,\\ref{sec:extinction} we make an $I$-band extinction map of our GB fields caused by dust. This is useful in estimating extinction free source magnitudes in the simulation. In \\S\\,\\ref{sec:OPTICALDEPTH} we describe the simulation used to estimate our detection efficiency and the resultant microlensing optical depth. Discussion and conclusions are given in \\S\\,\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  We have re-analyzed the sample of subtracted images that were derived from the real-time DIA of GB observations obtained by MOA during 2000 (\\cite{bon01}).  In this analysis we have found 28 microlensing event candidates in our 12 GB fields. The DIA is more suitable than DoPHOT analysis for our purpose, since the former method can detect the luminosity variation at any position, even where no star was previously identified.    We have used a Monte Carlo simulation to estimate our event detection efficiencies. By using these efficiencies and timescales of our 28 detected events, we have estimated the optical depth towards the GB for events with timescales within the range $0.3<t_{\\rm E}<200$ days as $ \\tau_{0.3}^{200} = 2.63_{-0.58}^{+0.72} \\times 10^{-6}, $ where the statistical uncertainty of the optical depth have been estimated by Monte Carlo simulation only using the Poisson statistics. This $\\tau$ is overestimated and these errors in $\\tau$ are underestimated. We found the systematic bias and large uncertainty in the OUTPUT $t_{\\rm Eout}$ in our simulation. By taking these bias and uncertainty in $t_{\\rm Eout}$ into account, we get  $$ \\tau_{0.3}^{200} = 2.59_{-0.64}^{+0.84} \\times 10^{-6}. $$ The GB microlensing optical depth, in which the disk source stars component ($f_{\\rm disk}=23\\%$) is taken into account, is given by  $$ \\tau_{\\rm bulge} =  3.36_{-0.81}^{+1.11}\\times 10^{-6}[0.77/(1-f_{\\rm disk})]. $$ This value is  consistent with the previous observations of $3.3_{+1.2}^{-1.2} \\times 10^{-6}$ from $9$ events by DoPHOT analysis (\\citealt{uda94}), $3.9_{+1.2}^{-0.9} \\times 10^{-6}$ from $13$ events in a clump giant subsample by DoPHOT (\\citealt{alc97a}) and $3.23_{+0.52}^{-0.50} \\times 10^{-6}$ from $99$ events by DIA (\\citealt{alc00}), while slightly higher than $2.0^{+0.4}_{-0.4} \\times 10^{-6}$ (\\citealt{pop00}) and $2.23 ^{+0.38}_{-0.35} \\times 10^{-6}$ (\\citealt{pop02}).  This observed $\\tau$ must be regarded as  a lower limit of the true value, since our observations are only sensitive to events with $0.3< t_{\\rm E} <200$ days. Nevertheless this value is still higher than those predicted by most Galactic models whose  mass and inclination of the bar are consistent with other observations. Even the smallest inclination angle and a large bar mass could not reproduce $\\tau$ greater than $\\sim 2.5\\times 10^{-6}$ (\\citealt{pea98}). In \\cite{eva02} Freudenreich's model can reproduce the high optical depths toward the GB $\\sim 2.5 \\times 10^{-6}$ which is about 1 $\\sigma$ level of our estimate.  The smallest inclination produces the largest optical depth. However the bar inclination has been reported to be in a wide range between $10^\\circ-45^\\circ$ summarized in Table 7 in \\cite{alc00}. The optical depth also depends on the mass of the bulge or bar. However various observations provide conflicting values $M_{bulge|bar}=0.7-2.8 \\times10^{10}$ $M_\\odot$ (\\cite{zha96}; \\cite{dwe95}; \\cite{hol98}; \\cite{han95b}; \\cite{blu95}; \\cite{zha96}). Very small inclinations ($\\theta=11^\\circ$ with $M_{bar}=2.0 \\times10^{10}M_\\odot$ (\\citealt{zha96}) and $\\theta=12^\\circ$ with $M_{bar}=2.5 \\times10^{10}M_\\odot$ (\\cite{gyu99})) or a very heavy mass ($\\sim 3.6\\times 10^{10}M_\\odot$ with $\\theta = 20^\\circ$ (\\cite{gyu99})) are required to account for the observed optical depth in this analysis and in \\cite{alc00}. \\cite{bin01} estimate the minimum total mass in baryonic matter within the Solar circle to be greater than $\\sim 3.9\\times 10^{10}M_\\odot$ from $\\tau = 2.0 \\times 10^{-6}$, and such high baryonic contribution is consistent with implications from hydrodynamical modeling and the pattern speed of the Galactic bar.   Such a high mass would imply low halo MACHO fractions (\\citealt{gat96}). A massive bulge puts tight constraints on the contribution of the disk to the rotation curve at small radii. A small disk, however, leaves more room for the halo. Since microlensing results towards the LMC fix the MACHO content in the halo, a massive halo implies a smaller MACHO fraction.  The uncertainties in the Galactic bar orientation, the bar mass and the stellar mass function are still large. So the optical depth we have derived might yet be explained by other models. Further discussion can be seen in \\cite{alc00}.     We showed our observed timescale distribution which is not biased by blending. Though our statistics are smaller, this seems to be consistent with the one previously presented by the MACHO group (\\citealt{alc00}). The number of short timescale events was quite small, in spite of our high detection efficiency for events down to $t_{\\rm E} \\sim 0.3$ days.    A significant number of long timescale ($t_{\\rm E} >70$ days) events have been detected, as reported previously by the MACHO group (\\citealt{alc00}). These could not be explained by any current Galactic model (\\citealt{han96}). These might be a heavier remnant component, such as white dwarfs, or some dynamically cold component. Either way, we need more observations to investigate the mass function and Galactic structure in greater detail.    We have shown how efficiently MOA can detect the high magnification events in which the probability of detecting extrasolar planets is high and find $50\\sim 60 \\%$  of all detected events have high magnification ($u_{\\rm min} <0.1$). This fraction is much higher than $10\\%$ from the DoPHOT analysis (\\citealt{alc97a}) and $30\\%$ from recent DIA analysis (\\citealt{alc00}) by the MACHO group. This is because our sampling  rate is higher ($5\\sim6$ times\\,day$^{-1}$) than theirs. These results support our belief that high frequency observations and analysis using DIA, that MOA is currently carrying out, can detect high magnification microlensing events very efficiently, even with a small telescope.     \\vspace{5cm}"
    },
    "0207/astro-ph0207097_arXiv.txt": {
        "abstract": "With a redshift of $z \\approx 1.7$, SN 1997ff is the most distant type Ia supernova discovered so far. This SN is close to several bright, $z = 0.6-0.9$ galaxies, and we consider the effects of lensing by those objects on the magnitude of SN 1997ff. We estimate their velocity dispersions using the Tully-Fisher and Faber-Jackson relations corrected for evolution effects, and calculate, applying the multiple-plane lensing formalism, that SN 1997ff is magnified by $0.34 \\pm 0.12$ mag. Due to the spatial configuration of the foreground galaxies, the shear from individual lenses partially cancels out, and the total distortion induced on the host galaxy is considerably smaller than that produced by a single lens having the same magnification.  After correction for lensing, the revised distance to SN 1997ff is $m-M = 45.5$ mag, which improves the agreement with the $\\Omega_M =0.35,\\Omega_\\Lambda = 0.65$ cosmology expected from lower-redshift SNe~Ia, and is inconsistent at the $\\sim 3\\sigma$ confidence level with a uniform gray dust model or a simple evolution model. ",
        "introduction": "One of the strongest lines of evidence for an accelerating universe is the observation that \\( z \\approx 0.5 \\) type Ia supernovae (SNe~Ia) are \\( \\sim 0.25 \\) mag fainter than predicted by an \\( \\Omega_\\Lambda =0 \\) cosmology (Riess et al. 1998; Perlmutter et al. 1999; Riess 2000). However, there are several effects such as gray dust (Aguirre 1999) or luminosity evolution (Drell, Loredo, \\& Wasserman 2000) that could mimic the effect of a cosmological constant.  Riess et al. (2001, hereafter R2001) have recently concluded that SN 1997ff, one of the two SNe discovered by Gilliland, Nugent, \\& Phillips (1999) in the Hubble Deep Field North (HDFN; Williams et al. 1996), has a redshift of \\( z \\approx 1.7, \\) making it the highest redshift probable SN~Ia identified so far. (The SN~Ia classification is based on the nature of the host galaxy, an evolved, red elliptical; also the observed colors and temporal evolution of SN 1997ff are most consistent with those of SNe~Ia.)  This object is particularly interesting because at \\( z > 1 \\) the preceding epoch of deceleration causes objects to appear brighter than a persistence of astrophysical effects invoked as alternative hypotheses to explain the SN~Ia faintness at \\( z \\approx 0.5 \\). R2001 measured a distance modulus for SN 1997ff which is more than a magnitude brighter ($1.25 \\pm 0.32$ mag brighter assuming the spectroscopic indication of $z = 1.755$) than the predictions from uniform gray dust or toy evolution models, disfavoring these alternatives to a cosmological constant.  Lewis \\& Ibata (2001) stated that SN 1997ff could be magnified by a factor of $ \\sim 1.4$ by two ``elliptical systems\" close to its line of sight.  R2001 anticipated a magnification of $0.3$ mag from these same sources, in good agreement with Lewis \\& Ibata.  As we see below, these are in fact two spiral galaxies (\\#709 and \\#720 in Table 1) which magnify SN 1997ff by only $\\sim 0.2$ mag, about half of the value estimated by Lewis \\& Ibata, and close to the upper limit inferred by R2001 for these two galaxies from the induced shear on the host-galaxy shape and orientation.  M{\\\" o}rtsell, Gunnarsson, \\& Goobar (2001) also tackled this problem using rest-frame $B$-band and ultraviolet luminosities to infer the masses of the lenses. After performing numerous ray-tracing simulations, they concluded that the range of possible magnifications was too large to obtain a firm value for the distance modulus of SN 1997ff based on our current knowledge of the lensing galaxies.  Here we use the superb photometric and positional information provided by the HDFN to estimate velocity dispersions for the galaxies lensing SN 1997ff using the Tully-Fisher (TF) and Faber-Jackson relations, and correcting for evolution effects.  With this information, we calculate, using the multiple-plane lensing formalism (Schneider, Ehlers, \\& Falco 1992, hereafter SEF), the magnification of SN 1997ff and the shear induced on the host galaxy, which is fully compatible with its expected intrinsic shape.  The structure of the paper is as follows. Section 2 describes the environment of SN 1997ff, and identifies the main lensing foreground galaxies. The multiple-plane lens formalism is discussed in \\S~3, while \\S~4 presents our estimates of the magnification and shear. Section 5 summarizes our main results and conclusions. ",
        "conclusions": "The magnification of SN 1997ff, \\( \\mu \\approx 1.4 \\), is high but compatible with the expectation from simulations. For instance, Barber et al. (2000) find \\(\\mu >1.4 \\) for \\( 5\\% \\) of the lines of sight to \\( z=2 \\). However, we would not attribute such magnification to a selection effect; the dispersion of observed magnitudes due to lensing is only $\\sim$0.2 mag, no larger than the intrinsic dispersion of SNe~Ia, and it has been shown to cause negligible selection bias (Riess et al. 1998).   It is unlikely that lensing by large-scale structure could significantly affect our results. Dalal et al. (2002) show that most of the scatter in the magnification comes from scales of less than $1'$, corresponding to the range considered here. The expected residual root-mean-square on larger scales is $\\sim 0.07$ mag, smaller than our uncertainty. It is noteworthy that by combining an ``individual lenses'' approach like the one followed here, with a shear analysis which estimates the large-scale magnification as proposed by Dalal et al., it could be possible to significantly reduce the uncertainty in supernova distances due to lensing. Dark haloes (without optically visible galaxies) may be a possible problem with this approach, but so far there is very little, if any, evidence for their existence. The major source of error in our calculations is the dispersion in the TF and Faber-Jackson relations at high redshift, but in the future improved calibrations using near-infrared data should reduce this scatter considerably.  The amount of magnification we expect for SN 1997ff is similar to the values determined by Lewis \\& Ibata (2001) and R2001, and roughly agrees with the results of M{\\\" o}rtsell et al. (2001) for realistic values of their input parameters. However, our use of the multiple lens-plane formalism, individual galaxy K-corrections, and TF and Faber-Jackson relationships, corrected for evolution, provides an estimate that is more accurate and robust than previous work.  In summary, SN 1997ff, with \\( z\\approx 1.7 \\), is the highest redshift SN~Ia discovered so far. It is shown here that gravitational lensing by nearby foreground galaxies is likely to have magnified SN 1997ff by \\( 0.34 \\pm 0.12\\) mag. Due to the spatial configuration of the foreground galaxies, the shear from the individual lenses partially cancels out, and the total distortion induced on the host galaxy is less than half of that produced by a single lens with the same magnification.  After correcting for lensing, and assuming $z_s = 1.755$, the distance modulus to SN 1997ff is $m-M = 45.49 \\pm 0.34$ mag, in better agreement with the $\\Omega_M =0.35,\\Omega_\\Lambda =0.65$ cosmology expected from lower-redshift SNe~Ia, but inconsistent with uniform gray dust or simple evolution models as an explanation for the dimming of $z < 1$ SNe~Ia at the $3\\sigma$ confidence level. It is noteworthy that our new result also agrees within $1.8\\sigma$ with the predictions of conformal gravity cosmological models (Mannheim 2001).  Since the study of high-z SNe~Ia cannot avoid the effects of gravitational lensing, future use of SNe to constrain cosmological parameters will need to consider and ultimately contend with the effects of lensing."
    },
    "0207/astro-ph0207574_arXiv.txt": {
        "abstract": "Recent results for Galactic and Magellanic Cloud Wolf-Rayet stars are summarised based on line blanketed, clumped model atmospheres together with UV, optical and IR spectroscopy. The trend towards earlier WN and WC spectral types with decreasing metallicity is explained via the sensitivity of classification diagnostics to abundance/wind density, such that WR mass-loss rates are metallicity dependent. Pre-supernovae masses for WC stars are determined, in reasonable agreement with CO-cores of recent Type-Ic SN. ",
        "introduction": "This article will focus on recent determinations of physical parameters for Galactic and Magellanic Cloud WR stars from UV to IR diagnostics. Observationally, the {\\it Far-Ultraviolet Spectroscopic Explorer (FUSE)} has provided an impressive database of $\\lambda$912--1187\\AA\\ spectroscopy for Galactic and Magellanic Cloud WR stars (see Willis et al. these proc.) to  supplement previous UV datasets obtained with the {\\it International Ultraviolet Explorer (IUE)} and {\\it Hubble Space Telescope (HST)}. At longer wavelengths, {\\it Infrared Space Observatory (ISO)} observations of WR stars have now been analysed. Much of the recent observational progress with Wolf-Rayet stars has involved the acquisition of high quality optical spectroscopy for individual stars beyond the Magellanic Clouds (Drissen, these proc.), plus X-ray spectroscopy of single and binary WR stars (e.g. Skinner et al. 2002) neither of which topics will be discussed here.    Theoretical developments in the last few years have been more steady, with the (laborious) implementation of line blanketing into codes by elements other than CNO and Fe, which had already been discussed at the last hot star beach symposium, IAU Symp. 193. The major change has been the widespread use of such codes to analyse individual stars within a range of galaxies. At present, there are a variety of model atmosphere codes which consider sphericity and line blanketing and fall into two main types, outlined below.  CMFGEN (Hillier \\& Miller 1998) and the Gr\\\"afener et al. (2002) code make use of variants of the super-level approach to incorporate the effect of tens of thousands metal lines on the atmospheric structure within the radiative transfer code. CMFGEN can now simultaneously consider blanketing by individual ions of up to 30 elements, including C, N, O, Ne, Si, S, Ar, Ca, Fe and Ni (Hillier, these proc.), whilst Gr\\\"afener et al. consider CNO, Si plus Fe-group elements (Sc to Ni) grouped together in a single generic atom. This approach suffers the least number of approximations, but remains computationally demanding. Recent test calculations for early-type WC stars show (perhaps surprisingly!) good consistency between these two codes, including ionizing fluxes.  Alternatively, Schmutz (1997) and ISA-wind (de Koter et al. 1993, 1997) use separate codes to solve the radiative transfer problem and line blanketing, the latter making use of Monte Carlo techniques. The method had the great computational advantage that complete intensity-weighed effective opacity factors can be calculated separately from the transfer problem. On the negative side, the ionization and excitation equilibrium of metal species is approximate, dictating which lines are efficient at capturing photons for each point in the atmosphere. Test calculations for a late-type WN star between CMFGEN and ISA-wind show excellent agreement in derived stellar parameters, but rather poorer agreement for ionizing fluxes (Crowther et al. 1999). \\begin{table} \\caption{\\small Recent revisions in the derived stellar parameters (clumped in bold with $f$=0.1) for HD~165763 (WR111, WC5) due to the incorporation of metal line blanketing.\\label{table1}} \\begin{center} \\begin{tabular}{lcllcl} \\hline $T_{\\ast}$ & $\\log L$ & $\\log \\dot{M}$ &Elements & Blank?& Reference \\\\ kK         &  $L_{\\odot}$   & $M_{\\odot}$yr$^{-1}$ &included& &\\\\ \\hline 35         & 4.6                  & $-4.6$ & He & no& Schmutz et al. 1989\\\\ 59         & 5.0                  & $-4.4$ & He, C & no & Hillier 1989\\\\ 90  &  5.3            & {\\it\\bf $-$4.8} & He, C, O, Fe & yes& Hillier \\& Miller 1999 \\\\ 85   & 5.45           & {\\it\\bf $-$4.9} & He, C, O, Fe-group & yes & Gr\\\"{a}fener et al. 2002\\\\ \\hline \\end{tabular} \\end{center} \\end{table}    The main effect of blanketing is to re-distribute extreme UV flux to longer wavelengths, reducing the ionization balance in the atmosphere, such that higher stellar temperatures (and luminosities) are required to match observed line profile diagnostics relative to unblanketed studies. This is illustrated in Table~1 for the prototypical Galactic early-type WC star HD~165763 (WR111) whose derived stellar luminosity has increased by a factor of 5--7 over the past decade. Differences in luminosities for HD~165763 between the recent studies of Hillier \\& Miller (1999) and Gr\\\"afener et al. (2002) most likely result from the inclusion of additional blanketing elements, which CMFGEN now routinely handles. Recent revisions to temperatures and luminosities of O supergiants have acted in the opposite sense, relative to previous plane-parallel unblanketed model analyses, such that common techniques are now employed throughout for O and WR stars (e.g. Crowther et al. 2002ab).  Clumping is now routinely, albeit approximately, handled in WR model atmospheric studies via radial dependent volume filling factors, $f$. Constraints on $f$ can be obtained from comparisons between red electron scattering wings and observations (e.g. Hillier 1991), although exact determinations generally prove elusive due to line blending, particularly in WC stars. Generally, $f\\sim$0.05--0.25 provide reasonable matches to observed line profiles (e.g. Hamann \\& Koesterke 1998), such that global mass-loss rates are reduced by a factor of $1/\\sqrt{f} \\sim 2-4$ relative to smooth models. The majority of line profiles behave rather insensitively provided $\\dot{M}/\\sqrt{f}$ remains constant with some exceptions (e.g. Herald et al. 2001). In WC stars, the line centre of UV resonance lines of C\\,{\\sc iii-iv} reacts to changes in $f$, providing additional constraints on the filling factor (Crowther et al. 2002a). ",
        "conclusions": ""
    },
    "0207/astro-ph0207268_arXiv.txt": {
        "abstract": "{We investigate how central black holes (BHs) in galactic dark halos could affect strong gravitational lensing. The distribution of integral lensing probability  with image separations are calculated for quasars of redshift 1.5 by foreground dark matter halos. The mass density of dark halos is taken to be the Navarro-Frenk-White (NFW) profile such that, when the mass of a halo is less than $10^{14} M_{\\sun}$, its central black holes or a bulge is included as a point mass. The relationship between the masses $M_{\\bullet}$ of supermassive black holes and the total gravitational mass $M_{\\mathrm{DM}}$ of their host galaxy is adopted from the most recent literature. Only a flat $\\Lambda$CDM model is considered here. It is shown that, while a single black hole for each galaxy contributes considerable but not sufficient lensing probabilities at small image separations compared with those without black holes, the bulges (which are about $100$--$1000$ times larger in mass than a typical black hole) would definitely contribute enough probability at small image separations, although it gives too high probabilities at large separation angles compared with lensing observations. ",
        "introduction": "Cold Dark Matter (CDM) has become the standard theory of cosmological structure formation. The $\\Lambda$CDM variant of CDM with $\\Omega_m=1-\\Omega_{\\Lambda}\\approx 0.3$ appears to be in good agreement with the available data on large scales (Primack \\cite{primack}). On smaller (sub-galactic) scales, there seem to be various discrepancies, such: N-body CDM simulations which give cuspy halos with divergent profiles towards the center (Navarro, Frenk and White \\cite{nfw96}, \\cite{nfw97}, NFW hereafter); bar stability in high surface brightness spiral galaxies which also demands low-density cores; CDM models which yield an excess of small scale structures; formation of disk galaxy angular momentum, which is much too small in galaxy simulations. Issues that have arisen on smaller scales have prompted people to propose a wide variety of alternatives to CDM, such as warm dark matter (WDM) and self-interacting dark matter (SIDM). Now that problems arise from galaxy-size halos and centers of all dark matter halos, high-resolution simulations and observations are the final criterion. Recent highest-resolution simulations appear to be consistent with NFW (Klypin, \\cite{klypin}; Power et al. \\cite{power}) until scales smaller than about 1 kpc. Meanwhile, a large set of high-resolution optical rotation curves has recently been analyzed for low surface brightness (LSB) galaxies. One can also conclude that the NFW profile is a good fit down to about 1 kpc. Although further simulations and observations, including measurement of CO rotation curves (Bolatto et al., \\cite{bolat}), may help to clarify the nature of the dark matter, it now appears that WDM and SIDM are both probably ruled out, while the small-scale predictions of $\\Lambda$CDM may be in better agreement with the latest data than appeared to be the case as recently as a year ago.  In addition to direct simulations and observations, gravitational lensing provides another powerful probe of mass distribution in the universe. Since mass within small scales only deflect light rays slightly, it is difficult to extract mass information from a single lensing event, and thus statistical gravitational lensing is needed even for ``strong\" gravitational lensing of small halos(Turner, Ostriker \\& Gott \\cite{turner}; Narayan \\& White \\cite{narayan}; Cen et al. \\cite{cen}; Kochanek \\cite{kochanek}; Wambsganss et al. \\cite{wambs95}; Wambsganss, Cen, \\& Ostriker \\cite{wambs98}; Porciani \\& Madau \\cite{porci}; Keeton \\& Madau \\cite{keeton}). Li \\& Ostriker (\\cite{li}) first used the semi-analytical approach to analyze gravitational lensing of remote quasars by foreground dark halos in various cold dark matter cosmologies. The mass function of dark halos they used is alternatively given by singular isothermal sphere (SIS), the NFW profile, or the generalized NFW profile. They found that none of these models can completely explain the current observations: the SIS models predict too many large splitting lenses, while the NFW models predict too few small splitting lenses, so they proposed that there must be at least two populations of halos in the universe: small mass halos with a steep inner density slope and large mass halos with a shallow inner density slope. The author conclude that a combination of SIS and NFW halos can reasonably reproduce the current observations. Similarly, Sarbu et al. (\\cite{sarbu}) investigated the statistics of gravitational lenses in flat, low-density cosmological models with different cosmic equations of state $\\omega$. It was found that COBE-normalized models with $\\omega > -0.4$ produce too few arcsecond-scale lenses in comparison with the JVAS/CLASS radio survey, a result that is consistent with other observational constraints on $\\omega$.  When attention is attracted to  alternatives of CDM dark matter density profile at small scales, another kind of dark matter --- super-massive black holes in the centers of most galactic halos is forgotten or ignored in this case, although the idea of detecting supermassive compact objects by their gravitational lensing effects was proposed very early (Press \\& Gunn \\cite{press73}, Wilkinson et al. \\cite{wilki}) and the lensing effects of Schwarzschild black holes in the strong field regime have been discussed in detail (e.g., Virbhadra \\& Ellis \\cite{virbh}; Frittelli, Kling \\& Newman \\cite{fritt}; Bozza et al. \\cite{bozza}). On the other hand, cosmological voids can form directly after the collapse of extremely large wavelength perturbations into low-density black holes or cosmological black holes; such black holes can also be detected through their  weak and strong lensing effects (Stornaiolo \\cite{storn}). The observational evidence presented so far suggests the ubiquity of BHs in the nuclei of all bright galaxies, regardless of their activity, and BH masses correlate with masses and luminosities of the host spheroids and, more tightly, with stellar velocity dispersions (Magorrian et al. \\cite{magor}; Ferrarese \\& Merritt \\cite{ferra}; Ravindranath et al. \\cite{ravin}; Merritt \\& Ferrarese \\cite{merria}, \\cite{merrib}; Wandel \\cite{wandel}; Sarzi et al. \\cite{sarzi}). Most recent high-resolution observational data gives $M_{\\bullet}/M_\\mathrm{bulge}\\approx 10^{-3}$(Merritt \\& Ferrarese \\cite{merric}). Ferrarese (\\cite{ferra02}) further gave the relation between masses $M_{\\bullet}$ of supermassive black holes and the total gravitational mass of the dark matter halo in which they presumably formed \\begin{equation} \\frac{M_{\\bullet}}{10^8M_{\\sun}}\\sim 0.046\\left(\\frac{M_\\mathrm{DM}}{10^{12}M_{\\sun}}\\right)^{1.57}. \\label{bd} \\end{equation} In this paper, we investigate the contributions of galactic central black holes to lensing probabilities at small image separations. Since $\\Lambda$CDM cosmology and NFW profile are in good agreement with the available data of structure formation on almost all scales as mentioned above, we only chose these two models respectively as cosmology and mass density function in our calculations. We model the lenses as a population of dark matter halos with an improved version of the Press-Schechter (\\cite{press74}, PS) mass distribution function, and central BHs are considered for galaxy-size halos.  The paper is organized as follows: the lensing equation is given in Sect.~\\ref{s2}, lensing probabilities are calculated in Sect.~\\ref{s3}, and discussion and conclusions are provided in Sect.~\\ref{s4}. ",
        "conclusions": "\\label{s4} Our numerical results for lensing probability with image separations larger than $\\Delta\\theta$ in five different cases are shown in Fig.~\\ref{fig2}. In all cases, lensing probabilities keep nearly constant until $\\Delta\\theta\\sim 0.1$ arc seconds, and obvious dropdown takes place at about $1$ arc second if central black holes are included, which, of course, does not mean that the main lensing events have image separations larger than $0.1$ arc seconds. As a matter of fact, in the NFW case (without galactic central black holes, the full line in Fig.~\\ref{fig2}), the lensing probability drops quite slowly in the whole range of image separations: $\\Delta\\theta\\sim$ 0---10 arc seconds; such a tendency would extend even to $30$ arc seconds if it is plotted beyond this range, which implies a uniform distribution of lensing probability for its $log$ value among image separations. However, note that in the single black hole case (dotted line), the lensing probability drops to the same value of NFW at 2 arc seconds, which gives the influence range of a single black hole. In the range of $0\\sim 0.1$ arc seconds, the lensing probability for the single black hole case is about 3 times that for NFW. So, clearly, the contributions from central black holes cannot be omitted, although such contributions alone are indeed not enough to explain the observational data. As we have pointed out, there is always more than one black hole in a galactic bulge, and the collector of black holes would make a bulge itself `act like' a black hole. On the other hand, not all the mass of a bulge is concentrated in black holes, so if we treat a whole bulge as an extreme black hole, such a model would produce too many lenses at image separations larger than 3 arc seconds compared with the JVAS/CLASS radio survey. As mentioned above, we have sufficient reason to tune the fraction of a bulge mass to produce a `right' profile of lensing probability at larger image separations required by observational data, but this `sufficient reason' seems not make us produce sufficient lensing probabilities at smaller image separations, as shown by the dash-dot line in Fig.~\\ref{fig2}.  However, we can attribute sufficient lens events at small image separations to galactic central black holes or the bulge. On the one hand, since this paper focuses on whether galactic central black holes would contribute considerably to the lensing probability, we have not considered the effect of magnification bias, which would increase the final result provided here at all image separations. On the other hand, we have used an improved version of the PS halo mass function but not the `best' version. The shape of the mass function predicted by standard PS theory (the improved version) is in reasonable agreement with what is measured in numerical simulations of hierarchical clustering from Gaussian initial conditions only for massive halos,; less massive halos are more strongly clustered or less anti-biased than the standard PS predicted. Sheth \\& Tormen (\\cite{st}, ST) proposed a model that provides a reasonably good fit to the bias relation of less massive haloes as well as to that of massive halos. Note that central black holes are only found in galactic bulges,; ST's correction for mass function in the range of less massive halos would definitely change the lensing probability  discussed in this paper. Also note that one of the two images produced by a galactic central black hole is close to the lens center and very faint; however, VLBI experiments can detect its existence (Hirabayashi \\cite{hirab}; Ulvestad \\cite{ulves}), and further radio lensing surveys would have the ability to identify high flux density ratio of the two images. How and to what extent lensing magnification bias, flux ratio and modified PS mass function may change the final result will be discussed in another paper."
    },
    "0207/astro-ph0207212_arXiv.txt": {
        "abstract": "An 8 m successor to the Hubble Space Telescope (HST) would make incredible gains in the study of stellar populations, especially in the Local Group.  If diffraction-limited at 0.5 $\\mu$m, the ``Next HST'' could produce high-resolution imaging of faint sources over a wide field in 1 percent of the time needed with the HST.  With these capabilities, photometry of the ancient main sequence could be obtained for many sight-lines through Local Group galaxies, thus determining directly the ages of their structures and providing a formation history for the Local Group populations. ",
        "introduction": "One of the primary quests of observational astronomy is measuring the formation history of giant galaxies. Recently, renewed interest in formation via accretion of dwarf galaxies has been sparked by the discoveries of the Sagittarius dwarf galaxy falling into the Milky Way (Ibata et al. 1994) and of a tidal stream in the Andromeda halo (Ibata et al. 2001). With large investments of Hubble Space Telescope (HST) time, color-magnitude diagrams that reach the ancient main sequence can be constructed for selected fields in galaxies of the Local Group, thus providing accurate ages for their structures via the same techniques traditionally used to date the populations in Galactic globular clusters.  Unfortunately, these studies are limited by the long integration times needed to reach the main sequence at the distance of Andromeda (the nearest spiral to our own), and by the stellar crowding that can be addressed at the HST resolution. However, an 8 m optical-UV space telescope, diffraction-limited at 0.5 $\\mu$m, would crack this field wide open, because of a simple concept that is often overlooked in discussions of an HST successor: The time to reach a given signal-to-noise for background-dominated photometry of point sources scales as aperture to the fourth power, for a telescope that is diffraction-limited at a given wavelength. ",
        "conclusions": "For background-dominated photometry of stellar populations, an 8 m version of the HST would be 100 times faster than the current HST, and provide 3 times the resolution.  Although the increased resolution would allow photometry in fields that are currently impossible with the HST, the reduction in exposure time is not an advance that should be treated lightly.  Obviously, for a given field that can be resolved with the HST, one can always argue that requesting one hundred orbits of HST time would be easier than building an 8 m HST and observing for one orbit.  However, a factor of 100 in exposure time allows one to resolve the main sequence in dozens of sight-lines throughout the Local Group, a program that is simply impossible with the HST, no matter what fraction of its remaining lifetime is devoted to the problem.  Such a program would allow the direct age determination of structures (disks, bulges, halos) in Local Group galaxies, providing a formation history for these galaxies and their components."
    },
    "0207/astro-ph0207448_arXiv.txt": {
        "abstract": "We investigate a new implementation of the Smoothed Particle Hydrodynamics technique (SPH) designed to improve the realism with which galaxy formation can be simulated. In situations where cooling leads to the coexistence of phases of very different density and temperature, our method substantially reduces artificial overcooling near phase boundaries, prevents the exclusion of hot gas from the vicinity of cold ``clouds'', and allows relative motion of the two phases at each point. We demonstrate the numerical stability of our scheme in the presence of extremely steep density and temperature gradients, as well as in strong accretion shocks and cooling flows. In addition, we present new implementations of star formation and feedback which simulate the effect of energy injection into multiphase gas more successfully than previous schemes.  Our feedback recipes deposit thermal energy separately in cold dense gas and hot diffuse gas, and can explicitly reinject cold gas into the hot phase. They make it possible to damp star formation effectively, to reheat cold gas, and to drive outflows into the galaxy halo and beyond. We show feedback effects to be strongest in small mass objects where much of the gas can be expelled. After idealised tests, we carry out a first low resolution study of galaxy formation in a $\\Lambda$CDM universe. Feedback results in substantial and mass-dependent reductions in the total baryonic mass gathered onto the final object as well as in significant modulation of the star formation history. ",
        "introduction": "A detailed understanding of galaxy formation in cold dark matter universes remains a primary goal of modern astrophysics. Whereas on large scales the clustering of matter is determined almost solely by gravitational forces, a large number of other physical processes contribute to the dynamics on the scales relevant to galaxy formation. In order to gain insight into this problem, it is important to develop numerical methods which can reliably represent these physical processes, many of which occur on scales too small to be resolved by the simulations. The aim of this work is to describe a set of numerical tools that can be used to simulate galaxy formation within popular CDM models.  Smoothed Particle Hydrodynamics or SPH (Gingold \\& Moneghan 1977; Lucy 1977) is a particle-based technique for solving gas-dynamics which is often applied to astrophysical problems. This scheme is fundamentally Lagrangian, it can be easily combined with gravity solvers that use tree structures and it lends itself readily to the wide range of densities in galaxy formation problems. However, standard implementations of SPH have limited ability to resolve steep density gradients, and a number of numerical problems occur when particles are close to a region of very different density. These arise because the usual formulation of SPH assumes that the density gradient across the smoothing kernel of each particle is small. This is not true in many situations in which SPH is commonly used. As a result, low mass clumps of dense gas artificially ``evaporate'', hot diffuse gas is prevented from coexisting with dense ``clouds'' and radiative cooling is artificially enhanced in the diffuse gas which lies near such clouds. In most current implementations of SPH the nett result of these effects is excessive accretion onto the cold phase (Pearce \\etal 1999; Ritchie \\& Thomas 2000; Springel \\& Hernquist 2002a).  The method we propose to overcome these limitations is based on a new neighbour search that considers the thermodynamic properties of particles. Our scheme evaluates the appropriate local density more correctly than previous schemes for particles near a phase boundary. At the same time, it performs well in situations where the phases are far from pressure equilibrium, for example in shocks. The method will be outlined in Sec. 2.1.  In Sec 3.2, we compare its stability and convergence to those of the standard SPH implementation using a set of idealised galaxy formation simulations of varying resolution. We will see that the new implementation retains a larger fraction of hot gas, avoiding the artificial overcooling often caused by inappropriate density estimates. The same test problem also demonstrates the numerical stability of the scheme in the presence of steep density (and temperature) gradients as well as in an accretion shock.  As noted above, proper modelling of the formation and evolution of a galaxy requires many physical processes to be considered in addition to the already complex interaction of nonlinear gravitational evolution and dissipative gas dynamics. The interstellar medium (ISM) where gas exists in a wide range of density and temperature states, may be considered as a multiphase system resulting from the interplay of processes such as gravity, hydrodynamics, star formation, stellar photo-heating, shocking by supernovae and stellar winds, cosmic ray and magnetic field dynamics, chemical enrichment and dust formation (Field 1965, Cox \\& Smith 1974, McKee \\& Ostriker 1977, Ferrara \\etal 1995, Efstathiou 2000). Each process introduces its own length and time scales which often differ by orders of magnitude from those of the galaxy as a whole. As a result, a realistic description of the galactic environment is a severe challenge both for theoretical modelling and for numerical simulation.  Simulations of the interstellar medium which attempt to follow all or many of these effects are only possible at the cost of studying a very small portion of a galaxy. The formation of the galaxy as a whole cannot then be considered. A small simulated region makes it possible to resolve the few parsec length scale characteristic of the evolution of individual expanding supernova remnants (see e.g. Rosen \\& Bregman 1995, Wada \\& Norman 2001, Avillez 2000). In order to achieve sufficient resolution, a two dimensional geometry is often adopted and important ingredients, for example magnetic fields and cosmic rays, are often neglected. The results are thus still far from realistic. Our aim is different. We seek to follow the overall evolution of a forming galaxy. We therefore include small-scale processes only insofar as they affect evolution on scales larger than a few hundred parsecs. Such processes are included with an appropriate (and heuristic) ``sub-grid'' model. In particular, we focus our attention on star formation and on feedback from supernovae and stellar winds, adopting simple recipes similar to those commonly used in semi-analytic models for galaxy formation (White \\& Frenk 1991, Kauffmann, White \\& Guiderdoni 1993, Cole \\etal 1994, Somerville \\& Primack 1999).  In recent years, a number of authors have worked along similar lines, coupling smoothed particle hydrodynamics codes with simple star formation and feedback prescriptions in order to study galaxy formation and evolution (Katz \\& Gunn 1991, Katz 1992, Navarro \\& White 1993, 1994, Steinmetz \\& M\\\"uller 1994). Others have implemented similar recipes in grid-based hydrodynamics codes for the same purpose (Cen \\& Ostriker 1992, 1999). This work has shown that, except at extremely low resolution, the implementation of feedback as a localised heat source is ineffective in regulating star formation (Katz 1992). This is because most of the gas heated by supernovae is so dense that it radiates the injected energy immediately without significant effect on the dynamics. As a consequence, nothing prevents more distant and diffuse gas from cooling and adding itself to the rapidly star-forming, dense ISM. Such ``thermal feedback'' thus fails to produce the starburst-driven winds which can drive galactic fountains or blow gas out of weakly bound systems like dwarf galaxies. Although such winds are observed and are thought to play an important role in galaxy formation, a robust numerical scheme which can generate them is still lacking.  {\\it Ad hoc} solutions of varying complexity have been proposed to remedy this weakness of the standard method (Navarro \\& White 1993, Yepes \\etal 1997, Hultman \\&  Pharasyn 1999, Thacker \\& Couchman 2000, Springel \\& Hernquist 2002b). Our own scheme is a further attempt in this direction, based on an explicit separation of the protogalactic gas into diffuse and dense (star-forming) components. Feedback energy is used to heat the two components separately and to convert gas from the dense to the diffuse component. As a result, it becomes possible to cycle material between the two phases and to heat the diffuse phase directly. Our scheme is thus able to regulate star formation and to drive galactic fountains or winds. In our view, it contains fewer arbitrary elements and is more robust than earlier suggestions, and, in addition, it is closely related to the modification of SPH which we propose in order to avoid unphysical effects near phase boundaries.  Our methods for simulating stellar feedback are described in \\S 2.3. Their ability to suppress star formation, to reheat cold gas and to drive outflows from the galactic disk is demonstrated in \\S 3.3, where we also show how feedback has a relatively larger impact in small mass galaxies. Finally, in \\S 3.4 we use our methods for a preliminary study of the formation of galaxies in a $\\Lambda$CDM universe. Despite their relatively low numerical resolution, these simulations test our schemes in the more realistic context of hierarchical aggregation in an expanding universe. In this situation also we find that feedback suppresses star formation much more effectively than would be inferred from simulations employing more standard SPH methods. ",
        "conclusions": "In this paper we have proposed and tested two modifications of the standard algorithms used for SPH simulations of the formation of galaxies.  The first (MSPH) is designed to reduce artifacts which occur in the common (and often poorly resolved) configuration of cold, dense gas clouds embedded in a hot diffuse halo. The second (MFB) is a new implementation of feedback which allows supernova energy to be channeled effectively into the heating of diffuse gas and the evaporation of cold clouds.  When strong density jumps are absent, for example in most non-radiative problems, our MSPH scheme reduces to a standard SPH algorithm. In the presence of cooling a multiphase structure can arise, and our scheme then eliminates the artificial overcooling discussed by Pearce \\etal (1999); particles in the hot phase which happen to lie near a clump of cold gas have their density, and thus their radiative cooling rate, substantially overestimated by the standard SPH formula. In our scheme such cold, dense neighbours are not considered when calculating densities for particles in the diffuse phase. This modification also allows diffuse gas to take up a realistic spatial structure in the presence of an embedded cold component. This is not the case for standard algorithms (see Figs. (2) and (3)). Finally in strongly dynamic situations our MSPH scheme conserves energy and momentum to the same accuracy and is just as stable as standard algorithms when similar timestep criteria are used.  Since the work of Katz (1992) it has been recognized that implementations of feedback which simply inject supernova energy into the thermal reservoir of neighboring gas particles have little effect on the dynamics of SPH simulations; most of the energy is radiated before it can accelerate the gas. Many alternative schemes have been proposed (e.g. Navarro \\& White 1993, Yepes \\etal 1997, Hultman \\& Pharasyn 1999, Thacker \\& Couchman 2000, Springel 2000, Springel \\& Hernquist 2002b) but none is yet accepted as a proper representation of the unresolved ``microphysics''. Our MFB scheme is original in several respects and is designed to facilitate reproducing the observed properties of starbursts, while introducing as few {\\it ad hoc} elements as possible. We use the supernova energy with predefined efficiencies to heat the hot diffuse phase and to evaporate gas from cold clouds into the diffuse phase. As our tests show, this not only allows feedback to regulate star formation, but also generates winds or galactic fountains without dialling in their characteristics ``by hand'' and permits such flows to entrain significant amounts of cold interstellar material.  Our tests have concentrated on the idealised rotating, collapsing sphere of Navarro \\& White (1993) and on the formation of a single isolated galaxy and its environment in a $\\Lambda$CDM universe. We have used relatively small numbers of particles in these experiments both to facilitate testing and because these kinds of algorithms are often used to study galaxy formation within cosmologically ``representative'' regions (e.g. Katz. Weinberg \\& Hernquist 1996, Pearce \\etal 1999; Murali \\etal 2002, Springel \\& Hernquist 2002a,b); a large fraction of the ``galaxies'' then form from fewer than (say) 1000 gas particles. Our tests show that our proposed algorithms are numerically stable, and that for plausible choices of the heating efficiencies they reproduce the main qualitative features of observed star-forming galaxies; self-regulation of star-formation; bursting behaviour in small systems; the generation of fountains and winds with simultaneous inflow and outflow; the entrainment of disk gas by winds. The obvious next steps are to carry out much larger simulations both of the formation of individual galaxies and of representative volumes. The first can study the origin of the spatial, kinematic and chemical structure of galaxies, checking whether more realistic feedback can indeed solve the disk angular momentum problem (Navarro \\& White 1994, Navarro \\& Steinmetz 1997, Weil, Eke \\& Efstathiou 1998; Thacker \\& Couchman 2001). The second can study how galactic winds enrich and structure the intergalactic medium. We are currently pursuing projects in both these directions.  \\newpage"
    },
    "0207/astro-ph0207481_arXiv.txt": {
        "abstract": "Recent observations of excess radiation at extreme ultraviolet and hard X-ray energies straddling the well known thermal soft X-ray emission have provided new tools and puzzles for investigation of the acceleration of nonthermal particles in the intercluster medium of clusters of galaxies.  It is shown that these radiations can be produced by the inverse Compton upscattering of the cosmic microwave background photons by the same population of relativistic electrons that produce the well known diffuse radio radiation via the synchrotron mechanism.  It is shown that the commonly discussed discrepancy between the value of the magnetic field required for the production of these radiation with that obtained from Faraday rotation measures could be resolved by more realistic models and by considerations of observational selection effects.  In a brief discussion of the acceleration process it is argued that the most likely scenario is reacceleration of injected relativistic electrons involving shocks and turbulence.  The seed electrons cannot be the background thermal electrons because of energetic considerations, and a steady state situation may not agree with the details of the observed spectra.  Episodic generation of shocks and turbulence or episodic injection of relativistic electrons is a more likely scenario for acceleration. ",
        "introduction": "The intercluster medium (ICM) of some clusters of galaxies, in addition to the well studied thermal bremssstrahlung (TB) emission in the 2 to 10 keV soft X-ray (SXR) region, shows growing evidence for nonthermal activity.  The first such activity discovered was the diffused radio radiation classified either as relic or halo sources (see review by Giovannini \\& Feretti 2000).  This radiation is due to synchrotron emission by relativistic electrons of Lorentz factor $\\gamma \\sim 10^4$ in a magnetic field of strength $B\\sim \\mu{\\rm G}$.  More recently, radiations bracketing the thermal one have been discovered in form of excess flux at extreme ultraviolet (0.07-0.4 keV; EUV) and hard X-ray (20 to 80 keV; HXR) regions in several clusters by {\\it The Extreme Ultraviolet Explorer} (Lieu et al.  1996, 1999), and by {\\it Beppo}SAX and RXTE (Coma, Fusco-Femiano et al. 1999 and Rephaeli et al.  1999; A2256, Fusco-Femiano et al.  2000; and possibly A2199, Kaastra et al.  1999).  The reality of the EUV signals from some, but not all, of the clusters is still disputed (see the contributions by Bowyer and Bergh\\\"oefer in this proceedings and Kaastra et al. 2002).  Even though the presence of nonthermal electrons in the ICM was established decades ago very little theoretical treatment of the acceleration mechanism was carried out (see e.g.  Schlikeiser, Siervers \\& Thiemann 1987) until the discovery of the EUV and HXR radiations.  Since then there has been numerous discussions of the possible acceleration mechanisms.  Given the meager amount of the data, detailed calculations of the energy sources  and the exact mechanisms of the acceleration  may be premature.  Consequently, I will emphasize the general physical characteristics and not the numerical details of the problem.  Recently I have analyzed the merits and shortcomings of the various radiative processes in detail and described some possible scenarios for the acceleration of the electrons responsible for the nonthermal radiations (Petrosian, 2001, {\\bf P01} for short).  In what follows I summarize these results.  In \\S 2 I will summarized the observations and compare emission mechanisms and describe what can be surmised about the spectrum of the radiating electrons.  In \\S 3 I discuss  several  ways to resolve the difficulties with the value of the magnetic field and in \\S 4 I will discuss possible acceleration scenarios for the production of the required spectrum of the nonthermal electrons.  A brief summary is given in \\S 5. ",
        "conclusions": "The nonthermal activity in the ICM of some galaxies is well established and a population of 0.1 to 10 GeV electrons in a magnetic field of 1-2 G and the bath of the CMB photons can account for the radio and the EUV-HXR emission via synchrotron and inverse Compton processes, respectively.   The most likely acceleration process is stochastic acceleration by a turbulent ICM.  The source of the accelerated electrons cannot be the background ICM plasma.  There must be injection of relativistic electrons, perhaps throughout most of the ICM, which are then subject to energy loss by the Coulomb and IC processes and to reacceleration by turbulence.  This process cannot be time independent and the injection process or the acceleration mechanisms must be episodic producing electron and photon spectra similar to those observed for periods extending $< 10^9$ yr.  Nonthermal bremsstrahlung is not a viable candidate for production of the HXRs but emission of bremsstrahlung photons by 0.1 to 10 GeV electrons may be important contributor in the GeV range where GLAST sensitivity is the highest.  Other processes may also contribute in this range. For example, as described by Blasi in these proceedings, the decay of pions produced by the cosmic ray protons will produce a broad spectrum around 0.1 GeV. Another source of GeV photons is the IC scattering of the optical and SXR photons (instead of CMB photons) by the GeV electrons responsible for radio and HXR radiations. Electrons with Lorentz factors around $10^3$ to $10^4$ will give rise to photons around 1 to 100 MeV and  1 to 100 GeV, respectively. The former should have a total flux equal to the HXR flux times the optical to CMB energy density ratio ($< 10^{-3}$, see Fig. 1 left panel)  or $\\nu f(\\nu)> 10^{-13}$ erg/cm$^2$. This is well above the anticipated threshold of GLAST in this energy range. The scattering against the SXR photon would give a slightly larger flux of photons with energies above few GeV, but here one is in the Klein-Nishina regime which will reduce the flux considerably, specially at higher energies, making detection by GLAST less likely."
    },
    "0207/astro-ph0207162_arXiv.txt": {
        "abstract": "We apply an eclipse mapping technique to observations of the eclipsing magnetic cataclysmic variable HU~Aqr. The observations were made with the S-Cam2 Superconducting Tunnel Junction detector at the WHT in October 2000, providing high signal-to-noise observations with simultaneous spectral and temporal resolution. HU~Aqr was in a bright (high accretion) state ($V=14.7$) and the stream contributes as much to the overall system brightness as the accretion region on the white dwarf. The stream is modelled assuming accretion is occuring onto only one pole of the white dwarf. We find enhanced brightness towards the accretion region from irradiation and interpret enhanced brightness in the threading region, where the ballistic stream is redirected to follow the magnetic field lines of the white dwarf, as magnetic heating from the stream$\\--$field interaction, which is consistent with recent theoretical results. Changes in the stream eclipse profile over one orbital period indicate that the magnetic heating process is unstable. ",
        "introduction": "HU Aqr is a member of the AM Her sub-class of cataclysmic variables, also called polars because of their highly polarised optical emission. Polars are characterised by a white dwarf primary with a strong magnetic field ($\\sim10-240$ MG) and no accretion disc (see Cropper 1990 and Warner 1995 for reviews). The secondary is a main sequence star which fills its Roche lobe, resulting in an overflow of material from the inner Lagrangian point to form an accretion stream. The stream is initially on a ballistic freefall trajectory, until at some radius the magnetic pressure of the white dwarf magnetic field is approximately equal to the ram pressure of the stream. At this point the stream is threaded onto the field lines of the white dwarf and accretes onto a small area near one or both of the magnetic poles.  HU~Aqr is an eclipsing polar with a period of $\\sim$~$125$~mins, lying just below the period gap (see e.g. Howell, Nelson \\& Rappaport 2001). Harrop-Allin, Hakala \\& Cropper (1999b) used the eclipse of the white dwarf primary by the secondary to establish the distribution of brightness along the stream given a number of parameters defining both the geometry of the system and the physical parameters of the model. The stream emission was calculated using model light curves which are optimized by a genetic algorithm, based on a method first employed on photometric eclipse profiles of HU Aqr by Hakala (1995). The technique was developed and improved by Harrop-Allin \\etalc\\ (1999a) and applied to observations of HU~Aqr in both a high and low accretion state (high- and low-mass transfer; Harrop-Allin \\etalc\\ 1999b and Harrop-Allin, Potter \\& Cropper 2001).  This method of eclipse mapping has also been applied to emission-lines (Sohl \\& Wynn 1999; Vrielmann \\& Schwope 2001; Kube, G\\\"{a}nsicke and Beurmann 2000). Kube \\etalc\\ (2000) used HST FOS spectra of (another polar) UZ For to map the accretion stream line emission at C{\\sc iv}\\ $\\lambda$\\ 1550\\AA, and Vrielmann \\& Schwope (2001) used the emission line light curves of H$\\beta$, H$\\gamma$ and He\\ {\\sc ii}\\ $\\lambda$4686\\AA\\ of HU Aqr. Both these line-emission methods use a three-dimensional stream so data from the entire orbit of the system is fitted and the model can reproduce the out-of-eclipse features, such as the pre-eclipse dips (Watson 1995). They can also reproduce the effect of differing brightness distributions between the irradiated and un-irradiated faces of the stream, and this effect is not reproduced in the model of Harrop-Allin \\etalc\\ (1999a and 1999b). However, the Harrop-Allin \\etalc\\ technique has the benefit of being sensitive to the total emission from the stream (not just the line emission). The technique also has the advantage of access to high signal-to-noise ratio data, which is crucial in producing robust fits.  The light curves presented here were obtained using the S-Cam2 Superconducting Tunnel Junction (STJ) camera, developed by the ESA Astrophysics Division at ESTEC. The camera is the second prototype of a new generation of detectors that record the energy as well as the position and time of arrival (to within $\\sim5\\ \\mu$s for this particular detector) of the incident photons (Rando \\etalc\\ 2000). The application of STJs to optical photon counting was first proposed by Perryman, Foden \\& Peacock (1993), and has since been applied to observations of the Crab Pulsar (Perryman \\etalc\\ 1999), the magnetic cataclysmic variable UZ For (Perryman \\etalc\\ 2001) and to quasar spectroscopy (de Bruijne \\etalc\\ 2002).  The detector itself consists of a liquid helium cooled array of $6\\times 6$ pixels, each being $25\\times 25$~$\\mu$m$^2$, corresponding to $0.6 \\times 0.6$ arcsec$^{2}$ per pixel and a field-of-view of about $4 \\times 4$ arcsec$^2$. The pixels are sandwiches of superconducting tantalum with a thin insulating layer in the middle and the whole device is cooled well below the superconductor's critical temperature (about 0.1T$_{c}$). An incident photon then perturbs the device equilibrium and as the energy gap between the ground state and the excited state is only a few meV, a large number of free electrons is created by each photon, this number being proportional to the photon energy. This is in contrast to normal optical CCD semiconductors where the band gap is $\\sim1$~eV, and photon absorption results in typically only one free electron being created.  The S-Cam2 instrument is particularly suited to observations of cataclysmic variables. The high-time resolution and simultaneous observations of spectral and intensity variations is ideal for eclipsing systems with orbital periods of the order of those in cataclysmic variables. The intensity variations over the rapid ingress and egress can be probed directly, and the rapid variations of the intensity of the accretion stream can be used to provide information on the possible mechanisms of emission along the stream. ",
        "conclusions": "We have carried out high signal-to-noise ratio observations of HU Aqr using S-Cam2 on the WHT on two nights. The system was in a high accretion state, and from the single sharp change in the eclipse profile and archive soft X-ray light curves in this state, we infer that matter was accreting at only one pole on the white dwarf. At the onset of the eclipse, the accretion stream is the source of more than half of the optical emission from the system. The system brightness was similar from orbit to orbit, and the eclipse duration remained constant, but the shape of the accretion stream eclipse changed significantly, ending at $\\phi=0.979$, $0.977$ and $0.982$.  We find eclipse durations which are unchanged from past optical studies (Harrop-Allin \\etalc\\, 1999b, Schwope \\etalc\\ 1997), but shorter than those deduced in the soft X-rays by Schwope \\etalc\\ (2001). However, the location of the accretion region is similar to that found by Schwope \\etalc\\ (2001), so that there is no evidence that the optical emission is from higher latitudes on the white dwarf than the soft X-rays. This indicates that the duration of the soft X-ray eclipses may be affected by absorption, as they suggested.  The duration of the egress of the accretion region in the optical is $8$~s, compared to $1.3$~s in soft X-rays (Schwope \\etalc\\ 2001). This indicates clearly that the region emitting cyclotron radiation is extended by a factor of $\\sim 5$ by comparison with the soft X-ray emitting region, which Schwope \\etalc\\ (2001) calculated as subtending an angle of $3^{\\circ}$.  We have found significant changes in the colour of the accretion stream from one eclipse to the next. This indicates that the threading region is hottest in the last of the eclipses by comparison with the previous two (see Section~\\ref{sec:colourvariations}). We have modelled the stream using the technique of Hakala (1995) and Harrop-Allin \\etalc\\ (1999a). This finds that most of the emission originates from two places, the region close to the white dwarf and in the threading region. By comparison with the models of Ferrario \\& Wehrse (1999) this indicates that magnetic heating is required in the threading region. The modelling clearly identifies an increase in brightness in the threading region for both cycles and an enhancement towards the white dwarf for cycle 29993. From this change in brightness of the heated regions in the model streams, and the varying stream eclipse profiles, we suggest that the magnetic heating in the threading region may be unstable. The implications of the highly variable stream trajectory and brightness profile should be recognised in future investigations of the stream properties. A systematic study of series of consecutive eclipses, with the highest possible signal-to-noise ratio, is required to investigate the characteristics of the magnetic heating and stream instabilities."
    },
    "0207/astro-ph0207591_arXiv.txt": {
        "abstract": "The recent discovery of a Gunn--Peterson (GP) trough in the spectrum of the redshift 6.28 SDSS quasar has raised the tantalizing possibility that we have detected the reionization of the universe.  However, a neutral fraction (of hydrogen) as small as 0.1\\% is sufficient to cause the GP trough, hence its detection alone cannot rule out reionization at a much earlier epoch. The Cosmic Microwave Background (CMB) polarization anisotropy offers an alternative way to explore the dark age of the universe. We  show that for most models constrained by the current CMB data and by the discovery of a GP trough (showing that reionization occurred at $z > 6.3$), MAP can detect the reionization signature in the polarization power spectrum. The expected 1-$\\sigma$ error on the measurement of the electron optical depth is around 0.03 with a weak dependence on the value of that optical depth. Such a constraint on the optical depth will allow MAP to achieve a 1-$\\sigma$ error on the amplitude of the primordial power spectrum of 6\\%. MAP with two years (Planck with one year) of observation can  distinguish a model with 50\\% (6 \\%) partial ionization between redshifts of 6.3 and 20 from a model in which hydrogen was completely neutral at redshifts greater than 6.3.  Planck will be able to distinguish between different reionization histories even when they imply the same optical depth to electron scattering for the CMB photons. ",
        "introduction": "\\label{sec:introduction}  How and when the intergalactic medium (IGM) was reionized is one of the long outstanding questions in cosmology, likely holding many clues about the nature of the first generation of light sources and the end of the cosmological ``Dark Age'' \\citep[see][for a review of our current understanding]{barkana01}.  The lack of strong HI absorption (the ``Gunn--Peterson '', GP, trough) in the spectra of high redshift quasars has revealed that the intergalactic medium (IGM) is highly ionized between redshifts $0\\lsim z \\lsim 6$.  On the other hand, the lack of a strong damping by electron scattering of the first acoustic peak in the temperature anisotropy of the cosmic microwave background (CMB) radiation has shown that the universe was neutral between the redshifts $30\\lsim z \\lsim 10^3$.  Together these two sets of data imply that most hydrogen atoms in the universe were reionized during the redshift interval $6\\lsim z \\lsim 30$.  The recent discovery by \\cite{becker01} of the bright quasar SDSS 1030+0524 in the Sloan Digital Sky Survey (SDSS) at redshift $z=6.28$ has, for the first time, revealed a full GP trough, i.e., a spectrum consistent with no flux at a substantial stretch of wavelength shortward of $(1+z)\\lambda_\\alpha=8850$\\AA.  This discovery has raised the tantalizing possibility that we are detecting reionization occurring near redshift $z\\sim 6.3$.  The lack of any detectable flux indeed implies a strong lower limit $x_{\\rm H}\\gsim 0.01$ on the mean mass--weighted neutral fraction of the IGM at $z\\sim 6$ \\citep{fan02, pentericci02}. On the other hand, because of the large opacities of the hydrogen Lyman series, it is considerably difficult to push this method much further. In particular, in order to prove that we are probing into the neutral epoch, one would like to directly infer $x_{\\rm H}\\approx 1$. While in principle it is possible to infer this from the lack of any flux in a high enough S/N spectrum, in practice the required integration times are implausibly long.  In this paper, we discuss an alternative way of probing deeper into the dark ages. CMB polarization anisotropy at large angles is very sensitive to the optical depth to electron scattering for the CMB photons \\citep[and references therein]{basko80,hogan82,zaldarriaga97b,haimanknox99}. In a model with no reionization, the polarization signal at large angles is negligible. However, CMB photons scattering in a reionized medium greatly enhance the polarization signal  -- making it very likely that such a signal (based on present CMB temperature anisotropy data and the GP trough) will be detected by the ongoing CMB satellite experiment MAP. We also show that Planck (and for large optical depths, MAP), will have the power to discriminate between different reionization histories even when they lead to the same optical depth.  The results from CMB polarization experiments will deepen our understanding of the physics of reionization. They will lead to better constraints on the models of reionization and, thereby, tell us much more about the first sources of light, and indeed, the first structures to form. ",
        "conclusions": "The remarkable discovery of Gunn-Peterson trough in the spectrum of the redshift $z=6.28$ quasar has brought into sharp focus the questions of how and when the universe was reionized. While representing a significant new discovery, the power of the hydrogen Gunn--Peterson trough to probe the dark age is limited, because a small neutral hydrogen fraction ($x_{\\rm H} \\sim 10^{-3}$) is enough to give rise to such a trough. CMB polarization is the most promising way to probe deeper into the dark age; our goal in this paper has been to quantify this statement for forthcoming CMB anisotropy experiments.  Contaminating astrophysical (foreground) sources of polarization may be present at significant levels, but there is very little we know about them.  In deriving our results, we have used a maximum of two channels per experiment, with the remainder assumed to be used as foreground monitors.  Foregrounds might be more detrimental than we have assumed; the only way to find out is to do multi--frequency CMB experiments.   Using current data we have shown that MAP will most likely detect the reionization feature in the polarization data. We have shown that MAP can measure the electron optical depth to 0.02-0.03 (1-$\\sigma$), regardless of the value of the optical depth and that this measurement is sufficient to break the degeneracy between the optical depth and the amplitude of primordial density perturbations. Further, MAP (with two years of observation) will be able to distinguish a model with only neutral hydrogen at $z>6.3$ from one that was about 50\\% ionized between $z=6.3$ and 20. Planck, with its much higher projected sensitivity (and one year of observation), will be able to distinguish an ionized hydrogen fraction of about 6\\% between the redshifts of 6.3 and 20 from the model with no ionized hydrogen at $z>6.3$. More strikingly, Planck (and MAP if the optical depth is larger than about 0.2) should be able to differentiate among reionization histories which lead to the {\\em same} optical depth for CMB photons.  These results will enormously deepen our understanding of the physics of reionization."
    },
    "0207/astro-ph0207134_arXiv.txt": {
        "abstract": "We study the effect of inclusion of muons and the muon neutrinos on the phase transition from nuclear to quark matter in a magnetised proto-neutron star and compare our results with those obtained by us without the muons. We find that the inclusion of muons changes slightly the nuclear density at which transition occurs.However the dependence of this transition density on various chemical potentials, temperature and the magnetic field remains quantitatively the same.\\\\ \\\\ PACS Nos: 97.60.Jd; 12.38.Mh; 97.60.Bw ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207302_arXiv.txt": {
        "abstract": "We report here on a sequence of 28 observations of the binary pulsar system \\psrbe{} at four radio frequencies made with the Australia Telescope Compact Array around the time of the 2000 periastron passage.  Observations made on 2000 Sep 1 show that the pulsar's apparent rotation measure (RM) reached a maximum of $-14800 \\pm 1800$~\\rotm{}, some 700 times the value measured away from periastron, and is the largest astrophysical RM measured. This value, combined with the dispersion measure implies a magnetic field in the Be star's wind of 6~mG.  We find that the light curve of the unpulsed emission is similar to that obtained during the 1997 periastron but that differences in detail imply that the emission disc of the Be star is thicker and/or of higher density. The behaviour of the light curve at late times is best modelled by the adiabatic expansion of a synchrotron bubble formed in the pulsar/disc interaction.  The expansion rate of the bubble $\\sim$12~km~s$^{-1}$ is surprisingly low but the derived magnetic field of 1.6~G close to that expected. ",
        "introduction": "\\psr{} was discovered in 1990 in a survey of the Galactic plane for pulsars \\cite{jlm+92} and was subsequently found to be in a 3.4-yr, highly eccentric orbit about a $\\sim$10 M$_\\odot$ Be star, \\be{} \\cite{jml+92}.  It remains the only known radio pulsar with a Be star companion.  The pulsar has a spin period of $\\sim$48~ms, a characteristic age of 0.33 Myr, and a moderate magnetic field of $3.3 \\times 10^{11}$~G.  The dispersion measure (DM) of 146.7~\\pow{cm}{-3}pc yields a distance of 4.5 kpc in the model of Taylor \\& Cordes (1993)\\nocite{tc93}, however scintillation \\cite{mjsn97} and optical \\cite{jml+94} measurements indicate a distance closer to 1.5 kpc.  The pulse profile, shown in Figure~\\ref{fig:pulseprofile}, has two almost equal intensity peaks both of which are highly (and almost orthogonally) polarised, with an average interstellar rotation measure (RM) measured away from periastron of +21 \\rotm{} \\cite{mj95}.  Optical observations \\cite{jml+94} show that \\be{} is of spectral type B2e, and thus has a mass of $\\sim$10 M$_{\\odot}$ and a radius R$_{c}\\sim$6 R$_{\\odot}$. Assuming a pulsar mass of 1.4 M$_{\\odot}$, the implied inclination angle of the binary orbit to the plane of the sky is 36\\degr. The H$\\alpha$ emission line shows that the Be star's emission disc extends to at least 20 R$_{c}$, just inside the pulsar's orbital radius at periastron.  Timing measurements have shown that the disc of the Be star is likely to be highly inclined with respect to the orbital plane \\cite{wjm+98}. \\begin{figure} \\centerline{\\psfig{figure=fig1.eps,width=8cm,angle=-90}} \\caption{Mean pulse profile at 1.4~GHz. Position angle is shown on top, and the total intensity (solid line), linear (dashed line) and circular (dash-dot line) polarisations are shown in the bottom panel.} \\label{fig:pulseprofile} \\end{figure}  The 1997 periastron was observed extensively by Johnston~et~al.~(1999, 2001)\\nocite{jmmc99,jwn+01} and a model for the transient emission was proposed by Ball et al. (1999)\\nocite{bmjs99} following earlier work by Melatos et al. (1995)\\nocite{mjm95}. In brief, unpulsed emission began at \\p{-22} and lasted until at least \\p{+100} (where $\\peri$ denotes the epoch of periastron). The spectral index of the unpulsed emission was in the range $-0.5$ to $-0.7$, a value indicative of synchrotron radiation. Peaks in the light curve at \\p{-10} and \\p{+20} coincide with the pulsar crossing through the emission disc of the Be star. The model therefore proposed that the synchrotron emission was generated in the shock between the relativistic pulsar wind and the outflowing disc of the Be star. Synchrotron losses, set by the magnetic field strength, then determined the decay time of the light curve. The decay time of the emission from the post-periastron disc crossing was substantially shorter than that of the pre-periastron decay.  The model predicted a frequency dependence for the end of the transient phase and verification of this late-time behaviour was a motivation for the 2000 periastron observations. Also, the improvements in the correlator and the data reduction software meant that the pulsar rotation measure (RM) and dispersion measure (DM) could be accurately obtained during this periastron passage. We therefore obtained data over 28 epochs from \\p{-46} to \\p{+113} during the 2000 periastron encounter. In section 2, we describe the observations. In section 3 we detail the results obtained and discuss their implications in section 4. ",
        "conclusions": "A series of observations of the 2000 periastron of the \\psrbe{} system were made at the ATCA at four frequencies. The pulsar binning mode available with the ATCA correlator meant that pulsed and unpulsed emission could be separated and we therefore obtained light curves for the pulsar and unpulsed emission independently. We were also able to obtain RMs and DMs for the pulsar.  The main features in the light curve from the 2000 periastron are similar to the 1997 periastron, providing confirmation of our earlier model. It does appear, however, that the pulsar eclipse lasted longer and the unpulsed emission began and finished later in 2000 than in 1997. This indicates that the disc has become larger and/or denser in 3.5~yr. As we obtained more data at late times than in the previous periastron encounters, we found that adiabatic decay of the synchrotron bubble provided the best fit to the data. The expansion rate of the bubble, $\\sim$12~km~s$^{-1}$, is significantly lower than expected but the magnetic field value of 1.6~G consistent with measurements in other Be stars.  In addition, at \\p{-46} the pulsar had the highest measured astrophysical RM of $-14800 \\pm 1800$ \\rotm{} implying a magnetic field of at least 6~mG in the Be star wind."
    },
    "0207/astro-ph0207464_arXiv.txt": {
        "abstract": "The Fourier transform of a dataset apodised with a window function is known as the Gabor transform. In this paper we extend the Gabor transform formalism to the sphere with the intention of applying it to CMB data analysis. The Gabor coefficients on the sphere known as the pseudo power spectrum is studied for windows of different size. By assuming that the pseudo power spectrum coefficients are Gaussian distributed, we formulate a likelihood ansatz using these as input parameters to estimate the full sky power spectrum from a patch on the sky. Since this likelihood can be calculated quickly without having to invert huge matrices, this allows for fast power spectrum estimation. By using the pseudo power spectrum from several patches on the sky together, the full sky power spectrum can be estimated from full-sky or nearly full-sky observations. ",
        "introduction": "The Cosmic Microwave Background (CMB) is one of our most important sources of information about the early universe \\cite{bond,jungman,review,durrer}. The pattern of the temperature fluctuations in the CMB contains information about a number of cosmological parameters. If the temperature fluctuations are Gaussian as predicted by most models of the early universe, all this information is stored in the angular power spectrum coefficients $C_\\ell$. For this reason, several experiments have been conducted to measure the CMB power spectrum. The COBE satellite discovered the fluctuations in 1992 \\cite{cobe}, and since then several ground based and balloon borne experiments \\cite{boom1,maxima1,boom2,maxima2,dasi1,dasi2} have been made to study the CMB at an ever increasing resolution. As the amount of CMB data from these experiments is rapidly growing, the task of extracting the power spectrum from the data is getting harder.\\\\  Analysing the CMB data from a given experiment consists of several steps as the data consists of several components not belonging to the CMB \\cite{comp1,comp2}. In this paper, we will concentrate on extracting the power spectrum from a CMB map with foregrounds removed. The standard method of extracting the power spectrum from a sky map is the method of maximum likelihood. This method gives the smallest error bars on the power spectrum estimates, but has the drawback that the number of operations needed to perform the estimation, scales as $N_\\mathrm{pix}^3$, where $N_\\mathrm{pix}$ is the number of pixels in the map. For experiments with high resolution the number of pixels can be up to several million and this method becomes infeasible using current computers \\cite{borrill}.\\\\  In \\cite{OhSpergelHinshaw}, it is shown how the likelihood analysis can be speeded up to scale as $N_\\mathrm{pix}^2$ with assumptions about azimuthal symmetry and uncorrelated noise. Another $N_\\mathrm{pix}^{3/2}$ method for large azimuthally symmetric parts of the sky with uncorrelated noise was presented in \\cite{pseudo}. The likelihood problem can also be solved exact in $~N_\\mathrm{pix}^2$ operations with correlated noise for special scanning strategies as demonstrated in \\cite{ringtorus1,ringtorus2}. In \\cite{bond,BJK,bartlett} it is shown how one can approximate the likelihood to speed up the calculations, but still an $~N_\\mathrm{pix}^3$ operation is needed. This has led people to find other estimators than the maximum likelihood estimator to extract the power spectrum. In \\cite{tegmark} an optimal estimator was found but the calculation scales as $N_\\mathrm{pix}^2$ times a huge prefactor. Recently some near optimal estimators have been found which can be calculated in $~N_\\mathrm{pix}^2$ operations \\cite{dore,szapudi,master} The data from the BOOMERANG \\cite{boom1,boom2} experiment was analysed using the MASTER method \\cite{master}. In this method, the power spectrum was extracted by a quadratic estimator based on the pseudo power spectrum (the power spectrum on the cut sky). A similar method was suggested by \\cite{amad} for the Planck surveyor. Here we propose to use the pseudo power spectrum ($\\tilde C_\\ell$) for likelihood estimation. This principle was also used in \\cite{pseudo} but then for large sky coverage so that the correlations between the $\\tilde C_\\ell$ coefficients could be neglected.\\\\  In this paper, we study the effect of {\\it Gabor transforms} on the sphere. Gabor transforms, or windowed Fourier transforms are just Fourier transforms where the function $f(x)$ to be Fourier transformed is multiplied with a {\\it Gabor window} $W(x)$ \\cite{gabor}. In the discrete case $f(x_i)$ can be a data stream. If parts of the data stream are of poor quality or is missing, this can be represented as $W(x_i)f(x_i)$ where the window $W$ is zero where there are missing parts. The window can also be formed so that it smoothes the edges close to the missing parts and in this way avoid ringing in the Fourier spectrum.\\\\  We will study the effect of Gabor transforms on the sphere and use it for fast CMB power spectrum estimation. The Gabor transform in this context is just the multiplication of the CMB sky with a window function before using the spherical harmonic transform to get the Gabor transform coefficients in this case called the {\\it pseudo power spectrum}. The window can be a top-hat to take out certain parts of the sky in the case of limited sky coverage. Another window can be a Gaussian Gabor window for smoothing the transition between the observed and unobserved area of the sky. The Gabor window can also be designed in such a way as to increase signal-to-noise by giving pixels with high signal-to-noise higher significance in the analysis. The use of the windowed Fourier transform was already studied in \\cite{hobson} in the flat-sky approximation. We show that some of their results are also valid on the sphere.\\\\  In the standard likelihood approach of power spectrum estimation, the pixels on the CMB sky or the spherical harmonic coefficients $a_{\\ell m}$ are used as elements in the data vector in which case the correlation matrix will have dimensions of the order $N_\\mathrm{pix}\\times N_\\mathrm{pix}$. A matrix of this size can not be inverted in a reasonable amount of time with current computers. We propose to use the pseudo power spectrum coefficients $\\tilde C_\\ell$ as elements of the data vector in the likelihood. In this case the size of the correlation matrix will at most be $l_\\mathrm{max}\\times l_\\mathrm{max}$ which can be inverted in a few seconds. The most time consuming part is the calculation of the elements of the correlation matrix of pseudo-$C_\\ell$.\\\\  In Section (\\ref{sect:gabor}) we will first describe the one dimensional Gabor transform and then define the Gabor transform on the sphere. We will define the pseudo power spectrum which is just the Gabor coefficients on the CMB sky. The kernel relating the full sky power spectrum and the pseudo power spectrum for a Gaussian and top-hat Gabor window will be discussed. Then in Section (\\ref{sect:lik}) we will use the pseudo power spectrum as input values to a maximum likelihood estimation of the full sky power spectrum. The probability distribution of the pseudo power spectrum coefficients will be assumed Gaussian and we will show that this is a good approximation at high multipoles ($\\ell>100$). Some examples of likelihood estimations of the power spectrum with different noise patterns will be shown. In Section (\\ref{sect:ext}) two extensions of the method will be discussed. First the use of the pseudo power spectrum from different Gabor windows centred at different points on the sphere simultaneously is demonstrated. In this way full-sky or nearly full-sky observations can be analysed. The second extension of the method is the use of Monte Carlo simulations to obtain noise properties in the case where this is faster than using the analytic expression or where the noise is correlated. Finally in Section (\\ref{sect:disc}) the results and further extensions are discussed. ",
        "conclusions": "\\label{sect:disc}  In this paper, we propose to use the spherical harmonic transform of the sky apodised by a window function, or Gabor transform, as a fast and robust tool to estimate the CMB fluctuations power spectrum. It is known that the coupling between modes resulting from the analysis on a cut sky affects the shape of the measured pseudo power spectrum and the statistics of the $C_\\ell$ coefficients (\\cite{pseudo,master} and reference therein). In the case of axisymmetric windows  we can compute analytically (in about $\\ell_\\mathrm{max}^3$ operations) the kernel relating the cut sky power spectrum to the full sky one for a Gaussian and top-hat profile we give an analytical relation between the spectral resolution attainable and the size of the sky window. Studying windows of different sizes, we show that for windows as small as $36$ degrees in radius, the measured power spectrum is undiscernable from the true one for $\\ell$ larger than about 50.  Noting that for large multipoles ($\\ell\\geq100$ for windows with radius larger than $36$ degree) the statistics of the pseudo-$C_\\ell$ coefficients measured in Monte Carlo simulations is close to Gaussian, we suggest the use of the pseudo power spectrum as input data vector in a likelihood estimation. For the first time, we show how the correlation matrix between the pseudo power spectrum coefficients obtained on an axisymmetric window of arbitrary profile can be computed rapidly for any input power spectrum, based on a recurrence relation. The computation of the correlation matrix needs a precomputation (independent of the power spectrum) of $\\ell_\\mathrm{max}N_m(N^\\mathrm{in})^2$ operations and each calculation of the correlation matrix with a given power spectrum takes $(N^\\mathrm{in})^2(N^\\mathrm{bin})^2$ operations, where $N^\\mathrm{in}$ is the number of input pseudo-$C_\\ell$ coefficients used, $N^\\mathrm{bin}$ is the number of estimated $C_\\ell$ bins and $N_m$ is a window dependent factor ($N_m\\approx200$ for a Gaussian window and $N_m\\approx400$ for a top-hat window). The noise correlation matrix can also be computed by recurrence. For axisymmetric noise this is very quick ($N_\\mathrm{pix}\\ell_\\mathrm{max}$ operations). For general non-uniform noise this takes some more time (between $(N^\\mathrm{in})^2\\sqrt{N_\\mathrm{pix}}\\ell_\\mathrm{max}$ and $(N_\\mathrm{in})^2\\sqrt{N_\\mathrm{pix}}\\ell_\\mathrm{max}^2$ operations dependent on the window profile and number of approximations). For a Gaussian window with a sharply varying noise profile and a patch of sky similar in size to the one observed by BOOMERANG (about $2\\%$ of the sky) it takes about a day on one single $500 \\mathrm{MHz}$ processor. This is the computationally heaviest part of the method but this has to be done only once. The inversion of the correlation matrix, which is the leading problem when doing likelihood analysis, is now overcome, as the size of the correlation matrix is so small that inversion is feasible. In the standard likelihood approach, the correlation matrix has dimensions $N_\\mathrm{pix}\\times N_\\mathrm{pix}$ which needs $N_\\mathrm{pix}^3$ operations to be inverted. In our approach, the size of the correlation matrix is $N^\\mathrm{in}\\times N^\\mathrm{in}$ which in our example $N^\\mathrm{in}=200$ is inverted in a few seconds.   By doing Monte Carlo simulations of different experimental settings, we shown that the likelihood estimator is unbiased. The error bars were found using the inverse Fisher matrix and compared to the error bars obtained from Monte Carlo. There was an excellent agreement between the two sets of error bars. In \\cite{master} it was shown that using a Gaussian apodisation suppresses the signal such that the error bars on the estimated power spectrum becomes larger. In this paper we have shown that using a different window than the top-hat window can be important for increasing the signal-to-noise ratio in data with non-uniform noise. We applied a Gaussian window to an observed disc which had the noise level increasing from the centre of the disc and outwards similar to what one can expect around the ecliptic poles in scanning strategies like the ones of MAP and Planck. In this case the Gaussian window has a high value in the centre where signal-to-noise per pixel is high and a low value close to the more noise dominated edges. We shown that for this noise profile using a Gaussian window increased the signal-to-noise ratio significantly over the top-hat window, showing that adapting the window to downweight noisy pixels gives better performance than a simple uniform weighting.  Finally two extensions of the power spectrum estimation method were discussed. First it was shown that for observed areas on the sky which are not axisymmetric, one can cover the area by several axisymmetric patches and make a joint analysis of the pseudo power spectrum coefficients from all the patches. Each of the patches can have a different window in order to optimise signal-to-noise in each patch. This method will be extended in a forthcoming paper where we will discuss the use of the method for analysing MAP and Planck data sets. We also shown that the calculation of the noise correlation matrix can be quicker by Monte Carlo simulations if the number of pixels in the observed area is huge (about $10^6$ pixels but dependent on the window shape). This may also be used in the case of correlated noise. We shown that a few thousand simulations are necessary to get the same accuracy in the power spectrum estimates as when using the analytic formula for the noise correlation matrix.  In \\cite{polarisation} we show that the power spectrum estimation method presented in this paper can easily be extended to polarisation. By extending the data vector in the likelihood to have also the pseudo-$C_\\ell$ from polarisation, one can in a similar way estimate for the temperature and polarisation power spectra jointly."
    },
    "0207/astro-ph0207187_arXiv.txt": {
        "abstract": "GRB~010921 was the first HETE-2 GRB to be localized via its afterglow emission.  The low-redshift of the host galaxy, $z=0.451$, prompted us to undertake intensive multi-color observations with the {\\it Hubble Space Telescope} with the goal of searching for an underlying supernova component.  We do not detect any coincident supernova to a limit 1.34 mag fainter than SN~1998bw at 99.7\\% confidence, making this one of the most sensitive searches for an underlying SN. Analysis of the afterglow data allow us to infer that the GRB was situated behind a net extinction (Milky Way and the host galaxy) of $A_V \\sim 1.8$ mag in the observer frame.  Thus, had it not been for such heavy extinction our data would have allowed us to probe for an underlying SN with brightness approaching those of more typical Type Ib/c supernovae. ",
        "introduction": "\\label{sec:introduction}  Since the discovery of gamma-ray burst (GRB) afterglows there has been growing evidence linking GRBs to massive stars: the host galaxies of GRBs are star-forming galaxies and the position of GRBs appear to trace the blue light of young stars \\citep{bkd02}; some of the host galaxies appear to be dusty with star-formation rates comparable to ultra-luminous infrared galaxies \\citep{bkf01,fbm+01}. On smaller spatial scales, there is growing evidence tying GRBs to regions of high ambient density \\citep{gw01,hys+01} and the so-called dark GRBs arise in or behind regions of high extinction \\citep{dfk+01a,pfg+02}.  However, the most direct evidence linking GRBs to massive stars comes from observations of underlying supernovae (SNe) and X-ray lines.  The presence of X-ray lines would require a significant amount of matter on stellar scales (e.g. \\citealt{pgg+00}), as may be expected in models involving the death of massive stars.  However, to date, these detections (e.g. \\citealt{pgg+00,rwo+02}) have not been made with high significance.  If GRBs do arise from the death of massive stars, then it is reasonable to expect associated SNe.  The GRB-SN link was observationally motivated by two discoveries: the association of GRB~980425 with the peculiar Type Ic SN 1998bw \\citep{gvv+98,kfw+98} and an excess of red light superposed on the rapidly decaying afterglow of GRB~980326 \\citep{bkd+99}.  However, these two discoveries were not conclusive.  The SN association would require GRB~980425 to be extra-ordinarily under-energetic as compared to all other cosmologically located GRBs and the case for GRB 980326 is weakened by the lack of a redshift for the GRB or the host galaxy.  Nonetheless, the two discoveries motivated searches for similar underlying SN components.  As summarized in section~\\ref{sec:conclusions}, suggestions of similar red ``bumps'' in the light curves of various other GRB afterglows have been made (to varying degrees of confidence).  However, there is little dispute that the well-studied red bump in the afterglow of GRB~011121 is most easily explained by an underlying supernova \\citep{bkp+02,gsw+02}. Furthermore, from radio and IR observations of the afterglow \\citep{pbr+02}, there is excellent evidence that the circumburst medium was inhomogeneous with ambient density $\\rho \\propto r^{-2}$, as expected from a massive star progenitor \\citep{cl00}; here, $r$ is the distance from the progenitor.  These developments are in accordance with the expectation of the ``collapsar'' model \\citep{woo93,mw+99}.  In this model, the core of a rotating massive star collapses to a black hole which then accretes matter and drives a relativistic jet.  Internal shocks within this jet first cause bursts of $\\gamma$-rays and then subsequently result in afterglow emission as the jet shocks the ambient medium.  It is important to appreciate that the SN light is primarily powered by radioactive decay of the freshly synthesized $^{56}$Ni whereas the burst of $\\gamma$-rays are powered by the activity of the central engine. In the current generation of collapsar models, there is sufficient flexibility to allow for a large dispersion of $^{56}$Ni and the energy of the engine. Thus, the next phase of understanding the GRB-SN connection\\footnotemark\\footnotetext{A class of models, known as ``supranova'' models, posit a supernova greatly in advance, many months, of the the GRB event \\citep{vs99}. The long delay was physically motivated to explain the X-ray lines as arising from a large spatial region.  The current data (e.g. GRB~011121) do not allow for such long delays.}  will benefit from (and require) observational measures of these parameters.   Motivated thus, we have an ongoing program of searches for SNe in GRB afterglows with the {\\it Hubble Space Telescope} (HST).  Here, we present a systematic search for a SN underlying GRB~010921.  In \\S\\ref{sec:observations} we present our observations and the details of photometry in \\S\\ref{sec:subphot}.  We fit afterglow models and constrain the brightness of an underlying SN in \\S\\ref{sec:discussion}. We then present an overview of previous such efforts and conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   Here we report the search for an underlying SN in the afterglow of GRB~010921. Thanks to the superb photometric stability of HST and the $N\\times(N-1)/2$ subtraction technique, we have been able to trace the light curve of the afterglow of GRB~010921 over two months. The resulting photometry is unbiased by aperture effects that are so prevalent in simple aperture and PSF-fitting photometry. We report two results.  First,  we find a jet break time of 35 days, using only optical data. Second, we find no evidence for an SN.  A SN, if present, must be fainter than SN~1998bw by $> 1.34$ mag at 99.7\\% confidence.  To our knowledge, to date, this is the most stringent limit for an underlying SN associated with a cosmologically located GRB.  As noted in \\S\\ref{sec:introduction}, the collapsar model as currently understood has little power in predicting the dispersion in the amount of $^{56}$Ni synthesized as compared to the energy in relativistic ejecta. Underlying SNe are directly powered by the former whereas the GRB is powered by the latter. Observations are needed to start mapping the distribution in these critical explosion parameters.  Progress can be expected with such observational inputs accompanied by further refinements in the model. Motivated thus, we summarize in Table~\\ref{tab:previous} the status of SN searches for all Table~\\ref{tab:previous} all known GRBs with redshift\\footnotemark\\footnotetext{Beyond a redshift of $\\sim$ 1.2, the distinctive and strong absorption blueward of 4000\\AA\\ is redshifted out of the optical bands.  The higher sensitivity of the optical bands thus favor searches for SNe below this redshift.} less than 1.2.  The most secure case for an SN is that for GRB~011121 \\citep{bkp+02,gsw+02}.  GRB~980326 shows a strong red excess at about a month but unfortunately a redshift is lacking.  GRB~970228 shows a less clear excess but benefits from a known redshift. Stated conservatively, a SN as bright as that of SN~1998bw can be ruled out in GRB~000911. In all cases, save that of GRBs~980326 and 011121, the presence of a host with a magnitude comparable to the brightness of the peak of the SN, makes it difficult to identify an SN component. As noted in \\S\\ref{sec:subphot}, ``bumps'' can arise from host contamination.  Combining HST and ground based measurements (as is the case for GRB~970228) is prone to considerable errors (\\S\\ref{sec:subphot}).  In summary, there is good evidence for an SN comparable in brightness to SN~1998bw in GRB~011121 \\citep{bkp+02}. For GRB~010921, using the HST observations reported here, we constrain any putative underlying SN to be 1.34~mag fainter than SN~1998bw.  In the collapsar framework, this absence could be most readily attributed to the well known dispersion of the peak luminosity of Type Ib/c SNe.  An alternative possibility is that there may be more than one type of progenitor for long duration GRBs. Along these lines we note that \\citet{cl00} claim that some afterglows (e.g. GRB~990123) are incompatible with a $\\rho\\propto r^{-2}$ inhomogeneous circumburst distribution whereas other afterglows (e.g. GRBs~970228 and 970508) are better explained by invoking an inhomogeneous circumburst medium. Progress requires both searches for underlying SNe as well as characterizing the circumburst medium via modeling of the early-time afterglow (e.g. GRB~011121, see \\citealt{pbr+02}).   Finally, we note that the afterglow of GRB~010921 (and any coincident SN) was extincted by $A_V^{\\rm MW} \\approx 0.5$ mag of dust in the foreground, and $A_V^{\\rm host} \\approx 1$ mag of dust in the host galaxy (Table~\\ref{tab:fit}). Thus, in the future, using ACS aboard HST it should be possible to extend SN searches to at least 3 mag fainter than SN~1998bw, at which point it will be possible to detect more typical SNe Ib/c coincident with GRBs."
    },
    "0207/astro-ph0207378_arXiv.txt": {
        "abstract": "In this paper we report on an \\em XMM--Newton\\em~observation of the ultraluminous infrared QSO Mrk~1014.  The X-ray observation reveals a power-law dominated ($\\Gamma \\approx$ 2.2) spectrum with a slight excess in the soft energy range. AGN and starburst emission models fit the soft excess emission equally well, however, the most plausible explanation is an AGN component as the starburst model parameter, temperature and luminosity, appear physically unrealistic. The mean luminosity of Mrk~1014 is about 2 $\\times$ 10$^{44}$ erg s$^{-1}$. We have also observed excess emission at energies greater than 5 keV.  This feature could be attributed to a broadened and redshifted iron complex, but deeper observations are required to constrain its origin.  The light curve shows small scale variability over the $\\sim$11 ks observation.  There is no evidence of intrinsic absorption in Mrk~1014.  The X-ray observations support the notion of an AGN dominated central engine. We establish the need for a longer observation to constrain more precisely the nature of the X-ray components. ",
        "introduction": "Mrk 1014 ($V = 15.7$; $z = 0.163$), with a far-infrared (FIR) luminosity in excess of 10$^{12}$L$_{\\odot}$ (Yun 2001), is one of the brightest members of the class of `warm' ultraluminous infrared galaxies (Sanders et al. 1988a) which from their warm IRAS 25$\\mu$m/60$\\mu$m colors and optical spectra, are believed to host powerful AGN. Sanders et al. (1988b) report on the detection of CO(1--0) emission and suggest that Mrk~1014 may be an important link in the evolution of ultraluminous infrared galaxies into UV-excess quasars.  Spectroscopy using the \\em Infrared Space Observatory (ISO)\\em~adds quantitative information to this picture. The low resolution ISOPHOT-S spectrum presented by Rigopoulou et al. (1999) is dominated by AGN-like continuum emission, with no evidence for the aromatic `PAH' emission typical for starbursts. The ratio of 7.7$\\mu$m PAH feature and local continuum is $<$0.63, well below starburst values of $\\sim$3. The diagnostics proposed by Genzel et al. (1998) and Tran et al. (2001) then suggest an upper limit of $\\approx$25\\% for the starburst contribution to the infrared luminosity of Mrk~1014. Sturm et al. (2002) present ISO-SWS mid-IR fine structure line data for Mrk~1014. The source is at the limit of the the SWS sensitivity, and only tentative detections are obtained for the low excitation [Ne\\,II]12.8$\\mu$m line and the high excitation [O\\,IV]25.9$\\mu$m line, with a ratio consistent with AGN dominance.  \\em Hubble Space Telescope (HST)\\em~NICMOS observations (Scoville et al. 2000) reveal twisting spiral isophotes beneath the dominant QSO nucleus, indicating either a starburst spiral disk or tidal debris. Optical observations of the host galaxy portray a prominent tidal arm with several knots (Surace \\& Sanders 2000; Surace et al. 1998). The tidal arm is composed primarily of intermediate age stars ($\\sim$ 1 Gyr) with very little contribution from older stars (Canalizo \\& Stockton 2000). Nolan et al. (2001) determine that the majority of the flux is associated with an instantaneous burst of star formation at approximately 12~Gyr, however, Canalizo \\& Stockton (2000) find several regions of recent ($\\sim$ 0.2 Gyr) star formation accounting for up to 30\\% of the total luminous mass along the line of sight. Both age estimations are too high to be associated with on-going starburst activity ($\\sim$10 Myr), which produces X-ray emitting hot gas.  Mrk~1014 was observed with \\em XMM--Newton\\em~(Jansen et al. 2001) as part of the guaranteed time programme. The objective of the study was to obtain a high quality X--ray spectrum of an AGN dominated ultraluminous infrared galaxy, and determine if starburst activity is a necessary component to explain the spectral energy distribution.  In the following section we will present the observations and discuss the data reduction. In section 3, we will describe the spectral models and discuss their significance. A description of the light curve will be presented in section 4. In section 5 we will summarize our findings. A value of the Hubble constant of $H_0$=$\\rm 70\\ km\\ s^{-1}\\ Mpc^{-1}$ and a cosmological deceleration parameter of $q_0 = \\rm \\frac{1}{2}$ have been adopted throughout. ",
        "conclusions": "The 0.3--8 keV spectrum of the Mrk~1014 was fitted with a number of models to determine if the soft energy emission above a predominant power-law is driven by an AGN central engine or starburst activity. The \\em XMM-Newton\\em~observation reveals a power-law dominated spectrum which may be up to thirty times more luminous than the soft X-ray excess. The existence of an X-ray starburst is highly unlikely due to the low temperature and high luminosity found in the models. Other evidence, such as a variable light curve and the soft X-ray to FIR flux ratio suggest an AGN powered continuum in Mrk~1014. A longer observation with \\em XMM-Newton\\em~would allow us to better constrain the nature of the soft energy component.  In addition, a longer look will allow us to probe the interesting high energy excess above 5 keV and determine if it is connected to a broadened and redshifted iron complex. Clearly, more multi-frequency observations are required to disentangle and compare the starburst and AGN contributions in different energy bands."
    },
    "0207/hep-ph0207145_arXiv.txt": {
        "abstract": "Within the framework of a five-dimensional brane world with a stabilized radion, we compute the cosmological perturbations generated during inflation and show that the perturbations are a powerful tool to probe the physics of extra dimensions.  While we find that the power spectrum of scalar perturbations is unchanged, we show that the existence of the fifth dimension is imprinted on the spectrum of gravitational waves generated during inflation. In particular, we find that the tensor perturbations receive a correction proportional to $(HR)^2$, where $H$ is the Hubble expansion rate during inflation and $R$ is the size of the extra dimension. We also generalize our findings to the case of several extra dimensions as well as to warped geometries. ",
        "introduction": "The expanding Universe, especially if it underwent a primordial inflationary phase \\cite{review}, represents the most powerful probe of small distance scales at our disposal.  Present-day astronomical length scales were extremely tiny at early epochs and were sensitive to short-distance physics.  This simple observation has recently generated a lot of excitement about the possibility of opening a window on transplanckian or stringy physics in Cosmic Microwave Background (CMB) anisotropies \\cite{tp}.  Unfortunately, in the absence of a quantum theory of gravity, uncontrollable nonlinear effects may dominate at transplanckian distances, and the behavior of the cosmological perturbations and crucial related issues such as the definition of the vacuum remain unknown. This makes it difficult to predict on firm grounds the signatures of transplanckian physics on present-day cosmological scales \\cite{tpf}.  In this paper we will demonstrate that cosmological perturbations generated during inflation may nevertheless provide a powerful probe of another important aspect of many modern theories of particle physics: the existence of extra dimensions.  The presence of extra dimensions is a crucial ingredient in theories explaining the unification of gravity and gauge forces.  A typical example is string theory, where more than three spatial dimensions are necessary for the consistency of the theory. It has recently become clear that extra dimensions may be very large and could even be testable in accelerator experiments.  In theories with $n$ compactified extra dimensions with typical radii $R$, the four-dimensional Planck mass, $M_{P}$, is just a derived quantity, while the fundamental scale is the gravitational mass, $M_*$, of the $(n+4)$-dimensional theory. The mass scale $M_*$ is a free parameter and can range from a TeV to $M_{P}$, with $M_{P}^2\\sim M_*^{n+2}\\, R^n$. The size of extra dimensions can range from macroscopic scales down to Planckian distances.  In general there is a large hierarchy between the size of extra dimensions, $R$, and $M_*^{-1}$, with $R\\gg M_*^{-1}$. This means that perturbations that are currently observable on cosmological scales might have been generated at early times on scales much smaller than the size of extra dimensions, but still on scales larger than the fundamental Planck mass so that the $(4+n)$-dimensional Einstein equations should describe gravity and the behavior of the quantum vacuum is more certain. This provides a unique probe of the physics of extra dimensions without the necessity of dealing with unknown effects at energies larger than $M_*$.  This is particularly relevant in brane-world scenarios where gravity propagates in a higher-dimensional space while our visible Universe is a three-dimensional brane in the bulk of extra dimensions \\cite{Arkani-Hamed:1998rs}.  In this paper we initially assume a five-dimensional world where our visible Universe is a three-dimensional brane located at a given point in the fifth dimension. We consider the simplest possibility that inflation is a brane effect, {\\it i.e.,} it is driven by a scalar field living on our three-dimensional brane, and study the effects of the transdimensional physics on the spectrum of the primordial density perturbations produced during the epoch of inflation.  Our findings indicate that despite the fact that the power spectrum of scalar perturbations remains unchanged, the existence of the fifth dimension is imprinted on the spectrum of gravitational waves generated during inflation. The tensor spectrum receives a correction proportional to $(HR)^2$, where $H$ is the Hubble rate during inflation and $R$ is the size of the extra dimension. Generalizing our results to the case of more than one extra dimension and to warped geometries, we show that the numerical coefficient of the correction term depends upon the details of the spacetime geometry of the extra dimensions.  In four-dimensional single-field models of inflation there exists a consistency relation relating the amplitude of the scalar perturbations, the amplitude of the tensor perturbations, and the tensor spectral index. We compute the correction to such a consistency relation from transdimensional physics. Surprisingly enough, we find that at lowest order in the slow roll expansion, the four-dimensional relation is quite robust and does not suffer corrections from extra-dimensional physics, at least in not the cases addressed in this paper.  Some similar conclusions have been reached in Refs.\\ \\cite{lmw,mwbh} for a particular five-dimensional setup in which the expansion law on the brane has a non-standard expression. Instead, we will focus on the case where the radius of the extra dimension is stabilized, leading to an ordinary Friedmann law. So any effect should be attributed to the non-trivial geometry along the extra dimension, rather than any modified cosmology on the brane. However, the formalism used in Ref.\\ \\cite{lmw} has many similarities with ours.  Our paper is organized as follows. In Section II we study the five-dimensional background with a stabilized radius. In Section III we compute the power spectrum of the tensor modes generated during inflation, while in Section IV we calculate the power spectrum of scalar perturbations.  Section V is devoted to the consistency relation and Section VI to a generalization of our findings to more than one extra dimension and to warped geometries. Finally, in Section VII we draw our conclusions.  The paper also contains an Appendix where we collect the background and perturbed Einstein equations. ",
        "conclusions": "In this paper we have initiated the investigation of the effects of transdimensional physics on the spectrum of the cosmological density perturbations generated during a period of primordial inflation taking place on our visible three-brane. We have shown that the size of the transdimensional effects are of order $(HR)^2$, where $H$ is the Hubble parameter during inflation and $R$ is the typical size of the extra dimensions (or, more precisely, the inverse of the Kaluza--Klein mass gap at zero temperature).  The corrections appear in the power spectrum of the tensor modes. The coefficient of the corrections depends upon the compactification geometry, the number of extra dimensions, and if they are flat or warped.  As we have already stressed in the Introduction, our treatment should be unaffected by (unknown) quantum effects which might arise at distances below $M_*^{-1}$ as long as the size of extra dimensions is larger than $M_*^{-1}$.  Our results may be generalized in different ways. First of all, our set up is the simplest we could imagine: only one extra dimension and inflation taking place on the brane.  One can envisage the possibility of putting the inflaton field responsible for inflation in the bulk. In such a case we expect a different form of corrections. In particular, the power spectrum of scalar perturbations should be modified, thus possibly changing the consistency relation.  One can also relax our working assumption of keeping the radii of extra dimensions fixed. In this case there might be significant corrections to the slope of the power spectra since having a dynamical radion field during inflation amounts to change the Hubble rate during inflation.  In our paper we have also assumed that the energy density $\\rho$ on the brane is smaller than about $ M_{P}^2 / R^2$, or equivalently, that the Hubble radius is larger than the radii of compactification. Deviations from the standard four-dimensional Friedmann law are present in the opposite regime and large deviations from the standard results for the power spectra of density perturbations should appear.  Finally, we have assumed that deep in the ultraviolet regime, at distances much smaller than the horizon length, the initial vacuum is the traditional Bunch--Davies vacuum containing no initial particles in the spectrum.  This is a reasonable assumption at physical momenta $k/a_0$ much larger than $R^{-1}$, but still smaller than the fundamental scale $M_*$. Of course, for momenta $k/a_0\\gg M_*$, unknown quantum effects may take over and change drastically the properties of the vacuum. This would lead to corrections scaling as powers of $H/M_*$ as suggested by the analysis performed in the four-dimensional cases \\cite{tp}. Nevertheless, whenever $M_*R$ is sufficiently large, the computable corrections discussed in this paper dominate over these unknown quantum corrections."
    },
    "0207/astro-ph0207536_arXiv.txt": {
        "abstract": "We investigate the formation of planetesimals via the gravitational instability of solids that have settled to the midplane of a circumstellar disk.  Vertical shear between the gas and a subdisk of solids induces turbulent mixing which inhibits gravitational instability.  Working in the limit of small, well-coupled particles, we find that the mixing becomes ineffective when the surface density ratio of solids to gas exceeds a critical value.  Solids in excess of this precipitation limit can undergo midplane gravitational instability and form planetesimals.  However, this saturation effect typically requires increasing the local ratio of solid to gaseous surface density by factors of two to ten times cosmic abundances, depending on the exact properties of the gas disk.  We discuss existing astrophysical mechanisms for augmenting the ratio of solids to gas in protoplanetary disks by such factors, and investigate a particular process that depends on the radial variations of orbital drift speeds induced by gas drag.  This mechanism can concentrate millimeter sized chondrules to the supercritical surface density in $\\leq {\\rm few} \\times10^6$ years, a suggestive timescale for the disappearance of dusty disks in T Tauri stars.  We discuss the relevance of our results to some outstanding puzzles in planet formation theory: the size of the observed solar system, and the rapid type I migration of Earth mass bodies. ",
        "introduction": "\\label{sec:intro} The planetesimal hypothesis states that terrestrial planets and the icy cores of gas giants formed in a disk by the accretion of smaller solid bodies, called planetesimals. Planetesimals are defined as primitive solids of kilometer size or larger.  Such a scale is meaningful because it delimits where pairwise gravity dominates gas drag, leading to the runaway or oligarchic growth of protoplanets  \\citep{lis93}. Because planetesimal formation must occur, at least for the giant planets, while gas is still present in the solar nebula, the coupling of gas and solids may be crucial to understanding the process.  The prevailing view for planetesimal formation posits agglomerative growth from sub-micron-sized ``interstellar'' grains to km-sized bodies (cf. the review of Lissauer 1993).  However, building planetesimals through pure solid-state sticking forces has many problems.  First, in the inner solar system, interior to the ``snow line,'' the only solids available for the formation of the terrestrial planets were rocks.  Everyday experience tells us that dry silicate particulates of millimeter and larger size, such as sand, do not stick at almost any speed of attempted assemblage.  Measurements of interparticle collisions in microgravity experiments reinforce these intuitive impressions; indeed, $\\mu$m sized dust particles disrupt larger aggregates upon collision at relative velocities $>$ m/s \\citep{bw00}.  For collisions between fluffy mm-sized agglomerations of $\\mu$m sized grains, disruption also occured at speeds $>$ m/s. However at lower speeds, down to 15 cm/s, only restitution (bouncing) was observed, without any sticking \\citep{bm93}.  Relative velocities considerably higher than 1 m/s arise from chaotic motions in a turbulent disk or from differential orbital drift in a laminar disk \\citep{wc93}.  For mm-sized crystalline silicates, such as chondrules (see below), significant shattering occurs only at collision speeds $>$ several km/s \\citep{jth96}.  Thus, compact pieces of rock that make up the bulk of the material of chondritic meteorites, whose parent bodies are primitive asteroids (rocky planetesimals), are unlikely either to agglomerate or to fragment at the collisional velocities common in the nebular disk when there is still appreciable gas present to exert appreciable drag on the solids (the well-coupled limit, see Appendix).  Second, at temperatures significantly below their melting points, even ices in the outer solar system may not be much more sticky. Experiments by \\citet{sup97} found that water frost has spring-like properties and can only induce sticking for collision speeds $< 0.5$ cm/s.  Furthermore, if ice balls could grow by continued agglomeration until disruptive tidal forces became stronger than the cohesive strength of crystalline ice, we should expect many icy particulates in Saturn's rings to acquire sizes of order $\\sim 10-100$ km.  In fact, except for the occasional embedded moonlet (whose origin may lie in the fragmentation of yet larger bodies rather than from the assemblage of ordinary ring particles), the particulates in Saturn's rings have a maximum size $\\sim 5$ m \\citep{zmt85}, intriguingly close to the value implied if collective self-gravity were the only available force to assemble the largest bodies \\citep{shu84}.  Third, if we examine the most primitive meteorites, the ordinary and carbonaceous chondrites, we find no evidence for a continuous range of particulates spanning a range of sizes from less than a micron to, say, a meter or more (say, comparable to the size of the entire meteorite).  Instead, we find the largest mass fraction of such objects to be composed of chondrules: quasi-spherical, once molten, inclusions of millimeter and smaller sizes that give evidence of once having been intensely and briefly heated (perhaps multiple times). This fact suggests that sticky agglomerative events, such as those experienced by compound chondrules \\citep{rub00}, required special transient heating events to overcome the obstacle that liquids are not stable thermodynamic phases for any temperature at the low pressures prevalent in protoplanetary disks.  In any case, from the meteoritic record, the early solar system failed to generate by primitive processes any compact particulates in excess of a few centimeters in size (the largest refractory inclusions).  (See \\citet{sea01} for a promising, although unconventional, mechanism for producing the refractory inclusions and chondrules in chondritic meteorites.) Electrostatic attraction could have played a role in building looser aggregates if the individual particulates acquired significant levels of charge \\citep{mc01}.  This effect has been seen in zero-gravity experiments with particle densities well above the threshold for gravitational instability.  In order to be a relevant growth mechanism, it must be shown that tribocharging, the balance between collisional charging and ion/electron discharging, yields electrostatic attraction at much lower particle densities.  Fourth, even if a mechanism could be found to grow chondrule-sized particulates to meter-sized bodies, one would have to worry about the rapid inward orbital drift associated with gas drag that would carry such bodies from 1 AU into the protosun on a time scale of only $\\sim 10^2$ yr \\citep{wei77}.  In contrast, mm- and km-sized bodies have gas-drag drift times in excess of $10^5$ yr. Only by the direct assemblage of chondrules and related objects into planetesimals, avoiding intermediate steps, can one prevent a rapid loss of solid material from the solar nebula by gas drag.  Such a direct-assemblage mechanism exists in the gravitational instability (GI) proposal put forth by \\citet{gw73}.  A similar theory was advanced independently by \\citet{saf69}.  In the Goldreich-Ward theory, particulate settling yields a subdisk of solids that is thin and non-dispersive enough to make overdense regions undergo runaway local contraction.  This occurs when Toomre's criterion for axisymmetric GI in a rotating disk (the nonaxisymmetric case is similar) is satisfied: \\begin{equation} Q_{\\rm p} \\equiv {\\Omega c_{\\rm p} \\over \\pi G \\Sigma_{\\rm p}} < 1, \\label{eq:Q} \\end{equation} where $\\Omega$ is the angular Keplerian rotation rate, and $c$ and $\\Sigma$ are the velocity dispersion and surface density.  Throughout we use ``p'' and ``g'' subscripts to refer, respectively, to the particle and gas components of the disk.  The Goldreich-Ward instability should not be confused with the mechanism of \\citet{bo00}, who considers the formation of coreless gas giant planets from GI of gas disks.  Toomre's criterion (\\ref{eq:Q}) is equivalent (within factors of order unity) to the ``Roche'' limit which has been derived specifically for the case of stratified fluids\\citep{gl65,sek83}.  Sekiya finds GI to occur when the particulate plus gas space-density, $\\rho = \\rho_{\\rm g} + \\rho_{\\rm p}$, at a distance $r$ from a star of mass $M_\\ast$ exceeds a certain critical value in the midplane: \\begin{equation} \\rho > \\rho_{\\rm R} \\equiv 0.62 M_\\star/r^3. \\label{Roche} \\end{equation} At a distance $r=1$ AU from the Sun, $\\rho_{\\rm R} = 4\\times 10^{-7}$ g/cm$^3$, which implies a critical space density of rock that is roughly seven orders of magnitude less than the internal density of compact rock.  Thus, in the stages preceding actual planetesimal formation, we may treat the collection of solids as an additional ideal ``gas'' co-mixed with the real gas of the system.  Operating at a radius of $r=1$ AU, the self-gravitating disturbance with the most unstable wavelength creates $\\sim 5$ km planetesimals in $\\sim 10^3$ yr \\citep{ys02}.  The process occurs on a time scale longer than orbital periods ($\\sim$ 1 yr in the zone for terrestrial planet formation) because of the need to damp spin-up and random velocities during the contraction to planetesimal densities.  But the important point remains, that by leapfrogging intermediate size regimes, GI avoids the rapid inspiral of meter-sized bodies.  Unfortunately, a powerful argument has been developed against the GI scenario, which has led largely to its abandonment by modern workers in the field \\citep{wei95}.  A review of the difficulty is necessary before before we can justify a renewed attack on the basic idea.  Even in an otherwise quiescent disk, midplane turbulence may develop to stir the particulate layer too vigorously to allow sufficient solid settling to the midplane.  Without such settling the criterion (\\ref{Roche}) cannot be satisfied.  The problem lies in the vertical shear possessed by disks with a highly stratified vertical distribution of solid to gas.  The particulate-dominated sub-disk, which possesses near-Keplerian rotation, revolves somewhat faster than the surrounding gas disk, which has non-vanishing support against the inward pull of the Sun from gas pressure in addition to centrifugal effects.  The magnitude of the resulting velocity differential, $\\Delta v_\\phi = \\eta v_{\\rm K} = \\eta r \\Omega$, is proportional to $\\eta$, which roughly equals the ratio of thermal to kinetic energy of the gas: \\begin{equation} \\label{eta}\\eta \\equiv -{(\\partial P / \\partial r)\\over (2 \\rho_{\\rm g} r \\Omega^2)} \\sim (c_{\\rm g}/v_{\\rm K})^2, \\end{equation} where $P$ is the gas pressure and $c_{\\rm g}$ is the isothermal sound speed.  In the popular model of the minimum solar nebula (hereafter MSN, see \\S\\,\\ref{sec:prop}), $\\eta \\simeq 2 \\times 10^{-3} (r/\\mathrm{AU})^{1/2}$, and $\\Delta v_\\phi \\simeq 50 \\,\\mathrm{m/s}$ at $r=1$ AU  (Hayashi 1981).   Turbulent eddies with a characteristic velocity equal to the available velocity differential, $\\sim\\Delta v_\\phi$, would then prevent GI, since the Toomre $Q$ criterion requires that the particle random velocity be much smaller: $c_{\\rm p} < 7 \\cm/{\\mathrm s} \\ll \\Delta v_\\phi$, for instability (Weidenschilling 1995).  These general arguments are supported by numerical simulations \\citep{cdc93} which calculate the steady state properties of two-phase (gas and particulate) turbulence in the midplane of a MSN disk.  Particulates with internal densities of normal rock and with sizes from $10$ to $60 \\cm$ (which are assumed to have grown by other mechanisms and are moderately coupled to gas motions) acquire space densities too low for GI by an order of magnitude or more.  Their computational methods use a mixing length prescription to relate diffusivity and velocity shear through an extrapolation from laboratory studies of boundary layers.  While quite sophisticated, this approach does not directly address the mechanism by which the existence of the vertical shear generates turbulence.  Using a linear stability analysis, \\citet{sek98} confirmed that GI in a turbulent dust layer is impossible, unless the ratio of dust to gas {\\it surface} densities is significantly enhanced over normal cosmic values.  Part of the purpose of the current paper is to understand the physical basis of Sekiya's conclusion.  But for the present, let us merely note that there exist several possibilities for enhancing the solid/gas ratio in protoplanetary disks, either globally or by local concentration.  First, in the X-wind model (Shu, Shang, \\& Lee 1996), chondrules are created from material at the magnetically truncated inner edge of the disk, called the X-point.  Solids and gas are launched from the X-point in a bipolar outflow.  While the gas escapes in a collimated jet, solids of roughly millimeter size fall back to the disk at planetary distances, thus increasing the disk's solid/gas ratio.  Since $1/3$ of all material which passes through the X-point is launched in an outflow, while the remaining $2/3$ accretes onto the protostar, rocky material in the disk could be augmented by a total amount as much as $4 \\times 10^{30}~{\\rm g}$, or 30 times the amount of rock in a standard MSN model.  Such an enhancement factor is more than sufficient, as we shall see, to promote GI in the sub-disk of solids.   In point of fact, because of efficiency considerations in the manufacture of chondrules and refractory inclusions and their irradiation to produce short-lived radionuclides (see, e.g., Gounelle et al. 2001), there are reasons to believe that an amount of rock not much larger (but perhaps a few times larger) than that contained in a MSN was recycled by the X-wind to the disk.  We shall find that what is important is the ratio of solid to gas {\\it surface densities} in the disk.  Once we accept that this ratio need not be cosmic (e.g., rock to gas = $4\\times 10^{-3}$), then there exists no a priori theoretical objection to a revival of the Goldreich-Ward mechanism for forming planetesimals.  Second, since solids tend to settle towards the midplane, the surface layers of a disk should become relatively gas-rich.  Thus, any mechanism that removes material from the surface of a stratified disk would increase the solid to gas ratio computed in terms of vertically projected column densities.  Possibilities for such surface removal include (1) photoevaporation, which dominates in the loosely bound outer disk \\citep{sjh93}; (2) layered accretion, which occurs if only the surface layers of a disk are sufficiently ionized to support magneto-rotational turbulence \\citep{gam96}; and (3) stripping by stellar winds, which is probably more effective near than far from the star \\citep{hyj00}.  The amounts of solid to gas enhancements achievable by these processes are difficult to predict, but the timescales for the dominant processes are typically $< {\\rm few} \\times 10^6$ years, and thus likely to be relevant to the evolution of protoplanetary disks.  Third, gas drag can also lead to local enhancement of particulate concentrations by a variety of mechanisms.  (1) Isotropic turbulence with a Kolmogorov spectrum concentrates particles in numerical \\citep{se91} and laboratory \\citep{fke94} experiments.  Extrapolation to the high Reynolds numbers of protoplanetary disks implies concentration of chondrules by factors of up to $10^5$ \\citep{cuz01}. However, these estimates do not take into account the redispersal of the concentrated pockets of solids if the turbulent eddies are intermittent and do not maintain fixed centers.  (2) Similarly, disk vortices could concentrate chondrules as well, but they are more effective for meter sized bodies \\citep{fmb01}.  There also remains the issue whether vortices will rise spontaneously in protoplanetary disks if there are no natural stirring mechanisms.  (3) Secular instabilities associated with gas drag might concentrate particulates, even without self-gravity \\citep{gp00}.  Goodman and Pindor assumed that gas drag acts collectively on a particulate ``sheet'', a valid approximation when turbulent wakes overlap.  Unfortunately, wake overlap seems unlikely unless the particles are fairly closely packed, in which case GI would already be effective.  In a related context, \\citet{war76} has shown in an underappreciated study that viscous drag modifies the standard GI criterion through the introduction of an additional instability that can occur at values of $Q \\gg 1$. However, this additional instability, being secular in nature, has a much smaller growth rate than the usual Goldreich-Ward mechanism.  (4) In \\S\\,\\ref{sec:drift}, we develop the simplest concentration mechanism for particulates: radial migration due to gas drag.  The fundamental assumptions required for the results of this paper are the existence, at some epoch of the disk's evolution, of (1) relatively quiescence in the midplane regions, even though the surface layers may be undergoing active accretion \\citep{gam96}, and (2) compact solids with chondrule-like properties that are well, but not perfectly, coupled to the gas through mutual drag.  Under these conditions, we argue (1) that vertical shear can only induce a level of midplane turbulence that has limited ability to stir solids, with the critical value of the surface density of solids being given roughly by $\\Sigma_{\\rm p,c} \\sim \\eta r \\rho_{\\rm g}$; and (2) that gas drag alone can lead to a global {\\it radial} redistribution of solids so that the local surface density of solids, $\\Sigma_{\\rm p}$, can exceed the critical value, $\\Sigma_{\\rm p,c}$.  Therefore, whether a recycling of solids occurs by the X-wind mechanism or a radial redistribution of solids occurs by simple gas drag, we conclude that the planet-forming zones of the primitive solar system can achieve the requisite conditions for the formation of planetesimals on a time scale comparable to the typical lifetime, $\\sim 3\\times 10^6$ yr, that has been inferred for the disks of T Tauri stars \\citep{hll01}.  However, the margin for success is not large, and it could be that planet formation is a less universal phenomenon than it has been widely touted to be, and that it has a much greater diversity of outcomes (including a complete failure to form any planets) than suspected prior to the discovery of extrasolar planets (Mayor and Queloz 1996, Marcy and Butler 1998).  Indeed, the very fact that planetesimal formation may involve a threshold phenomenon, namely the existence of a nontrivial critical surface density, $\\Sigma_{\\rm p,c} \\sim \\eta r \\rho_{\\rm g}$, implies that low-metallicity systems should be much less likely to form planets than high-metallicity systems, a correlation which seems already to be present in the empirical literature (Gilliland et al. 2000, Laughlin 2000).  This paper is organized as follows.  In \\S\\,\\ref{sec:prop} we present the basic properties of our disk models.  After reviewing the techniques of \\citet{sek98} for deriving density distributions of well-coupled particles in \\S\\,\\ref{sec:tech}, we physically interpret, in \\S\\,\\ref{sec:sat}, the midplane density singularities which occur in these profiles as evidence that midplane turbulence can only stir a finite amount of material.  This allows us, in \\S\\,\\ref{sec:enh} to calculate the solid/gas enhancements required for GI in various model disks.  We show that aerodynamic drift can concentrate particles of a given size radially in the disk on cosmogonically interesting time scales in \\S\\,\\ref{sec:single}, and we generalize to distributions of particle sizes in \\S\\,\\ref{sec:dist}.  In \\S\\,\\ref{sec:gotsat} we evaluate whether this concentration mechanism provides enough enhancement to induce GI.  Closing remarks are made in \\S\\,\\ref{sec:disc}. ",
        "conclusions": "}  In this paper we have shown that planetesimals can form by midplane gravitational instabilities despite Kelvin-Helmholtz stirring if the ratio of particle to gas surface densities is increased above cosmic abundances.  We have also presented a simple drift mechanism which can provide most or all of the particle enhancement required for this to occur.  This mechanism has the advantage that it must operate in passive protoplanetary disks, and does not depend on assumptions about accretion physics.  There are some attractive features of this scenario not yet discussed.  Curiously, the mass in our planetary system appears to be significantly truncated outside of 40 AU, \\ie in the Kuiper belt region and beyond \\citep{tb01}, whereas T-Tauri disks typically extend to several hundred AU.  The size disparity is even greater if Uranus and Neptune migrated from a location interior to Saturn's orbit \\citep{tdl99}.  The process of drift induced enhancement offers an explanation.  The outer disk is drained relatively of its solid resources by inward particulate drift, and thereby becomes or remains inhospitable to the formation of planetesimals.  Another troubling aspect of planet formation theories is the ``Type-I'' migration of Earth-sized and larger (but not large enough to open a gap) bodies due to density waves torques exerted on the gas disk \\citep{war97}.  When the resonant torques are assumed to damp locally in the disk, Earth-mass bodies (at 5 AU) migrate inwards in $10^5$ years in a MSN disk.  This timescale is inversely proportional to the mass of the body, but increases sharply once the gap-opening mass, typically $10-100 M_\\oplus$ is reached.  However the drift speed is proportional to $\\Sigma_{\\rm g}$.  Thus if planetesimals form due to depletion of gas below MSN values (or equivalently the enhancement of solids in a very low mass disk), then subsequent earth-mass cores suffer less migration.  Gas depletion by one order of magnitude would significantly increase the survival odds of a nascent planetary system, and still leave enough gas for the formation of giant planet atmospheres.  To summarize, in conventional cosmogonies with unit sticking probabilities and cosmic abundances of solids and gas in a MSN \\citep{lis93}, planetesimal formation is easy, occurring on a time scale of $\\sim 10^4$ yr or less, while giant planet formation is hard, requiring time scales in excess of the typical lifetimes of T Tauri disks, $\\sim 3 \\times 10^6$ yr.  Overall surface density enhancements (gas and dust) above MSN values can speed up post-planetesimal growth \\citep{tdl02}, but this solution would exacerbate the problem of the Type-I migration of planetary embryos (Ward 1997).  It would also require a finely tuned mechanism to remove the considerably greater amount of extra gas and solids from the solar system.  In an unconventional cosmogony, where sticking probabilities are zero (except for the special mechanisms that produce chondrules and refractory inclusions near the protosun), where gas is depleted, and solids are enhanced relative to standard MSN values, the conditions for the formation of planetesimals and giant planets might be intimately tied to the gas-dust evolution of the nebular disks of T Tauri stars.  The two processes would then naturally acquire similar time scales, related by a single continuous process of gravitational growth in a gas-dust disk.  Attractive byproducts of this unconventional approach would be a corresponding alleviation of the problems of Type I migration and gas-disk dispersal, as well as a possible understanding of why the Kuiper belt marks a sudden apparent truncation of the primitive solar system."
    },
    "0207/astro-ph0207646_arXiv.txt": {
        "abstract": "{\\small We have studied the spectral and timing behaviour of the atoll source 4U 1608--52. We find that the timing behaviour of 4U 1608--52 is almost identical to that of the atoll sources 4U 0614+09 and 4U 1728--34. Recently Muno, Remillard \\& Chakrabarty (2002) and Gierlinski \\& Done (2002) suggested that the atoll sources trace out similar three--branch patterns as the Z sources. The timing behaviour is not consistent with the idea that 4U 1608--52 traces out a three--branched Z shape in the color--color diagram along which the timing properties vary gradually.} ",
        "introduction": "In this work we use all available data from RXTE's PCA to simultaneously study the spectral and timing properties in the transient low mass X-ray binary 4U 1608--52. We calculate a color--color diagram. As the energy spectrum of a source changes, it moves through the diagram. To study the timing we calculate Fourier power density spectra and fit them with the multi--Lorentzian fit function; a sum of Lorentzian components plus an occasional power law to fit the very low frequency noise \\cite{bpk02,vstr02}. It has been recently proposed \\cite{muno02,gier02} that the atoll sources trace out similar three--branch patterns as the Z sources; one of our goals in this work is to test this hypothesis.\\\\ ",
        "conclusions": ""
    },
    "0207/astro-ph0207470_arXiv.txt": {
        "abstract": "The presence of a small amount of hydrogen is expected in most single degenerate scenarios for producing a Type Ia supernova (SN~Ia). While hydrogen may be detected in very early high resolution optical spectra, in early radio spectra, and in X-ray spectra, here we examine the possibility of detecting hydrogen in early low resolution spectra such as those that will be obtained by proposed large scale searches for nearby SNe~Ia. We find that definitive detections will require both very early spectra (less than 5 days after explosion) and perhaps slightly higher amounts of hydrogen than are currently predicted to be mixed into the outer layers of SNe~Ia. Thus, the non-detection of hydrogen so far does not in and of itself rule out any current progenitor models. Nevertheless, very early spectra of SNe~Ia will provide significant clues to the amount of hydrogen present and hence to the nature of the SN~Ia progenitor system. Spectral coverage in both the optical and IR will be required to definitively identify hydrogen in low resolution spectra. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are now considered to be among the best ``standardizable candles'' available \\citep{philetal99}, and they are the tool of choice for observational cosmology \\citep{perletal99,riess_scoop98,riess00,goldhetal01}. However, the progenitor system remains in doubt \\citep[for a review see][]{prog95}. The most widely accepted current view is that a C+O white dwarf accretes hydrogen (or perhaps helium) from a companion star until it reaches the Chandrasekhar mass and explodes due to the thermonuclear burning of C+O. While the merging of two white dwarfs (the double degenerate scenario) remains a viable progenitor for SNe~Ia, recent theoretical work has focused on the accretion of material via Roche lobe overflow from a MS or subgiant companion (hydrogen cataclysmic variables) or via wind accretion from a red-giant companion (symbiotic stars). A recent search for radio emission appears to rule out the peculiar SN~Ia 1986G as having a symbiotic star progenitor \\citep{eck95}. \\citet{maxted00} have found a candidate double degenerate system containing a 0.5~\\msol\\ white dwarf and the companion is more massive that 0.97~\\msol, so that the total mass of the system exceeds the Chandrasekhar mass.  While it is clear that double degenerate progenitors do exist, it is still not clear whether the merger of a double degenerate system leads to collapse \\citep[due to electron capture,][]{sainom98,mochkoasi97,TY96} or explosion. It is also unclear whether sufficiently massive double degenerate systems exist in sufficient numbers to produce the observed sample of SNe~Ia \\citep{prog95}. Therefore the detection of hydrogen in a SN~Ia would greatly expand our knowledge of the progenitor system. The question of the identity of the progenitor system remains a significant hurdle for the general acceptance of the observational cosmology results; is a gap in our knowledge of the binary stellar evolution; and impedes our understanding of the chemical evolution of galaxies. There have been attempts to detect hydrogen via looking for narrow H$\\alpha$ lines \\citep{cum94d96}, in the radio \\citep{eck95}, and in X-rays \\citep{SP92A93,schlegelrpp95}.  While all of these methods should be pursued, it seems likely that we will obtain large numbers of very early low resolution optical spectra of nearby SNe~Ia from dedicated searches that will begin taking data soon. Thus, we examine the possiblity of detecting hydrogen in early SN~Ia spectra. ",
        "conclusions": "The mixing of hydrogen and solar composition material into the outer layers of a Type~Ia supernova model has little effect on the spectra, unless the solar material replaces a substantial fraction of the \\snia\\ C+O ejecta. In fact the non-detection of hydrogen in SNe~Ia is expected in most of the present progenitor scenarios.  The effects of mixing in the outer layers diminish as the photosphere recedes into deeper layers of the ejecta. In our tests the most detectable results were obtained by replacing half of the mass of the unburned C+O in the outer layers with solar abundance material (for a total of $2.7 \\times 10^{-2}$~\\msol\\ of hydrogen). These models produced reasonably strong signals 5 and 10 days after the explosion, but are less prominent at day 15, still 5 days before maximum light. One fifth of that mass of mixed hydrogen is able to produce identifiable signals at 5 days after explosion  and approximately half that mass is required 10 days after the explosion.  The effects of additional electrons from the hydrogen atoms may also be a useful diagnostic when coupled with detailed analysis of high quality data. Replacing half of the C+O with solar material, but to different depths, we found that only at 5 days after explosion did any model other than mixing to the bottom of the C+O layer provide any chance of detection. This is consistent with a picture based on the mass of the mixed hydrogen.  Approximately 0.02~\\msol\\ of solar composition material swept up from the wind of or ablated off of the companion is needed for a detection at 10 days after explosion (10 days before maximum light). The deeper it is mixed the later it will be detectable.  For quantities of mixed hydrogen near the detection threshold, the \\ion{C}{2} feature in the red emission peak of \\ion{Si}{2} complicates the identification of \\halpha\\ absorption. Confirmation requires several epochs to trace the development of the spectral features, and ideally another signal, such as narrow circumstellar hydrogen lines. Even earlier spectra would be advantageous.  Additional spectral signatures that would prevent confusion of high-speed carbon ejecta that could mimic the \\halpha\\ from mixed hydrogen are needed. Figure~\\ref{fig:paschena} shows the region around the P$\\alpha$ line.  While the P$\\alpha$ line shows a pretty strong signature at the early epochs the P$\\beta$ line is not easily discerned, thus, combined IR and optical data would be likely be required for a definitive detection. \\citet{bowersetal97} and \\citet{hern98bu00} compared the IR data for SNe~Ia and most of the data is obtained at late times (SN 1998bu has data at about 13 days after maximum); the data are obtained much later than we desire and without detailed spectral analysis it would be difficult to draw any conclusions. SN~1999ee \\citep{hamuy99ee99ex02} does have very early IR spectra (approximately 11 days before maximum). It shows no sign of any feature near P$\\beta$, and the region around P$\\alpha$ is not covered spectroscopically. Nevertheless, early IR spectra are obtainable and should be vigorously persued.  The results of \\citet{marietta00} suggest that only about $10^{-4}$~\\msol\\ of hydrogen would be stripped from the companion at $v > 15,000$~\\kmps, however, they find that $10^{-3}$~\\msol\\ of hydrogen would be stripped from the companion at $v > 10,000$~\\kmps. Thus, our results indicate that very early photospheric spectra of SNe~Ia might be able to detect the presence of hydrogen, and certainly could be used as a signal to activate high resolution, radio, and X-ray programs.  Continued work in the radio, X-ray, and high resolution optical spectra can verify and reinforce any putative hydrogen detections."
    },
    "0207/astro-ph0207316_arXiv.txt": {
        "abstract": " ",
        "introduction": "}} The frequency of both lunar and Martian meteorites on the Earth indicates that the transfer of planetary material is common in the solar system. Vigorous hydrologic or tectonic cycles, past or present, prevent most nearby planetary bodies from serving as long-term repositories of this material. The Moon is an important exception, however. Strategically located within the inner solar system, the Moon has theoretically collected material from all of the terrestrial planets since its formation. Lacking an atmosphere and widespread, long-lasting volcanism, the Moon has potentially preserved meteorites from Mercury through the asteroid belt. While the lack of an atmosphere prevents a soft landing on the lunar surface, its low gravity means particles with small velocities with respect to the Moon will experience relatively low impact velocities. Moreover, unlike on other terrestrial planets, Martian, Venusian and Terran meteorites, blasted off their respective planets 3.9 Ga during the Late Heavy Bombardment, should still exist on the surface of the Moon. Such meteorites are likely to contain uniquely preserved remains of these planets that are not available elsewhere in the Solar System. In particular, Terran meteorites on the Moon may provide a substantive geological record for ancient Earth, corresponding to or predating the period for which the earliest evidence for life exists.  In considering the argument in favor of searching the Moon for Terran meteorites, the SNC meteorites represent obvious analogs. A considerable amount of Martian geoscience is built upon little more than a dozen samples. From them have been inferred important constituents of the atmosphere, mantle and core, the extent of interaction between the Martian hydrosphere and lithosphere, constraints on Martian water abundance, and the nature of a Martian carbon sink, while the lightly shocked condition of the meteorites has stimulated developments in the understanding of impact physics (McSween 1994). Terran meteorites have the potential to provide similar information, extending and broadening Earth's geologic record for a time period that has otherwise left little or no physical evidence. The rocks' elemental composition and mineralogy (in particular, hydration) could be used to constrain characteristics of the early crust and mantle, the global oxidation state, the extent of planetary differentiation, and the availability of water. Volatile inclusions sampling noble gases, carbon dioxide and molecular nitrogen could clarify atmospheric origin and evolution and, along with the meteorite mineralogy, could provide substantive constraints on early atmospheric concentrations (Bogard and Johnson 1983). Direct measurement of the timing, extent and planetary effects of the heavy bombardment by careful dating of Terran meteorites is also possible and would perhaps be the most robust and significant scientific reward of this project.  In addition to the scientific benefits listed above, which alone justify searching the Moon for Terran meteorites, a fraction of Earth-derived material on the Moon could contain geochemical and biological information, in the form of isotopic signatures, organic carbon, molecular fossils, biominerals or even, theoretically, microbial fossils. Again, analogy to SNC meteorites is instructive. Questionable interpretations of structures within the Martian meteorite ALH84001 as microbial fossils (McKay {\\em et al.} 1996) and, more recently, evidence supporting deposition there of magnetite by biomineralization (Thomas-Keprta {\\em et al.} 2001; Friedmann {\\em et al.} 2001) have been used to argue that this meteorite contains vestiges of ancient Martian life. While such interpretations remain highly controversial, they support the general principle that Terran meteorites should be examined for potentially novel evidence concerning early Earth life. Such evidence could substantiate or extend a contested fossil record that begins 3.5 Ga (Awramik 1982; Schopf and Packer 1987; Buick 1990; Brasier {\\em et al.} 2002) and geochemical evidence from even earlier periods, between 3.7 and 3.85 Ga (Schidlowski 1988; Mojzsis {\\em et al.} 1996; Rosing 1999; but see Moorbath 2001 and Fedo and Whitehouse 2002).  In addition, the Moon may preserve material not only from Earth, but also from Venus and the asteroid belt. The only attainable record of Venus' early surface geology, otherwise catastrophically erased 700 million years ago (McKinnon {\\em et al.} 1997), is probably on the Moon. Similarly, a record of the type, characteristics and origins of the heavy bombardment impactors themselves may be available on the Moon. Such a record would clarify not only the geological history of Earth, but also its chemical and biological history -- especially since these impactors were potentially major sources of biotic precursors on early Earth (Pierazzo and Chyba 1999).  Mars is presently the focus of attention with regard to the search for early signs of life outside Earth. Ironically, the Moon may be the better place to search for the remains of both early Martian and early Terran life. Most significantly, the Moon lacks the water capable of carrying contaminants into the interior of rocks through cracks. While gardening from micrometeorites is less severe on the surface of Mars, the deeply buried regolith on the Moon provides some protection.   Finally, the Moon is also a perfect testbed for targeted sample return.  For these reasons, we determined the likelihood that early remains of Earth, Mars, and Venus have been preserved on the Moon in high enough concentrations to motivate a search mission. While others (Gladman 1997; Halliday {\\em et al.} 1989) provide estimates for the transfer efficiency and total number of Martian meteorites impacting the Earth, there are no estimates for the abundance of Terran, Martian, and Venusian meteorites on the Moon. In Section 2, we consider the transfer of impact ejecta from Earth to the Moon's surface immediately following an impact event.  By considering both the slow moving, Earth-bound ejecta and the high velocity ejecta that achieves orbit around the Sun, we estimate the total transfer efficiency of the material.  Through a separate numerical simulation, we estimate a rough transfer efficiency for Venusian material.  Additionally, we compute the range of expected impact velocities on the surface of the Moon. In Section 3 the results of Section 2 will be combined with mass flux estimates to calculate concentrations of Terran and Venusian meteorites on the Moon.  In addition, the contribution from Martian meteorites is estimated from the literature. Section 4 includes a discussion of the survivability of the material delivered to the Moon via impacts over a range of velocities. ",
        "conclusions": "}}  In this paper, we have explored the Moon as an ideal location to search for remnants of the early solar system, particularly samples of Earth not currently available to researchers.  The Moon's proximity and relatively unaltered surface makes research missions viable, and the Terran samples on the surface are unavailable anywhere else in the solar system.  We have argued that Terran materials are abundant and near the surface, with a significant fraction retaining their geochemical and biological signatures for detailed analysis.  In addition, since the majority of Terran samples date from the end of the Late Heavy Bombardment, the samples in the lunar ``attic'' are a unique probe of the early conditions on Earth, and potentially contain clues to the earliest forms of life.  The Terran material is delivered in one of two ways, through direct transfer or orbital transfer.  We have found that orbital transfer of material is more efficient for the time since the end of the Late Heavy Bombardment, with direct transfer of material only comparable in early epochs. The amount of Terran material, 7 ppm, is sufficiently large to consider a search mission.  Before any such mission is attempted, the current stock of lunar material (approximately 400 kg worth) should be searched for Terran material.  Given a concentration of 7 ppm, there should be roughly 3 grams of Earth material in the current lunar samples.  Since lunar fines contain more than 10 million particles per gram, a technique of infrared spectroscopy coupled with microscopic imagery could distinguish hydrated silicates (common on Earth) from the dry, unhydrated lunar fines.  Such material could then be isotopically analyzed to confirm its terrestrial origin.  While this is not likely to yield much in the way of information about the early Earth, it would act as a proof of concept and a baseline for future missions.  Recovery and identification of Terran samples from the Moon represents a significant challenge.  Still, returning a sample from the Moon bearing the remains of Earth life may be orders of magnitude easier than returning, say, Martian samples from Mars bearing the remains of Martian life.  In this sense, lunar missions represent an interesting proving ground for these types of endeavors.  The risk of contamination and relative scarcity of Terran material makes sample return missions difficult.  Therefore, robotic missions need to be developed capable of finding Terran material using advanced in situ measuring devices to help identify samples and largely driving the analysis on the Moon.  The existence of a facility on the Moon to recover and analyze samples would guard against any contamination from Earth.  However, for a complete analysis of the material, we suggest the best way to conduct these studies is on site measurements by human observers - in essence, a return to the Moon.  \\newpage  \\begin{center} {\\bf ACKNOWLEDGEMENTS} \\end{center} We thank Tom Quinn, Joachim Stadel, and Derek Richardson for providing the dynamics simulation software, PKDGRAV, and other analysis and visualization tools used in the numerical simulations.  We also thank Don Brownlee, Toby Smith, and Katherine Shaw for discussions concerning the fate of Terran ejecta on the surface of the Moon. Finally, we appreciate Chris Stawarz and Zoe Leindhardt for their assistance in the numerical modeling.  The paper benefited greatly from a thorough review by Brett Gladman, and additional reviews by Norm Sleep and Jay Melosh.  Also, thanks to Jay Melosh for helping confirm our new derivation of Eq. 38.  This research was supported by the NSF-IGERT trainingship in Astrobiology, an NDSEG fellowship, and the NASA Astrobiology Institute.  \\appendix \\renewcommand{\\thesection}{{\\small {\\bf Appendix \\Alph{section}.}}}"
    },
    "0207/astro-ph0207120_arXiv.txt": {
        "abstract": "Images obtained in $g', r'$, and $i'$ with the Gemini Multi-Object Spectrograph (GMOS) on Gemini North are used to investigate the metallicity and stellar content of the M31 dwarf spheroidal companion galaxy And V\\@. Red giant branch (RGB) stars are traced out to radii in excess of $126\\arcsec$ from the galaxy center, indicating that And V extends over a diameter approaching 1 kpc. The mean $(g'-i')$ color of the RGB does not change with radius. Based on the slope of the RGB in the $(i', g'-i')$ color-magnitude diagram (CMD), we conclude that the metallicity of And V is [Fe/H] = $-2.2 \\pm 0.1$. This is lower than earlier estimates, and places And V squarely on the relation between metallicity and integrated brightness defined by other dwarf spheroidal and dwarf elliptical galaxies. In contrast to many of the Galaxy's dwarf spheroidal companions, there is no evidence for a statistically significant population of luminous asymptotic giant branch stars near the center of the galaxy. ",
        "introduction": "There is a growing body of evidence that galaxies form in a hierarchal manner, with low mass systems serving as the basic building blocks of galaxy formation. Not only are hierarchal galaxy formation models able to reproduce a wide range of the observational properties of nearby galaxies (e.g. Somerville \\& Primack 1999; Cole et al.\\ 2000), but they also provide a natural explanation for extended stellar streams (e.g. Ibata et al. 2001; Martinez-Delgado et al. 2001) and the overall structural properties of the Galactic halo (Bullock, Kravtsov, \\& Weinberg 2001). The merging process continues to the present day, and it has been suggested that the ultimate fate of the Local Group may be a merger that produces an elliptical galaxy (Forbes et al.\\ 2000).  Studies of present-day dwarf galaxies can provide insights into the systems from which larger galaxies formed. For example, there are indications that the Galactic halo could not have formed from the accretion of systems with chemical enrichment histories similar to those of the low mass dwarf spheroidal companions of the Milky-Way (Shetrone, C\\^{o}t\\'{e}, \\& Sargent 2001). One possible explanation is that the chemical evolution of the present-day dwarf spheroidal companions of the Milky-Way has been affected by tidal interactions with the Galaxy (e.g. Mayer et al. 2001), with the result that these systems evolved differently from those that merged early-on to form the Galactic halo.  CDM-based simulations predict many more low-mass systems in the Local Group than have been detected (e.g. Kauffmann, White, \\& Guiderdoni 1993; Klypin et al.\\ 1999), and this provides a strong motivation to search for heretofore undiscovered dwarf galaxies. Armandroff et al.\\ (1998) discuss such a search for low surface brightness companions of M31, and report the detection of the dwarf galaxy And V\\@. Using a color-magnitude diagram (CMD) obtained from follow-up CCD images, Armandroff et al.\\ (1998) conclude that the metallicity and mean central surface brightness of And V are superficially unremarkable, in the sense that these quantities are similar to those of two other companions of M31: And I and And III.  The apparent similarity to And I and And III notwithstanding, with a mean abundance [Fe/H] $= -1.5$ and absolute magnitude M$_V = -9$, And V falls well off of the relation between metallicity and M$_V$ that is otherwise well-defined by observational data (e.g.\\ Caldwell 1999). The [Fe/H] -- M$_V$ relation is generally thought to be imprinted during the early stages of galaxy evolution, and is a measure of the extent of chemical enrichment before the onset of winds that purge a galaxy of its interstellar medium. This relation is of fundamental astrophysical importance; for example, it is likely the physical basis behind the color-magnitude relation of early-type galaxies (Kodama \\& Arimoto 1997), and is a basic prediction of CDM-dominated galaxy formation models (e.g.\\ Dekel \\& Silk 1986, Cole et al.\\ 2000).  Why does And V depart from the [Fe/H] -- M$_V$ relation? A reasonable starting point for answering this question is to check the existing metallicity determination. The CCD data used by Armandroff et al.\\ (1998) to construct the CMD of And V were recorded during $1\\arcsec$ seeing conditions, and are restricted to the upper $\\sim 1.5$ mag of the red giant branch (RGB). Their CMD shows a scatter of $\\pm 0.1$ mag, and includes only a fraction of the stars in the galaxy. Given the potential significance of And V, we decided to re-investigate the metallicity of this galaxy using (1) deep multi-color images recorded during sub-arcsec seeing conditions and (2) spectra of individual stars that include the near-infrared Ca II triplet, which is a metallicity indicator (e.g.\\ Armandroff \\& Da Costa 1991). In the present paper we discuss the photometric measurements of stars in And V; the spectroscopic observations will be discussed in a second paper (Da Costa et al.\\ in preparation). Based on the slope of the giant branch in the CMD, we find that [Fe/H] $= -2.2\\pm 0.1$, indicating that And V is more metal-poor than previously thought. ",
        "conclusions": "Deep $g'$, $r'$, and $i'$ images obtained with GMOS on the Gemini North telescope have been used to investigate the photometric properties of bright giants in the dwarf spheroidal galaxy And V\\@. The data include stars that are three magnitudes in $i'$ fainter than the RGB-tip, making the current paper the deepest photometric study of this galaxy to date. Based on the slope of the upper RGB on the $(i', g'-i')$ CMD, we conclude that And V has a metallicity [Fe/H] $= -2.2 \\pm 0.1$, which is significantly lower than was estimated by Armandroff et al.\\ (1998). With the previous metallicity estimate of [Fe/H] $= -1.5$, And V had the largest deviation from the trend between [Fe/H] and M$_V$ among dwarf galaxies plotted in Figure 4 of Caldwell (1999); however, our revised metallicity places And V squarely on this relation.  A number of factors, including environment, likely contribute scatter to the relation between M$_V$ and [Fe/H], and studies such as ours provide data that will ultimately allow the intrinsic dispersion about the [Fe/H] -- M$_V$ relation to be measured. Our revised metallicity for And V tightens the relation between integrated brightness and metallicity for low mass systems, and one implication of such a tight relation is that the M/L ratio of dwarf systems must not differ by large amounts at a given M$_V$. The chemical contents of dwarf elliptical galaxies in the M81 group appear to be defined by total baryonic mass (Caldwell et al.\\ 1998), rather than other properties such as the central stellar concentration, possibly suggesting that the [Fe/H] -- M$_V$ relation has its origins in global galaxy properties. A caveat is that the structural properties of low mass galaxies in hierarchal systems may change with time; for example, tidal effects may alter the structural properties of dwarf systems (e.g.\\ Cuddeford \\& Miller 1990, Mayer et al.\\ 2001).  There is an HI cloud close to And V on the sky but, with $v_{\\odot} = -176 \\pm 1$ km/sec (Blitz \\& Robishaw 2000), the cloud has a velocity that is very different from that of And V, for which Evans et al.\\ (2000) measure $v_{\\odot} = -403 \\pm 4$ km/sec, indicating that the cloud is likely a chance superposition. The absence of an HI reservoir for And V may not be surprising since, with a distance of 118 kpc from the center of M31, this galaxy falls within the 250 kpc radius where Local Group dwarf spheroidals appear to have low HI contents (Blitz \\& Robishaw 2000). In fact, we do not find a statistically signicant population of upper AGB stars near the center of And V, suggesting that the galaxy did not experience significant star formation during intermediate epochs. At present, the only dwarf spheroidal companion of M31 to show evidence for star formation during intermediate epochs is And II (Da Costa et al.\\ 2000), whereas many of the dwarf spheroidal companions of the Milky-Way contain intermediate-age populations. It thus appears that the star-forming histories of the Milky-Way and M31 dwarf spheroidal systems may have been systematically different."
    },
    "0207/astro-ph0207299_arXiv.txt": {
        "abstract": "We examine the expected signal from annihilation events in realistic cold dark matter halos. If the WIMP is a neutralino, with an annihilation cross-section predicted in minimal SUSY models for the lightest stable relic particle, the central cusps and dense substructure seen in simulated halos may produce a substantial flux of energetic gamma rays. We derive expressions for the relative flux from such events in simple halos with various  density profiles, and use these to calculate the relative flux produced within a large volume as a function of redshift. This flux peaks when the first halos collapse, but then declines as small halos merge into larger systems of lower density. Simulations show that   halos contain a substantial amount of dense substructure, left over from the incomplete disruption of smaller halos as they merge together. We calculate the contribution to the flux due to this substructure, and show that it can increase the annihilation signal substantially. Overall, the present-day flux from annihilation events may be an order of magnitude larger than predicted by previous calculations. We discuss the implications of these results for current and future gamma-ray experiments. ",
        "introduction": "Dark matter is omnipresent in the universe. Most of it is non-baryonic, and a favoured candidate is a weakly interacting massive particle (WIMP), often generically taken to be the lightest stable relic particle surviving from when the universe was supersymmetric. The freeze-out of such a neutralino of mass $m_\\chi$ occurs at $kT\\sim m_\\chi/20$, and the annihilation cross-section determines the current value of the CDM density $\\Omega_{\\rm cdm}h^2.$ For minimal SUSY, for example, one can derive the relation between annihilation cross-section and particle mass, and compute the various annihilation products, including continuum gamma rays from $\\pi^0$ decays and line gamma rays from rare quark decays (c.f.\\ Bergstr\\\"om 2000 for a recent review). While cosmological observations specify $\\Omega_{\\rm cdm}h^2\\approx 0.1,$ scanning over minimal SUSY parameter space results in an uncertainty in the gamma-ray emissivity of several orders of magnitude.   Our galactic halo is a logical place to look for evidence of annihilations. Unfortunately, for a uniform dark halo with a realistic density profile, even the most optimistic models fall short of the observed diffuse high-galactic-latitude gamma-ray flux, as measured by EGRET, by an order of magnitude or more (Ullio et al.\\ 2002). In fact, high-resolution numerical simulations show that the dark halo has considerable substructure (Klypin et al.\\ 1999; Moore et al.\\ 1999). This substructure may boost the annihilation flux substantially. If this is indeed the case, then the isotropic diffuse background flux from the many small halos that merged in the past into our halo and others may also become significant. There is considerable uncertainty among cosmological halo simulators, however, about the quantitative role of substructure in our dark halo, and of the concentration of the substructure and of the dark matter itself towards the centre of the galaxy.   No reliable estimates have been given up till now of the properties of halo substructure on very small scales, and in particular of their dependence on halo mass and their evolution with cosmological epoch. We have developed a semi-analytical model of halo formation which is capable of following the key physics of tidal disruption of substructure during merging. In principle this approach has arbitrarily high resolution. Hence we are able to provide robust calculations of the annihilation flux generated within our own halo, but also especially of the component generated during the evolution of structure at early times, and visible as an isotropic gamma-ray background at the present day.   In this paper, we calculate the flux produced by annihilations in simple CDM halos, relative to the flux produced in a uniform background. We correct this result for halo substructure, and integrate over a large volume to determine the cosmological background from WIMP annihilation as a function of redshift. The outline of this paper is as follows. In section 2, we define a dimensionless flux multiplier $f$ that accounts for the enhanced rate of two-body interactions produced by inhomogeneities in the dark matter distribution, and determine its value for simple virialised halos. In section 3 we calculate $f$ for cosmological volumes, using analytic estimates of the halo mass function and halo concentrations, and determine its redshift dependence. Finally, in section 4 we study the contribution to $f$ from substructure within virialised halos, and calculate $f$ for a set of realistic halos generated using a semi-analytic model of halo substructure. Throughout this paper we assume a Lambda-CDM (LCDM) cosmology with a cosmological constant $\\Lambda_0 = 0.7$, a matter density $\\Omega_{{\\rm m},0} = 0.3$ and a Hubble parameter $H_0 = h \\times\\ 100 {\\rm km\\,s^{-1}}$, with $h = 0.65$. ",
        "conclusions": "We have calculated the amount by which structure in dark matter on subhalo, halo and cosmological scales amplifies the expected signal from neutralino annihilation. The overall magnitude of the effect remains uncertain at several levels, some more important than others. Uncertainty in the fundamental density profile which characterises halos, or more specifically in the slope of their density profiles at small radii, can change the flux multiplier for simple halos by a factor of 20--25, assuming concentrations typical of galaxy halos (see figure \\ref{fig:1}). There is evidence in recent simulations for some excess mass in the inner regions, relative to the NFW profile (Power et al.\\ 2002), so a reasonable estimate for the flux multiplier may be 3--4 times the NFW value used for most of the calculations in this paper, although studies of the Milky Way (e.g.\\ Binney \\& Evans 2001) suggest less dark matter in its central regions than these profiles would imply. There is no direct evidence for a pure Moore profile extending to very small scales, although if this is the case the resulting flux from simple halos will be 20--25 times the NFW value used here.  The mass function of dark matter halos, averaged over large volumes, is also uncertain, particularly at the low mass end. The mass function determined from simulations is better fit by the form proposed by Sheth and Tormen (ST) than by the traditional Press and Schechter (PS) mass function. Using the ST mass function to predict the background flux reduces its amplitude by about 1.5 over PS. We assume that the scale invariance of dark matter substructure will be broken at very small masses, perhaps close to Jeans mass at recombination, for instance. At a minimum, there should be some maximum density for dark matter halos, or equivalently some maximum for $\\sigma(M)$ at early times. Varying the lower mass limits over a reasonable range ($10^4$--$10^8\\,M_{\\odot}$) changes the flux by a factor of 2 or so for simple halos, but the contribution from substructure may increase this difference to as much as a factor of 40.  Overall, these uncertainties combine to produce a total uncertainty of more than an order of magnitude in the flux multiplier. A conservative model, at the bottom of this range, consists of an NFW profile, with a ST mass function and a mass limit of $10^6\\,M_{\\odot}$. For this model, the present-day flux multiplier is $6\\times 10^5$. We estimate that a more likely model has a profile with more mass in its inner regions, and lower limit to the mass function of $10^5\\,M_{\\odot}$. For this model, the present-day flux multiplier is $5\\times 10^6$. For an extreme model, with a Moore profile limited only by annihilation, and a mass limit close to the Jeans mass at recombination, the present-day flux multiplier will be $\\simeq 1\\times 10^8$.  It is worth noting that with these large flux multipliers, some neutralino candidates can already be ruled out. The 86 GeV neutralino considered in Bergstr\\\"{o}m, Edsj\\\"{o}, \\& Ullio (2001), for instance, produced almost 10\\% of the gamma-ray background observed by EGRET at 1 GeV, assuming a flux multiplier of $\\simeq 2\\times 10^6$, while a 166 GeV neutralino produced more than 50\\% of the flux observed at 10 GeV. Both these candidates are thus possible in our most conservative model, while the higher-energy candidate is marginally excluded in our favoured model, and both are ruled out in our extreme model."
    },
    "0207/astro-ph0207250_arXiv.txt": {
        "abstract": "The Magellan active optics system has been operating continuously on the Baade 6.5-m since the start of science operations in February 2001.  The active optical elements include the primary mirror, with 104 actuators, and the secondary mirror, with 5 positional degrees of freedom. Shack-Hartmann (SH) wavefront sensors are an integral part of the dual probe guiders. The probes function interchangeably, with either probe capable of guiding or wavefront sensing.  In the course of most routine observing stars brighter than 17th magnitude are used to apply corrections once or twice per minute.  The rms radius determined from roughly 250 SH spots typically ranges between 0\\farcs05 and 0\\farcs10. The spot pattern is analyzed in terms of a mixture of 3 Zernike polynomials (used to correct the secondary focus and decollimation) and 12 bending modes of the primary mirror (used to compensate for residual thermal and gravitational distortions).  Zernike focus and the lowest order circularly symmetric bending mode, known affectionately as the ``conemode,\" are sufficiently non-degenerate that they can be solved for and corrected separately. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207585_arXiv.txt": {
        "abstract": "The abundance of holmium (\\emph{Z} = 67) in the Sun remains uncertain.  The photospheric abundance, based on lines of Ho II, has been reported as $+0.26 \\pm 0.16$ (on the usual scale where log(H) = 12.00), while the meteoretic value is $+0.51 \\pm 0.02$.  Cowan code calculations have been undertaken to improve the partition function for this ion by including important contributions from unobserved levels arising from the (4f$^{11}$6p + 4f$^{10}$(5d + 6s)$^{2}$) group.  Based on 6994 computed energy levels, the partition function for Ho II is 67.41 for a temperature of 6000 K.  This is $\\approx$1.5 times larger than the value derived from the 49 published levels.  The new partition function alone leads to an increase in the solar abundance of Ho to log(Ho) $= +0.43$.  This is within 0.08 dex of the meteoretic abundance.  Support for this result has been obtained through LTE spectrum synthesis calculations of a previously unidentified weak line at $\\lambda_{\\odot}3416.38$.  Attributing the feature to Ho II, the observations may be fitted with log(Ho) = +0.53.  This calculation assumes log$(gf) = 0.25$ and is uncertain by at least 0.1 dex. ",
        "introduction": "The abundances of most nonvolatile elements in CI chondrites are in close agreement with those determined for the solar photosphere (Grevesse and Sauval, 1998; hereafter GS).  The lanthanide holmium is one of only a few elements with differences of the order of 0.3 dex or more.  GS give log(Ho$_{\\mathrm{CI}}$)$-$log(Ho$_{\\odot}$) $= +0.25$.  The solar abundance of holmium reported by GS is $+0.26 \\pm 0.16$ and is based on an unpublished analysis of Ho II, the dominant ion in the Sun, by Daems, Bi\\'{e}mont, and Grevesse (hereafter DBG) in 1984.  Four lines at $\\lambda$3343.6, 3398.9, 3456.0, and 3474.2 {\\AA} were analyzed.  All of the features are blended, with the holmium lines constituting only minor contributors to the observed spectrum (N. Grevesse, private communication).  The scatter about the mean abundance in this study arises partly from problems of blending and partly from the uncertainty in the adopted $gf$-values from Gorshkov and Komarovskii (1979). The partition function for Ho II was calculated using the known atomic levels at the time; at 6000 K, the value was found to be 43.8 (N. Grevesse, private communciation).  The published level structure for Ho II is seriously incomplete (cf. Martin, Zalubas, and Hagan, 1978; hereafter MZH): only 49 levels are known, many with no term designations and a few with multiple J-values.  Wyart, Koot, and van Kleef (1974) note that a satisfactory analysis of this ion requires the calculation of the (4f$^{11}$6p + 4f$^{10}$(5d + 6s)$^{2}$) group.  Of particular importance are the contributions from the 4f$^{10}$6s$^2$ and 4f$^{10}$5d6s configurations which are expected to begin below 12 000 cm$^{-1}$ (Brewer, 1971).  Only eight measured levels have energies lying below this threshold.  Much of the difference between the meteoritic abundance of holmium and that found by DBG in the Sun may be due to an underestimate of the partition function for Ho II. Cowan-code calculations were therefore undertaken to supplement the published data.  In addition, a fresh examination of the solar spectrum was made in an attempt to identify additional, weak, \\emph{unblended} features that might be attributable to Ho II and used in spectrum synthesis calculations to establish an independent estimate of the abundance of this element.  One line in particular, at 3416.38 {\\AA,} was found to be amenable to such analysis.  In the following section, we describe the energy level calculations and give the partition functions that result from them for temperatures in the range 3000 to 34 000 K.  The third section addresses our search for other Ho II lines suitable for analysis by spectrum synthesis techniques, while in the fourth, the results of our study of the $\\lambda$3416 line are presented.  The final section of the paper briefly summarizes the current state of our knowledge of the abundances of the lanthanide rare-earth elements in the Sun vis-a-vis the CI chondrites; it also includes some remarks about our recent re-examination of the abundance of the volatile element indium in the Sun. ",
        "conclusions": "Uncertainties in the spectrum synthesis not withstanding, the results presented herein lend credence to the view that the abundance of holmium in the Sun follows the pattern of other rare-earth elements in showing general agreement with the meteorites. In Figure 4 we plot the logarithmic CI abundances from the compilation of McDonough and Sun (1995) minus the solar abundances from GS (with some small modifications to reflect recent work on the light elements by Holweger (2001)). Because McDonough and Sun have not been involved with the reconciliation of solar and meteoritic abundances, they may be considered an independent source for the meteoritic values, which are, in any case, very close to those reported by GS.  (For example, McDonough and Sun give the logarithmic abundance of holmium in the meteorites as +0.49, 0.02 dex lower than the value quoted by GS but within the latter's stated error.)  The abundance differences are plotted versus elemental condensation temperatures taken from Lodders and Fegley (1998).  \\begin{figure} \\resizebox{10cm}{!}{\\rotatebox{270}{\\includegraphics{figure4.eps}}} \\vspace{1.5cm} \\caption{ Logarithmic CI minus solar photospheric abundances vs. elemental condensation temperatures.  Points are indicated by the chemical symbols of the elements. The position of holmium prior to and after the present work is shown by the heavy arrow. Recent adjustments in the abundances of lutetium and terbium are shown by the dashed arrows.  Among the heavy and rare-earth elements, only tungsten~(W) remains seriously discordant.} \\end{figure}  As may be seen in Figure 4, with recent revisions to the solar abundances of lutetium (Bord, Cowley, and Mirijanian, 1998), terbium (Lawler, \\emph{et al.}, 2001) and now holmium, among the heavy and rare-earth elements, only tungsten remains seriously anomalous.  The authors are currently investigating the prospects of being able to improve the photospheric abundance of this element using the techniques described in this paper.  Among the volatile elements, indium ($Z = 49$) presents the greatest discordance.  We have recently re-evaluated the solar abundance of this element in a manner similar to that described in this paper (cf. Bord and Cowley, 2001).  Specifically, we synthesized a 2.25 \\AA~segment of the solar spectrum centered on the position of the resonance line of In I at $\\lambda4511.3$.  The hfs constants were taken from Jackson (1957, 1958), and the VALD log($gf$) = $-0.210$ was adopted.  Our best fit to the solar data yielded log(In) = 1.56.  This is about 0.1 dex smaller than the value found by Lambert, Mallia, and Warner (1969), the difference arising almost entirely from revisions in the oscillator strength of the line.  We estimate that the formal errors in the computation are at least $\\pm0.2$ dex, half of which is due to residual uncertainties in the oscillator strength.  The remainder of the error is contributed by uncertainties associated with the placement of the continuum which amount to about 1\\% in this region.  We have thus confirmed that the Sun is overabundant in indium by a factor of more than 5, with log(In$_{\\mathrm{CI}}$)$-$log(In$_{\\odot}$) $= -0.74$.  This difference is reflected in Figure 4 and remains a challenge to our understanding.  With no other unblended indium lines available, further refinements in the solar abundance of this element appear remote.  Moreover, the meteoritic abundance rests on 24 separate analyses made at three different labs and is uncertain by only 6\\% (Anders and Ebihara, 1982).  Given that indium should be one of the last of the chalcophiles to leave the gas phase, conditions in the solar nebula may have led to its incomplete condensation and to a depletion of this species in the CI meteorites.  If this were to be the case, one might expect to see a similar depletion of cadmium ($Z = 48$) in these meteorites relative to the Sun.  However, the cadmium abundance in the CI meteorites is within 0.01 dex of that found for the solar photosphere (Youssef, D\\\"{o}nszelmann, and Grevesse, 1990), so the mystery persists.  These exceptions not withstanding, it is quite remarkable - indeed, puzzling perhaps -  the degree to which the abundances of the CI meteorites agree with those of the solar photosphere.  Given the complex chemical histories of these meteorites in which aqueous alteration has likely played a significant role, it is surprising that the bulk atomic compositions of these bodies seem to have remained largely unfractionated.  Continuing efforts to refine the abundances of trace elements in the Sun to permit more precise comparisons between the abundance patterns in our star and these meteorites hold promise for providing greater insight into the history of both."
    },
    "0207/astro-ph0207517_arXiv.txt": {
        "abstract": "Considerable work has been devoted to the question of how best to parameterize the properties of dark energy, in particular its equation of state $w$. We argue that, in the absence of a compelling model for dark energy, the parameterizations of functions about which we have no prior knowledge, such as $w(z)$, should be determined by the {\\it data} rather than by our ingrained beliefs or familiar series expansions. We find the complete basis of orthonormal eigenfunctions in which the principal components (weights of $w(z)$) that are determined most accurately are separated from those determined most poorly. Furthermore, we show that keeping a few of the best-measured modes can be an effective way of obtaining information about $w(z)$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207498_arXiv.txt": {
        "abstract": "We present Infrared Space Observatory (ISO) observations of $^{16}$OH and $^{18}$OH toward Sgr B2 with a spectral resolution of $\\sim$35  km s$^{-1}$. The OH J=5/2$\\rightarrow$3/2 and J=3/2$\\rightarrow$1/2 rotational lines of the $^2\\Pi_{1/2}$ ladder are seen in emission while the cross-ladder transitions (from the   $^2\\Pi_{3/2}$ J=3/2  to the J=1/2, 3/2 and 5/2 levels of the $^2\\Pi_{1/2}$ ladder), and the   $^2\\Pi_{3/2}$ J=5/2$\\leftarrow$3/2 and J=7/2$\\leftarrow$5/2 lines are detected in absorption. The $^{18}$OH  $^2\\Pi_{3/2}$ J=5/2$\\leftarrow$3/2 $\\Lambda$--doublet at $\\sim$120 $\\mu$m is also observed in absorption. All  OH $\\Lambda$--doublets are resolved (except the $\\sim$98 $\\mu$m) and show, in addition to the strong absorption at the velocity of Sgr B2, several velocity components associated to the gas surrounding Sgr B2 and to the foreground clouds along the line of sight.  No asymmetries in the line intensities of each doublet have been observed. We have modeled the observations using a non-local radiative transfer code and found that the OH absorption/emission must arise in a shell around Sgr B2 not resolved by the ISO/LWS beam. The gas density is moderate, with upper limits of 10$^4$ cm$^{-3}$ and $\\simeq$300 K in temperature. The  OH abundance is high, (2--5)$\\times$10$^{-6}$. We argue that a widespread photon dominated region explains the enhancement of OH abundance. ",
        "introduction": "The Sagittarius  B2 complex represents an interesting  burst of massive star formation in the Galactic Center (GC) and may be representative of other active galactic nuclei. Large scale continuum maps from radio to far--IR wavelengths show that Sgr B2 is the brightest and among the most massive clouds of the GC (Cox \\& Laureijs 1989, Pierce--Price et al. 2000; Scoville, Solomon \\& Penzias 1975). The Sgr B2 M source is the brightest far--IR condensation of the  complex with a diameter of $\\simeq$40$^{\\prime\\prime}$ at 100 $\\mu$m (Goldsmith et al. 1992) and has also the largest  gas column density (Lis \\& Goldsmith 1989). These studies have shown that the core is embedded in an extended and clumpy cloud of warm gas. However, the observed rich chemistry in the gas in front of Sgr B2 and its possible heating  mechanisms are far from  settled. Low velocity shocks have been invoked to explain the enhanced abundances of some species such as NH$_2$ or NH$_3$ (Flower, Pineau des For\\^{e}ts \\& Walmsley 1995) which are not observed in more quiescent regions. On the other hand, UV radiation can have an important effect on the gas in the outer layers. This radiation may be provided by evolved stars and by young massive stars in the envelope of Sgr B2 itself (Mart\\'{\\i}n--Pintado et al. 1999). In addition, the emission of several ions with high excitation potential (such as [OIII] and [NIII]) is extended in the Sgr B2 region (Goicoechea et al. 2003, in preparation). Thus, a widespread ionized component producing photodissociation regions (PDRs) in the envelope is also possible. In any of these scenarios or in a combination of both, O--bearing species such as H$_2$O, OH, H$_3$O$^+$ and atomic oxygen are decisive for the thermal balance of the dense molecular gas (Neufeld et al. 1995). In particular, the hydroxyl radical, OH, has been predicted to be abundant in both scenarios.   The structure of the OH levels is depicted in Figure 1. The fundamental rotational $\\Lambda$--doublet ($\\sim$119 $\\mu$m) was first detected in the direction of Sgr B2 in absorption against the thermal dust emission by Storey, Watson \\& Townes (1981). Each line of the doublet is so optically thick that absorbs completely the continuum radiation avoiding any check with the models. These lines are detected over a large area (9$^\\prime\\times$27$^\\prime$) around Sgr B2 (Cernicharo et al. 1999). A detailed study of the far-IR spectrum of OH has been carried out mainly in the Orion shocked region (Watson et al. 1985; Viscuso et al. 1985; Melnick et al. 1987, 1990; Betz \\& Boreiko 1989) but little is known toward GC clouds such as the gas in the line of sight toward Sgr B2.  From the theoretical point of view, Offer \\& van Dishoeck (1992; hereafter OfD) and Offer, van Hemert \\& van Dishoeck (1994; hereafter OfHD) have made detailed calculations of the collisional cross sections between OH and H$_2$ and of the expected intensity of the OH far--IR lines. They predicted strong asymmetries in the intensity of the OH $\\Lambda$--doublets due to important asymmetries in the collisional rates between OH and para-H$_2$.    In this letter, we present high--resolution observations of several OH $\\Lambda$--doublets toward Sgr B2 involving levels up to $\\sim$420 K. All  lines inside the $^2\\Pi_{1/2}$ ladder are seen in emission while all the remaining lines are detected in absorption. The data have been modeled using a non-local radiative transfer code in order to constrain some important physical parameters of the warm outer layers of Sgr B2 and of the foreground gas. ",
        "conclusions": "The typical OH abundance in dense quiescent clouds is (0.1-1)$\\times$ 10$^{-7}$ with OH/H$_2$O ratios in the range 1 to 10$^{-2}$ (Bergin et al. 1995). According to the chemical models, the major contribution to the enhanced  OH abundance comes from regions where water vapor is being rapidly photodissociated (Sternberg \\& Dalgarno 1995) or regions in which shock waves play a role (Draine et al. 1983). In the case of Sgr B2, the innermost regions of the cloud are completely hidden in the mid and far-infrared due to the large dust opacity. Almost all the observed far--IR OH  come from the external layers of the Sgr B2 cloud, i.e., from the surrounding gas at high T$_K$ and from the cold dark clouds and diffuse gas in the line of sight.   In order to estimate the OH abundance and the physical conditions leading to the observed OH line emission in the V$_{LSR}\\simeq+50$  km s$^{-1}$ component, we have modeled the first 20 OH rotational levels using the code developed by Gonz\\'alez-Alfonso \\& Cernicharo (1993). The hyperfine structure of OH is not included in the model. The collisional cross sections of OfHD have been used (kindly provided in electronic form by E.F van Dishoeck). We have adopted a spherical geometry for a cloud consisting of two components: a uniform core with a diameter of 25$^{\\prime\\prime}$ (3.2$\\times$10$^{18}$ cm for an assumed distance to Sgr B2 of 8.5 kpc), with a dust temperature of 30 K (see Goicoechea \\& Cernicharo, 2001) and opacity at 80 $\\mu$m of 5 ($\\tau_{\\lambda}=\\tau_{80}*(80/\\lambda({\\mu}m)$), and a shell of variable thickness and distance to the core. The shell was divided into 14 layers. All molecular transitions in the core are thermalized to the dust temperature due to the large dust opacity in the far--IR. In order to check the sensitivity of the results on the physical parameters, $n(H_2)$, T$_K$ and $N(OH)$ were varied from 10$^3$ to 10$^5$ cm$^{-3}$, 40 to 600 K, and 1$\\times$10$^{15}$ up to 1$\\times$10$^{17}$ cm$^{-2}$ respectively.   The observations put some constraints on the size and density of the shell. Excitation temperatures for the cross-ladder and $^2\\Pi_{3/2}$ transitions have to be below the dust temperature in order to reproduce the observed line absorptions. In addition, the shell thickness and its distance to the core cannot be large compared to the core size, otherwise limb effects introduce important re-emission that could cancel the absorption produced by the gas in front of the core. The results for some  models are reproduced in Figure 2. We have found that if the shell is placed at some distance from the core, then it is difficult to reproduce the observations as almost all lines appear in emission within the LWS beam (80$^{\\prime\\prime}$). Even if the shell is contiguous to the core but large in thickness, lines  also appear in emission. Models in Fig. 2 correspond to a total size of 42$^{\\prime\\prime}$ (5.3$\\times$10$^{18}$ cm) for the core+shell cloud.  The model labelled M$_1$ corresponds to T$_K$=40 K, $n(H_2)$=10$^4$ cm$^{-3}$ and $\\chi(OH)=2\\times10^{-6}$. The cross-ladder and $^2\\Pi_{3/2}$ transitions are in absorption while the $^2\\Pi_{1/2}$ intra-ladder transitions are predicted in emission. However, although the intensity of the absorbing lines is found to fit reasonably well the observations, emission lines are underestimated. If the OH abundance  is increased in order to reproduce the intensity of the emitting lines then, absorbing lines are poorly fitted as limb effects start to dominate. This effect is already  seen in the J=1/2$\\leftarrow$3/2 cross-ladder transition at $\\sim$79 $\\mu$m. In particular, the $^2\\Pi_{1/2}$ J=5/2$\\rightarrow$3/2  transition at $\\sim$98 $\\mu$m is strongly underestimated by model M$_1$. If the gas temperature increases, collisions start to pump the OH levels and the intensity of the $\\sim$98 $\\mu$m  line is better fitted. However, a new problem arises as strong asymmetries in the line intensities of each $\\Lambda$-doubling line do appear at high temperature  (no asymmetries have been observed within the S/N ratio of the data ). These asymmetries were already predicted by OfD and are due to a strong parity change propensity rule introduced by collisions with para--H$_2$. They can only be suppressed if the H$_2$ ortho--to--para (OTP) ratio is higher or if radiative pumping dominates. We have adopted a OTP ratio of 3. Model M$_2$ in Fig. 2 corresponds to T$_K$=150 K with otherwise the same parameters as M$_1$. The intensity of the  $\\sim$98 $\\mu$m doublet is considerably enhanced with respect to the low temperature model. The intensity asymmetry in this doublet is not visible in the data as the doublet is unresolved with the LWS/FP spectrometer. The effect of the IR pumping strongly reduces the asymmetries that could be present without the optically thick core. Finally, model M$_3$ (thick lines in Figure  2) represents T$_K$=300 K, $n(H_2)$=5$\\times$10$^3$ cm$^{-3}$  and $\\chi(OH)=3\\times10^{-6}$. Again, the strong asymmetry of the  $\\sim$98 $\\mu$m doublet is not  observable due to the lack of spectral resolution. However, the observed total intensity for the doublet agrees well with that predicted from the model. For the other lines, M$_3$ reproduces qualitatively well the observed OH pattern, except in the line wings in the 119 and 79 $\\mu$m doublets. Nevertheless, cold foreground gas (not included in the model) absorbs (like in M$_1$) at these wavelengths, reducing the  emission wing effects in these lines.  From our models it is clear that the gas surrounding Sgr B2 has to be warm (see Ceccarelli et al. 2002) and must have a moderate H$_2$ density. The OH abundance is rather high, (2--5)$\\times$10$^{-6}$. Assuming that M$_3$ represents a reasonable fit to the OH data, we derive $N($$^{18}$$OH)$=(6$\\pm$2)$\\times$10$^{13}$ cm$^{-2}$ for Sgr B2. Hence, the $^{16}$O/$^{18}$O isotopic ratio is 240--280, in excellent agreement with the value obtained by Bujarrabal, Cernicharo \\& Gu\\'elin (1983) from the OH 18 cm lines and a factor $\\simeq$2 lower than in the V$_{LSR}$ $\\simeq$0  km s$^{-1}$ component. Higher T$_K$ could even be possible but the actual collisional rates  introduce very strong intensity asymmetries in the cross-ladder transitions which are not observed. If the gas temperature is higher, then its density has to  be lower than the value used in M$_3$ (5$\\times$10$^3$ cm$^{-3}$). The reason is that higher temperatures produce strong emission in the OH high excitation transitions ($\\sim$71 and $\\sim$65 $\\mu$m for example, see Figure 2) which is not observed. For T$_K$=600 K, $n(H_2)$ has to be decreased to (1-2)$\\times$10$^3$ cm$^{-3}$ and $\\chi(OH)$ increased to (0.5-1)$\\times$10$^{-5}$ to obtain a crude fit the ISO data. A temperature gradient across the shell, going from T$_K$=40 K in the innermost regions to T$_K$=600 K in the external layers is probably a better representation to the physical structure of the edge of the Sgr B2 cloud.  The inferred $^{16}$OH column density in the shell, (1.5--2.5)$\\times$10$^{16}$ cm$^{-2}$, and that of H$_2^{16}$O determined also from LWS/FP data (Cernicharo et al. 1997; 2002, in preparation) leads to an OH/H$_2$O=0.1--1 abundance ratio. For comparison, shock models of Flower et al. predict $\\simeq$$10^{-3}$, while models of dense dark cores predict $\\simeq$10$^{-4}$. On the other hand, the maximun OH abundance in PDR models is $\\simeq$10$^{-5}$ with OH/H$_2$O ratios close to 10. The large abundance found in Sgr B2 suggests that its external shells are illuminated by a strong UV field producing a PDR, although low velocity shocks (v$_S\\sim$30 km $s^{-1}$), unresolved by the ISO data,  may be also present. Many species formed during the evolution of the cold gas in Sgr B2 are being now reprocessed in these regions. This warm gas is poorly traced by millimeter and submillimeter observations, but it represents the strongest contribution to the absorption/emission features in the far--IR spectrum of Sgr B2. In addition to $^{16}$OH and $^{18}$OH, we have searched for several related species such as  OH$^+$ or H$_2$O$^+$. No lines from these species have been found.  The future far--IR instruments on board next generation telescopes such as the {\\it{Herschel Space Observatory}} will allow a fast mapping  of many OH rotational lines in galactic and extragalactic sources making OH a very useful tool in deriving the physical properties of GC giant molecular clouds. This could help in the analysis of starburst galaxies like Arp 220."
    },
    "0207/astro-ph0207384_arXiv.txt": {
        "abstract": "{\\small We have studied time variability in the flux from the flat spectrum synchrotron radiation of the Blandford \\& K\\\"{o}nigl (1979) model for relativistic, conical jets. The resulting model has been applied to the flux variation of the flat spectrum of GRS 1915+105 observed by Fender \\& Pooley (2000). This comparison has highlighted a fundamental problem of the Blandford \\& K\\\"{o}nigl (1979) model in that it requires unphysically large electron densities to explain the flux levels observed from the flat spectrum of GRS 1915+105.} ",
        "introduction": "To explain the infrared and millimetre emission oscillations from GRS 1915+105, we have adopted the Blandford \\& K\\\"{o}nigl (1979, hereafter BK79)\\cite{BK79} model of a relativistic jet, containing power-law electron distributions. Adiabatic expansion gives the jet a conical shape with an electron density and magnetic field strength that decreases with radius. The jet emits partially self-absorbed synchrotron emission over a radial range covering several orders of magnitude, and hence from a wide range of optical depths. Therefore, the total jet emission consists of the summation self-absorbed synchrotron spectra from regions of the jet of decreasing optical depth (or increasing radius), and hence electrons at each radius, from the base to the top of the jet, emit a spectrum that is progressively shifted towards lower frequencies. This summation results in a `flat' region to the spectrum, which has a slope of zero if energy losses due to adiabatic decompression are ignored (as in the BK79 model). However, if some energy losses do occur then this region has a positive, inverted slope, the magnitude of which depends upon the magnitude of the energy loss.  The extent of the flat spectral region, $\\nu_{\\rm max}/\\nu_{\\rm min}$, is solely determined by the emission region size, which in our model is fixed by the jet velocity and time-scale of variability. At the lower frequency end of the flat spectrum there is a smooth transition to an optically thick synchrotron spectrum with $F_{\\nu} \\propto \\nu^{5/2}$, and the high frequency end smoothly terminates into an optically thin, $F_{\\nu} \\propto \\nu^{-5/8}$, synchrotron spectrum.  Full detail of the derivation of our model may be found in Collins, Kaiser \\& Cox (2002)\\cite{Collins02}. The emission spectrum predicted by this model is given in terms of the optical depth function, $\\tau_{\\nu}(r)$. \\begin{equation} \\label{eFN} F_{\\nu} = \\frac{8.8 \\times 10^{-18}}{({D_{\\rm j}}/ \\mbox{pc})^2} \\; b_0 \\left(\\frac{\\nu}{\\rm GHz}\\right)^{5/2} \\nonumber \\int_{r_{\\rm min}}^{r_{\\rm max}}\\left(\\frac{r}{r_0}\\right)^{3/2} \\left[1 - \\mathrm{e}^{-\\tau_{\\nu}(r)}\\right] \\,\\mathrm{d}r \\mbox{ mJy}, \\end{equation} where, \\begin{equation} \\label{eDepth} \\tau_{\\nu}(r) = 1.5 \\times 10^{-12} a_0 \\left(\\frac{r}{r_0}\\right)^{-25/8}         \\left(\\frac{\\nu}{\\rm GHz}\\right)^{-25/8}, \\end{equation} $D_{\\rm j}$ is the distance to the jet, $r_{\\rm max} / r_{\\rm min}$ defines the extent of the emission region, and all parameters have SI units unless stated otherwise.  The model has two parameters. The first, which we have named $a_0$, controls the optical depth as a function of radius, $\\tau_{\\nu}(r)$, and thus shifts the flat region of the spectrum to higher or lower frequencies. The second model parameter, which we have named $b_0$, only effects the flux normalisation. The $a_0$ parameter also has an effect upon the flux normalisation, and therefore we first fit the $a_0$ parameter to the spectral shape, before fitting $b_0$ to the flux normalisation. The model parameter, $a_0$, is related to the physical parameters of the system by $a_0 = w(r\\!=\\!r_0)k(r\\!=\\!r_0)B(r\\!=\\!r_0)^{17/8}$, where $k$ is the normalisation value of the electron density distribution, and $B$ is the magnetic field strength. The $b_0$ parameter is defined as $b_0 = w(r\\!=\\!r_0)B(r\\!=\\!r_0)^{1/2}$. The half-width of the jet $w(r\\!=\\!r_0)$ is fixed by observational constraints, and therefore from $a_0$ and $b_0$, we may determine $k(r\\!=\\!r_0)$ and $B(r\\!=\\!r_0)$.  To implement time dependence in this model as a variation in the flux level of the flat spectrum it would be simplest to make $b_0$ a function of time with $a_0$ fixed. However, physically this would require the electron density to vary by a process that exactly compensates variations in the magnetic field strength. Therefore we choose to make $a_0$ a function of time, and fix $b_0$, which may be simply interpreted, physically, as a variation to the electron density injected into the jet. Time dependence is then included into equation \\ref{eFN} by defining $r_{\\rm max} = v_{\\rm j} t$, where $v_{\\rm j}$ is the bulk velocity of the jet material, and by increasing the injected value of $k_0$ with time according to a Gaussian function.  From this model we can predict three observable parameters to which our model parameters may be fit; the time lag between the peak fluxes at each frequency of observation, the flux ratio of these peak fluxes, and the flux normalisation value. The value of the $a_0$ model parameter affects all three of these observables. The flux ratio is determined by the position of the flat spectral region with respect to the two frequencies of observation, which is defined by $a_0$. The time lags are also determined by the position of the flat spectral region, because the optically thin end of the flat spectrum is formed at small radii, whereas the optically thick end is formed at much larger radii. Therefore the closer that both frequencies of observation are to the optically thin end,  the smaller the time lag. ",
        "conclusions": ""
    },
    "0207/astro-ph0207667_arXiv.txt": {
        "abstract": "A detailed analysis of the very cool white dwarfs \\sdss\\ and \\lhs\\ is presented. Model atmosphere calculations with improved collision-induced absorptions by molecular hydrogen indicate that a pure hydrogen composition can be ruled out, and that the strong infrared absorption observed in these cool stars is better explained in terms of collisions of \\htwo\\ with neutral helium. It is shown that even though the overall shape of the observed energy distributions can be reproduced reasonably well with helium-rich models, the peak of the energy distribution near 6000 \\AA\\ is always predicted too narrow. The extreme helium-rich composition inferred for both objects is discussed in the broader context of the extremely cool white dwarfs reported in various surveys. ",
        "introduction": "The photometric and spectroscopic analyses of \\citet{brl97,blr01} have revealed that the coolest white dwarfs observed in the Galactic disk have effective temperatures in excess of $\\Te\\sim4000$~K \\citep[see, e.g., Fig.~21 of][]{blr01}. The absence of any cooler white dwarfs can be interpreted as the result of the finite age of the disk, which has been estimated from a determination of the luminosity function at 8 $\\pm$ 1.5 Gyr \\citep{lrb98}. This value is also consistent with the location of white dwarfs in a mass versus effective temperature diagram in which theoretical isochrones are overplotted \\citep[see Figs.~24 and 25 of][]{blr01}. Since the Galactic halo is believed to have formed many Gyr earlier than the disk, white dwarfs associated with the halo could be considerably older and thus cooler than those found in the disk. In addition, halo white dwarfs could be easily identified by their peculiar kinematics. This was first recognized by \\citet{ldm89} who identified 6 white dwarfs in their luminosity function sample that had tangential velocities consistent with a halo population ($v_{\\rm tan}\\ \\gta250$ km~s$^{-1}$). Several objects from that small sample are now believed to be too young (they are too warm and massive) to belong to the Galactic halo, however \\citep{fon01}.  More recently, white dwarfs with effective temperatures below 4000~K have been identified in various surveys \\citep{hambly99,harris99,ibata00,harris01,opp01a,opp01b,ruiz01,scholz02,farihi02}, many of which are believed to be associated with the halo population. Even though the model fluxes are successful at reproducing in detail the optical and infrared broadband photometric observations of white dwarfs above 4000~K \\citep{brl97,blr01}, the observed energy distributions of cooler white dwarfs are at odds with the predictions of model atmosphere calculations. For instance, the analysis of \\lhs\\ by \\citet{harris99} has shown that the observed optical and infrared energy distribution could not be reproduced adequately in terms of a pure hydrogen atmosphere, or a mixed hydrogen and helium composition. Similar conclusions were reached by \\citet{opp01b}. For WD 0346+246 (also analyzed by Oppenheimer et al.), \\citet{ber01} was even forced to introduce an {\\it ad hoc} source of opacity to reproduce the observed photometry at $B$, $V$, and $R$.  \\sdss\\ is another extremely cool white dwarf discovered by \\citet{harris01} in imaging data from the Sloan Digital Sky Survey. The optical and near-infrared spectrum of this object resemble that of \\lhs\\ \\citep[see Fig.~3 of][]{harris01}. Photometric observations reported by Harris et al.~covered only the optical $BVRI$ and a lower limit on the infrared $J$ magnitude. In this paper, we report new infrared photometric observations at $J$ and $H$ for \\sdss, and provide a thorough analysis of this object as well as of \\lhs, its almost identical twin. Both of these objects exhibit the strong infrared flux deficiency that results from collision-induced absorptions by molecular hydrogen. New calculations of this important source of opacity by Borysow and collaborators has led us to take a fresh look at the atmospheric properties of cool white dwarf atmospheres. In this paper, we thus explore in detail the effects of effective temperature, surface gravity, and atmospheric composition (hydrogen, helium, and heavier elements) on the predicted fluxes, and compare these emergent flux distributions with those of \\sdss\\ and \\lhs. ",
        "conclusions": "The extreme helium-rich composition inferred in our analysis for both \\sdss\\ and \\lhs\\ poses an obvious and challenging problem. Because the outer layers of cool white dwarfs are strongly convective, the hydrogen and helium chemical composition tends to be more or less homogeneous throughout the mixed convection zone. Since below $\\Te\\sim 12,000$~K, the mass of the deep helium convection zone is almost constant at $M_{\\rm He-conv}\\sim 10^{-6}~M_{\\ast}$ \\citep{tassoul90}, the small hydrogen abundances of only $\\nh\\sim10^{-5}-10^{-4}$ derived in our analysis imply a {\\it total} hydrogen mass of $\\sim5\\times10^{-12}$ \\msun. If this amount of hydrogen has been accreted from the interstellar medium over a cooling age of roughly 10 Gyr for a 3500~K white dwarf, the implied accretion rate would be only $\\sim 10^{-22}$ \\msun\\ yr$^{-1}$, a value that is completely unrealistic. For instance, the theoretical estimates of \\citet{wes79} suggest time-averaged accretion rates of $\\sim10^{-17}$ \\msun\\ yr$^{-1}$. Of course, the extreme helium-rich compositions derived here for \\sdss\\ and \\lhs\\ cannot be completely ruled out on the basis of these arguments alone since the problem of the accretion of hydrogen in cool white dwarfs is a long standing one, and no satisfactory explanation has yet been proposed to account for the persistence of helium-rich white dwarfs at low effective temperatures.  We note in this context that WD 0346+246 shown in Figure \\ref{fg:f9} has been interpreted by \\citet{opp01b} as a $\\Te=3750$~K white dwarf with an extremely small hydrogen abundance of $\\log \\nh=-6.4$. A reanalysis of this object by \\citet{ber01} with the improved \\htwo-He collision-induced opacities of \\citet{jorgensen} indicates an even lower hydrogen abundance of $\\nh\\sim10^{-9}$. Because this solution was improbable --- but not impossible --- from accretion considerations, \\citet{ber01} proposed an alternative solution with a higher hydrogen abundance of $\\nh\\sim0.8$, but was also forced to introduce in the model calculations a bound-free opacity from the Lyman edge associated with the so-called dissolved atomic levels of the hydrogen atom, in order to reduce the near-UV flux and match the optical photometric observations.  This additional pseudo-continuum opacity has not been included in the present calculations as the UV flux is already predicted too low for \\sdss\\ and \\lhs.  While it is probably safe to conclude that these cool white dwarfs have helium-rich atmospheres and effective temperatures below 4000~K, it is not yet possible to determine their atmospheric parameters and ages with great precision since the models fail to reproduce the photometric observations in detail.  The large parameter space explored in our analysis suggests that the source of this discrepancy lies in the physics included in our model atmospheres, which is either inadequate or incomplete. One avenue of investigation worth considering is non-ideal effects of the equation-of-state at the high atmospheric pressures that characterize helium-rich atmospheres. So far, these effects have only been estimated in pure hydrogen or pure helium atmospheres \\citep[][this paper]{bsw95,saumon99,rohrmann02}. It has also been demonstrated recently by \\citet{iglesias02} that the helium free-free opacity and Rayleigh scattering at high densities typical of those encountered here are seriously overestimated. The use of their improved treatment in model atmosphere calculations would most likely increase the flux where the collision-induced opacity is less important, i.e.~for wavelengths below 0.7 \\micron, precisely in the region where we observe the largest discrepancy. This could even solve the overluminosity problem of \\lhs\\ altogether as the calculated absolute visual magnitude for helium-rich white dwarfs would be brighter.  All white dwarfs below $\\Te=4000$~K for which a detailed analysis has been performed have been explained in terms of some mixed hydrogen and helium atmospheric compositions, or even a pure helium composition \\citep[][]{harris99,opp01b,ber01,farihi02}. There are still no white dwarfs below $\\Te=4000$~K that have been successfully and convincingly explained in terms of a pure hydrogen atmospheric composition. This result may not be completely unexpected since helium-atmosphere white dwarfs with their lower opacities have cooling time scales that can be considerably shorter than their hydrogen-atmosphere counterparts \\citep[see, e.g.,][]{hansen99, blr01}. Hence for a given population and at a given mass, the coolest white dwarfs are expected to have helium-dominated atmospheres. For instance, a cool helium-rich white dwarf at $\\Te=3250$~K (i.e., our temperature estimate for \\sdss) with an average mass of $\\sim 0.7$ \\msun\\ \\citep[see Fig.~22 of][]{blr01} has a cooling age of only 9.4 Gyr according to our evolutionary models with thin hydrogen layers \\citep[see also Fig.~24 of][]{blr01}; by comparison, a similar object with a thick hydrogen layer would have an age of 11.6 Gyr. This 9.4 Gyr estimate is entirely consistent with the age of the oldest objects in the Galactic disk analyzed by \\citet{blr01}, and the so-called ultracool white dwarfs reported in the literature may not be terribly old after all. Although WD 0346$+$246 has kinematics supporting halo membership \\citep{hambly99}, \\lhs\\ has a lower tangential velocity consistent with membership of either disk or halo, and \\sdss\\ has an even lower velocity inconsistent with halo membership \\citep{harris01}.  Despite the fact that we do not fully understand the accretion mechanism in cool white dwarfs, helium-atmosphere white dwarfs must eventually accrete ``some amount'' of hydrogen, at which point they will exhibit a strong infrared flux deficiency, even if the amount of hydrogen is extremely small according to our calculations. \\sdss\\ and \\lhs\\ are most likely such objects, and extremely cool hydrogen-atmosphere white dwarfs belonging to the old halo population of the Galaxy are yet to be identified."
    },
    "0207/astro-ph0207101_arXiv.txt": {
        "abstract": "{\\thanks{The work presented here is based in part on data obtained with the ESO facilities on La Silla (EFOSC/3.6-m)}\\ \\thanks{Figures 7 and 8 of this paper are available only in the electronic Journal.} We present results of a new, large survey for high-redshift radio-loud quasars, which targets quasars with $z>4$.  The survey is based on the PMN and NVSS radio surveys, optically identified using digitised UKST B, R and I plates.  Six new $z>4$ flat-spectrum QSOs have been discovered, and one previously known $z>4$ QSO rediscovered, based on their red optical colours.  The QSOs discovered in this survey are bright in both radio and optical bands; in particular PMN J1451-1512 ($z=4.763$, $\\rm I=17.3$, $R=19.1$) and PMN J0324-2918 ($z=4.630$, $\\rm R=18.7$) are very luminous.  PMN J1451-1512 at $z=4.763$ is also now the most distant radio-selected quasar.  In addition, 9 new quasars with $3.5<z<4.0$ were discovered during the survey. We present spectra and finding charts for the new quasars. We also derive a surface density of $\\rm 1.0\\pm 0.4\\times 10^{-3} sq\\ deg^{-1}$ for flat-spectrum $z>4$ QSOs with $\\rm S\\ge 72$mJy and $\\rm R< 21$mag. ",
        "introduction": "Radio selection remains one of the most efficient ways of finding high-redshift AGN. This approach has the further advantage of being less prone to selection effects than optical selection, since radio emission is unaffected by either intrinsic or extrinsic absorption due to dust. The specific aim of this work was to find optically bright, radio-selected high-redshift quasars. These can be used for unbiased studies of damped Lyman alpha systems at high redshift and other follow-up studies such as searches for associated high-redshift galaxy clusters.  We therefore began to carry out a large, systematic survey aimed specifically at $z>4$ QSOs. Our method involves the optical identification of flat-spectrum radio sources and the spectroscopic follow-up of the red stellar identifications. This approach exploits the fact that quasars at high redshift have redder optical colours than their low-redshift counterparts due to absorption by intervening HI (see figure 1 in Hook {\\it et al.}  1995), and has proved successful at finding high-redshift quasars in the past (Hook et al 1995, 1996, 1998).  Previous work using well-defined quasar samples has shown that $z>4$ radio-loud quasars are likely to be rare objects, both because the quasar population as a whole appears to decline at redshifts above $\\sim 2-3$ (Kennefick et al. 1996; Hawkins \\& Veron 1996, Schmidt, Schneider \\& Gunn 1995; Warren, Hewett \\& Osmer 1995) and because radio-loud quasars represent only about 10\\% of the full quasar population.  Specific studies of the radio-loud quasar population have shown that these objects are indeed rare at $z>4$.  Dunlop \\& Peacock (1990) presented strong evidence for a drop in the space density of radio-loud quasars between $z=2$ and $z\\approx 3$ based on radio-selected samples reaching $\\rm S_{2.7GHz} = 100mJy$.  More recently, significant progress has been made towards understanding the evolution of the radio-loud quasar population out to $z\\sim 4$ and the potential effects of absorption by dust, by the study of a completely identified, large area, flat-spectrum radio sample with $\\rm S_{2.7GHz}\\ge250mJy$ (Shaver et al. 1996; Wall et al., in preparation).  The low numbers of high-reshift quasars found in these studies demonstrates that there is a distinct drop-off in the space density of quasars at $z>3$. Thus for our new survey to be successful it must reach fainter radio flux density limits than the above surveys (to sample further down the luminosity function), and cover a significant fraction of the sky.  Here we present the first results of this new survey for high-redshift radio-loud quasars. As will be seen in section 2, the survey uses deeper radio and optical data than previous radio-loud quasar surveys, and covers a very large area in the Southern sky ($\\sim 7,500$ sq deg, comparable to that of the planned 10,000 sq deg of the Sloan survey). Our survey has produced seven $z>4$ flat-spectrum quasars, one of which was previously known. The survey complements the survey of Snellen et al (2001) which contains four flat-spectrum $z>4$ in the Northern sky, selected using a a similar method. ",
        "conclusions": "By using radio and optical multicolour data covering a significant fraction of the sky, we have produced a sample of high-redshift, optically bright radio-loud quasars.   Since this new quasar sample is well-defined, we can use it to estimate the surface density of $z>4$ quasars. To do this we consider objects with radio flux densities above the brightest limit of the PMN surveys, 72mJy. The SGC region of our survey has a high completeness of spectroscopic follow up (only two objects, or 1\\%, were not observed in this region) and also had a high completeness of the NVSS survey (0.955). There are four $z>4$ QSOs in the SGC region of the complete sample, which implies a surface density of $\\rm 0.92\\pm 0.5\\times 10^{-3} sq\\ deg^{-1}$ for $z>4$ QSOs with $\\rm S\\ge 72mJy$.  If the whole survey is considered, and the completeness for the SGC and NGC are taken as 94\\% and 61\\% respectively (as implied by the number of spectroscopically observed candidates in each region combined with the NVSS completeness factors derived in section 2) then the derived surface density is $\\rm 1.0\\pm 0.4\\times 10^{-3} sq\\ deg^{-1}$. This is fully consistent with the value determined from the SGC alone.  This value is similar to that of 1 per 1600sq deg ($6.3\\times 10^{-4}$) found by Snellen et al (2001) using a similar technique although with a slightly different radio and optical flux density limits ($\\rm R= 20$ compared to R=21 for the current survey, $\\rm S\\ge 30mJy$ compared to $\\rm S>72mJy$, an upper redshift limit defined by the red optical filter of $z\\sim4.5$ rather than $\\sim4.7$, and a radio spectral index cut at $-0.35$ rather than $-0.5$).  There were no quasars found in our survey with $z>4.76$ despite the fact that bright quasars with 4.9\\lapprox {\\it z}\\lapprox 6.3 should have been detectable on the I-plates. Allowing for the 86\\% completeness of spectroscopic follow up of the I-band sample, the effective area covered was 2930 sq deg. Therefore we derive an upper limit of $\\rm 3.4\\times 10^{-4} sq\\ deg^{-1}$ for the surface density of flat spectrum quasars with $\\rm I<19.5$, $\\rm S_{5GHz}\\ge 25mJy$ and 4.9\\lapprox {\\it z}\\lapprox 6.3.  Finally, the new sample of quasars presented in this paper represent some of the most luminous objects in the Universe and may also represent extreme peaks in the matter density distribution at high-redshift.  They are therefore ideal targets for various follow-up programs such as high-redshift absorption line studies and searches for associated high-redshift clusters."
    },
    "0207/astro-ph0207337_arXiv.txt": {
        "abstract": "The presence on the Virgo cluster outskirts of spiral galaxies with gas deficiencies as strong as those of the inner galaxies stripped by the intracluster medium has led us to explore the possibility that some of these peripheral objects are not newcomers. A dynamical model for the collapse and rebound of spherical shells under the point mass and radial flow approximations has been developed to account for the amplitude of the motions in the Virgo I cluster (VIC) region. According to our analysis, it is not unfeasible that galaxies far from the cluster, including those in a gas-deficient group well to its background, went through its core a few Gyr ago. The implications would be: (1) that the majority of the \\hi-deficient spirals in the VIC region might have been deprived of their neutral hydrogen by interactions with the hot intracluster medium; and (2) that objects spending a long time outside the cluster cores might keep the gas deficient status without altering their morphology. ",
        "introduction": "\\label{introduction}  A recent characterization of the large-scale 3D distribution of the neutral gas (\\hi) deficiency around the Virgo I Cluster (VIC) region by \\citet*[hereafter Paper~I]{Sol02} has shown that there are a significant number of galaxies with a dearth of atomic hydrogen at large Virgocentric distances. These peripheral gas-deficient objects, which can be observed both in the cluster front and in a probable background group well behind the cluster core, show gaseous deficiencies comparable in strength to those measured in the centers of Virgo and other rich galaxy clusters.  One of the mechanisms that can most naturally account for the observed reduction in the interstellar gas content of cluster galaxies is the ram pressure ablation caused by the rapid motion of galaxies through the dense intracluster medium (ICM). There is now compelling evidence for the decisive participation of this process in the gaseous deficiencies of spirals located in the centers of rich clusters, either directly from observations \\citep*{GH85,GJ87,DG91} ---including the discovery of shrunken gaseous disks \\citep{Cay94,Bra00} and the finding that \\hi\\ deficient spirals are on very eccentric orbits \\citep{Sol01}--- or from theoretical studies that have checked the efficiency of this mechanism \\citep*[][to name only a few]{SAP99,QMB00,Vol01}.  In the outer cluster regions the low density of the intergalactic medium calls, in principle, for alternative gas removal mechanisms, such as gravitational tidal interactions. Indeed, observational evidence suggests that processes of this kind might have played an important role on the evolution of the galactic population in distant clusters \\citep*[e.g.,][]{vDok99} due to favorable conditions for frequent low-relative velocity encounters among the galaxies in early epochs. Although in the VIC region low-relative-velocity galaxy-galaxy interactions may also be responsible for the gaseous deficiencies observed in some of the peripheral galaxies, it should not be forgotten that the dynamics of the Virgo region is dominated by large-scale non-Hubble radial streaming motions. In this context, it is plausible that some galaxies at large Virgocentric distances are on very eccentric orbits that carry them right through the cluster center with high relative velocities and are therefore liable to have suffered a strong interaction with the ICM.  One of the most influential studies of the Local Supercluster based on dynamical model calculations is the analysis by \\citet{TS84} of the infall of galaxies in the Virgo Southern Extension (or Virgo II cloud) toward the VIC. The lumpy distribution of galaxies in space led these authors to predict a very irregular infall rate which would be responsible for the secular evolution of the mix of morphological types in the cluster. It was suggested that the formation of the cluster took place at an early epoch ---when the universe was about one fourth of its present age \\citetext{R.~B.\\ Tully 2002, private communication}--- by a first generation of, probably, early-type galaxies. Afterwards, infall was reduced until very recently when the large spiral-rich Virgo II cloud has begun to fall into the cluster diluting the fraction of early-type systems. The fact that Tully \\& Shaya saw very few outwardly moving galaxies outside the $6\\degr$ VIC circle supported their argument that most, or perhaps all, spirals and irregulars in Virgo, mostly supplied by the Virgo II cloud, were recent arrivals.  All the findings of \\citet{TS84} were based on a data set that contained a limited number of galaxies with relatively uncertain distance estimates. Now, with a much larger sample and more accurate distances, we provide evidence that a substantial number of galaxies, with a wide range of clustercentric distances, are expanding away from Virgo. This suggests that these \\hfill galaxies \\hfill are \\hfill probably \\hfill reemerging \\hfill after \\hfill infall. If \\hfill our \\begin{center} \\vspace{-3mm} \\includegraphics[width=0.95\\columnwidth]{fig1.eps}% \\makeatletter\\def\\@captype{figure}\\makeatother \\figcaption{Voxel projection of the 3D distribution of \\hi\\ deficiency in the VIC region. The plot is in rectangular equatorial coordinates with distances given in Mpc. The xy-plane corresponds to $\\rm Decl.=0\\degr$, the x- and y-axis point to $\\rm R.A.=12$ and 18 hr, respectively, and the z-axis points to the north. The central dark spot is associated with the cluster, with M87 being right at its center. The other two enhancements are peripheral regions of \\hi\\ deficiency in the frontside and backside of the VIC. Our position is at the origin of the coordinate system.\\label{voxel}} \\end{center} interpretation of the situation is correct, then there is a more continuous influx of galaxies into Virgo than previously anticipated, so the hypothesis that the strong gaseous deficiencies currently seen in some peripheral objects were originated in a previous infall episode is worth exploring. ",
        "conclusions": ""
    },
    "0207/astro-ph0207271_arXiv.txt": {
        "abstract": "The variability in the infrared to millimetre emission from microquasar GRS 1915+105 is believed to be dominated by the system's relativistic jet. In this paper we develop a time-dependent version of the Blandford \\& K\\\"{o}nigl (1979) jet emission model and apply it to the oscillations in the infrared and millimetre emission from GRS 1915+105 observed by Fender \\& Pooley (2000). The resulting model provides a reasonable description of the observed flux oscillations from GRS 1915+105. From a fit of the observed time lag between the flux peaks in the infrared and millimetre emission together with the flux normalisation we were able to determine the model parameters for the GRS 1915+105 jet. We find that to achieve the observed flux levels with the model requires an unphysically large electron density within the jet. We therefore conclude that the Blandford \\& K\\\"{o}nigl (1979) model cannot explain these observations, either because it does not provide the correct description of the emission from microquasar jets, or because the observed emission variations do not originate in the jet. ",
        "introduction": "The microquasar GRS 1915+105 shows time variability in its emission at all observable wavelengths. Its high energy emission (X-rays) is believed to be created in the inner radii of the system's accretion disc, and its low energy (radio) emission is primarily from synchrotron emitting ejecta located at large radii ($r > 10^{11}$ m) within the system's jet. In this paper we investigate the cause of the variability in the infrared, and millimetre emission from GRS 1915+105 which is believed to originate in the inner regions of the relativistic jet.  Observations ranging from radio to infrared frequencies have revealed very large amplitude, quasi-periodic oscillations in the emission from GRS 1915+105 (Pooley \\& Fender 1997; Mirabel et al. 1998; Fender \\& Pooley 1998; Fender et al. 2002). In this paper we will concentrate on the oscillations detected at millimetre and infrared frequencies by Fender \\& Pooley (2000). These oscillations are characterised by the short time lags between flux peaks and a flux ratio, between these two frequencies, that is close to unity. This suggests a flat spectrum over a frequency range that encompasses three orders of magnitude. Although flat spectra had previously been observed from GRS 1915+105, these measurements were the first to suggest they extend to the infrared.  Flat spectra from compact radio sources such as microquasars have traditionally been explained in terms of partially self-absorbed synchrotron emission from the inner-most regions of a narrow, conical, relativistic jet (originally by Blandford \\& K\\\"{o}nigl 1979, hereafter referred to as BK79, but also by Hjellming \\& Johnston 1988, Falcke \\& Biermann 1995, Falcke 1996, and Falcke \\& Biermann 1999). This interpretation is supported by radio observations of blobs moving at apparent superluminal velocities outwards from a central elongated radio core (Mirabel \\& Rodr\\'{\\i}guez 1994; Rodr\\'{\\i}guez \\& Mirabel 1999; Fender et al. 1999; Dhawan, Mirabel \\& Rodr\\'{\\i}guez 2000). Detection of polarised emission strongly suggests the synchrotron nature of the emission (e.g. Fender et al. 2002). The temporal and spectral behaviour of the emission coming from the blobs is consistent with expectations from internal shocks in a relativistic jet flow (Kaiser, Sunyaev \\& Spruit 2000). These synchrotron-emitting blobs are therefore believed to be ejecta from a central compact, relativistic jet.  However, the BK79 model only represents a steady-state solution for emission from relativistic jets. Using the basic principles of the BK79 model we develop a time-dependent variation of this model in an attempt to explain the observed infrared and millimetre flux oscillations from GRS 1915+105. Whereas previous time-dependent jet emission models have only been concerned with the emission from the spherical ejecta blobs (e.g. van der Laan 1966, and Hjellming \\& Johnston 1988), the model presented in this paper concentrates solely upon the variability in the emission from the inner-most radii of the conical jet. In section \\ref{sBKModel} we demonstrate both how the steady-state BK79 model produces a flat synchrotron spectrum and how this is dependent upon the physical parameters of the jet. The model is then adapted to include time variability in section \\ref{sBKTEModel}, which allows a direct comparison between the observations and the model in section \\ref{sObs}. Finally we discuss the successes and failures of the model in explaining the observations in section \\ref{sDis}. ",
        "conclusions": "\\label{sDis}  \\subsection{Summary}\\label{ssSum}  In this paper we have demonstrated how the BK79 model produces a flat synchrotron spectrum from a conical jet geometry following the assumption that the emitting electron plasma is free of energy losses as it expands away from its origin. We have shown that the frequency range covered by this flat spectrum depends upon the radial extent of the emission region with respect to the radius where the electron plasma becomes optically thin to radiation at each frequency. This radius is determined by the value and radial dependence of the magnetic field strength and electron density within the jet. From observations of the frequency range that the flat spectrum extends over, and the observed flux from these frequencies, both the magnetic field strength and the electron density in the jet may be individually determined.  We also adapted the BK79 model to include time variability. This not only allows comparisons with light curve observations, and places stronger demands upon the hitherto unexplained re-energisation process, but also places more stringent constraints upon the model. Prior to the inclusion of time variability the extent of the emission region had to be assumed. The time scale of variability now sets the spatial extent of the emission region. Furthermore the model parameters may be derived from observation of the time lag between peaks at each frequency rather than the less reliable measure of the flux ratio between the peaks at each frequency. Time lags are caused by an optical depth effect that occurs whether or not the flat spectrum extends beyond the observed frequency range, and hence time lags are a more reliable indicator of the physics of the system than the flux ratio. The predicted time lag would be the same even with an inverted spectrum that fully covers the observed frequency range.  This time variability model was then applied to observations of the GRS 1915+105 microquasar, which revealed a discrepancy between the model light curve profile and the observed profile. Furthermore, the model calculated for the known parameters of GRS 1915+105 required extremely large electron densities to explain the observed flux. As we believe that the system parameters are sufficiently well constrained, the only conclusion must be that the model itself is flawed.  \\subsection{On the flux discrepancy}\\label{ssFlux}  As demonstrated in figure \\ref{f_dTvsFmm}, we have shown that the observed millimetre flux levels are only predicted by the models that also predict time lags, between the infrared and millimetre peaks, which are many times greater than those observed (for physically justifiable electron density values). We believe that the observed millimetre flux levels are an accurate representation of the true millimetre jet emission, with no significant background contamination, as the emission drops to zero between flux peaks. Therefore, the model is fundamentally incapable of predicting millimetre flux levels of the order of 100 mJy, together with time lags between flux peaks that are less than 100 seconds, which is what has indisputably been observed.  The bulk gas flow within the jet is relativistic. Therefore, if relativistic effects can cause the observed time lag to be considerably shorter than the time lag in the rest frame of the jet material, then the observed flux will be predicted from a lower electron density. Although time dilation has the opposite effect, relativistic Doppler shifts can in principle cause this effect. However, in the case of the GRS 1915+105 jet, the inclination of the jet axis to our line-of-sight is well constrained to $\\sim 70^{\\circ}$, and for such inclinations the maximum effect on the observed flux from relativistic Doppler shift and Doppler beaming is only 20 per cent. To explain the observations with believable electron densities the predicted flux needs to be $10^{4}$ times greater. Such a difference cannot easily be rectified by geometrical effects, nor by the current uncertainty in the distance to GRS 1915+105. Furthermore, the simplification in using the equivalent line-of-sight optical path length for a jet at an inclination of $90^{\\circ}$ does not significantly effect the results. Only for substantially smaller viewing angles and highly relativistic jet velocities do relativistic effects lead to a significant shortening of observed time lags compared to the jet rest frame.  The observations are almost certainly of synchrotron emission, as thermal emission would require even greater electron densities (Dhawan et al. 2000), and polarisation observations agree with this assertion (e.g. Fender et al. 2002). The flux discrepancy cannot be explained by the presence of the more extended jet beyond the region of the time variability, as the time variability still has to explain a large flux increase.  In conclusion it is impossible to explain synchrotron emission of this strength from a BK79 type jet with justifiable electron densities, when restricted to the observed time lags.  \\subsection{Is the flat spectrum flat?}\\label{ssFlat}  From a sample of just two frequencies we cannot conclusively say that the observed flux ratio is due to a flat spectrum which does not completely extend to the lower frequency. The other possibility is that the 'flat' region of the spectrum extends to cover both of the observed frequencies and has a slightly positive slope. Such a scenario would naturally arise if the energy losses in the jet due to adiabatic decompression were partially included. However, if adiabatic losses are fully included then the resulting inverted spectrum would have a slope that is significantly steeper than that observed. The steady-state jet model of Falcke (1996) includes partial energy losses within the jet, and, for the GRS 1915+105 jet inclination angle, predicts a spectral slope of $\\alpha \\approx 0.2$. The observations allow a spectral slope of $0.0 \\leq \\alpha \\leq 0.12$, due to the uncertainty on the infrared flux, so we cannot exclude the possibility of a fully extended, truly flat spectrum.  Hence, there are two possible scenarios to explain a non-unity flux ratio between the flux peaks of the two observed frequencies. Either the lower frequency observation is of the optically thick region (or the transition to this region) of the spectrum from a BK79 type jet, or it is of the 'flat' region of a jet with energy losses. There exists a simple observable difference between these two scenarios; whereas for the former case a relationship exists between the time lag and the flux ratio, for the latter case these two parameters are independent of each other. Although the loss of the relationship between the time lag and flux ratio will not in itself affect the predicted electron densities, its implication of energy losses within the jet could lead to a slightly lower predicted electron density.  The missing millimetre flux peak to coincide with the second observed infrared peak does seem to suggest that for the second ejection event the flat spectrum does not extend all the way down to the millimetre waveband. However, this explanation is applicable to either of the scenarios discussed above, and is perhaps more likely than a shift in the spectral slope.  \\subsection{Conclusion}\\label{ssConc}  We have developed a time-dependent version of the partially self-absorbed jet model of Blandford \\& K{\\\"o}nigl (1979). As expected from the original steady-state model, the time-dependent model gives rise to flat broadband spectra but cannot explain the large flux variations observed in GRS 1915+105 at millimetre and infrared frequencies without invoking unrealistically high electron densities. This result also holds true if the time variability is ignored. In this case the parameters of the steady-state model are determined through the observed flux ratio of the two frequencies at any given time, rather than from the observed time lag between flux peaks of the two frequencies. Therefore either this variable emission does not originate from the jet of GRS 1915+105, or an alternative model is required to explain microquasar jet emission.  Observations of the jet spectrum and time variability at different frequencies will help resolve the issue of whether the spectrum is flat, and whether we are observing optical depth effects. Observing a relationship between the flux ratio and the time lag would confirm that the flux ratio is due to an optical depth effect rather than due to a mildly inverted spectrum.  Recent observations of the infrared flaring behaviour of GRS 1915+105 (Rothstein \\&  Eikenberry 2002) have illustrated the complex nature of the time variability in the emission from relativistic jets. Hence, more complex models will ultimately be required to fully explain the observations."
    },
    "0207/astro-ph0207053_arXiv.txt": {
        "abstract": "s{There has never been a more exciting time in the overlapping areas of nuclear physics, particle physics and relativistic astrophysics than today. Orbiting observatories such as the Hubble Space Telescope, Rossi X-ray Timing Explorer (RXTE), Chandra X-ray satellite, and the X-ray Multi Mirror Mission (XMM) have extended our vision tremendously, allowing us to see vistas with an unprecedented clarity and angular resolution that previously were only imagined, enabling astrophysicists for the first time ever to perform detailed studies of large samples of galactic and extragalactic objects. On the Earth, radio telescopes (e.g., Arecibo, Green Bank, Parkes, VLA) and instruments using adaptive optics and other revolutionary techniques have exceeded previous expectations of what can be accomplished from the ground. The gravitational wave detectors LIGO, LISA VIRGO, and Geo-600 are opening up a window for the detection of gravitational waves emitted from compact stellar objects such as neutron stars and black holes.  Together with new experimental forefront facilities like ISAC, ORLaND and RIA, these detectors provide direct, quantitative physical insight into nucleosynthesis, supernova dynamics, accreting compact objects, cosmic-ray acceleration, and pair-production in high energy sources which reinforce the urgent need for a strong and continuous feedback from nuclear and particle theory and theoretical astrophysics.  In my lectures, I shall concentrate on three selected topics, which range from the behavior of superdense stellar matter, to general relativistic stellar models, to strange quark stars and possible signals of quark matter in neutron stars.} ",
        "introduction": "\\label{sec:intro}  A forefront area of research, both experimental and theoretical, concerns the exploration of the subatomic structure of superdense matter and the determination of the equation of state -- that is, the relation between pressure $P$, temperature $T$ and density $\\epsilon$ -- associated with such matter.\\cite{greiner96:wilderness}  Knowing its properties is of key importance for our understanding of the physics of the early universe, its evolution in time to the present day, compact stars, various astrophysical phenomena, and laboratory physics.  The high-temperature domain of the phase diagram of superdense matter is probed by relativistic heavy-ion colliders. Complementary to this, neutron stars contain cold superdense matter permanently in their centers (cf.\\ Fig.\\ \\ref{fig:cross}), which make them superb astrophysical laboratories for probing the low-density high-density domain of the phase diagram of superdense matter.\\cite{blandford92:b,weber99:book}  Neutron stars are dense, neutron-packed remnants of massive stars that blew apart in supernova explosions. They are typically about 10 kilometers across and spin rapidly, often making several hundred rotations per second.  The discovery rate of new rotating neutron stars, spotted as pulsars by radio telescopes \\cite{manchester77:a,lyne90:a}, is rather high. To date about 1400 pulsars are known. Depending on star mass and rotational frequency, gravity compresses the matter in the core regions of pulsars up to more than ten times ($\\sim 1.5$~GeV/fm$^3$) the density of \\begin{figure}[tb] \\begin{center} \\leavevmode \\psfig{figure=ns-crossection.ps,width=9cm,angle=-90} \\caption[]{Competing structures and novel phases of subatomic matter predicted by theory to make their appearance in the cores ($R\\lsim 8$~km) of neutron stars.\\protect{\\cite{weber99:book}}} \\label{fig:cross} \\end{center} \\end{figure} ordinary atomic nuclei, thus providing a high-pressure environment in which numerous subatomic particle processes plausibly compete with each other and novel phases of matter may exist. The most spectacular ones stretch from the generation of new baryonic particles (e.g., $\\Sigma, \\Lambda, \\Xi, \\Delta$) to quark ($u, d, s$) deconfinement to the formation of Boson condensates ($\\pi^-, K^-$, H-matter), as illustrated in Fig.\\ \\ref{fig:cross}. There are theoretical suggestions of even more exotic processes inside pulsars, such as the formation of absolutely stable quark matter, a configuration of matter even more stable than the most stable atomic nucleus, $^{56}\\rm{Fe}$! In the latter event, pulsars would be largely composed of pure quark matter, eventually enveloped in thin nuclear crusts (bottom-left portion in Fig.\\ \\ref{fig:cross}). No matter which physical processes are actually realized inside neutron stars, each one leads to fingerprints, some more pronounced than others though, in the observable stellar quantities. Paired with the unprecedented wealth of new observational pulsar data, it seems to be within reach for the first time ever to seriously explore the subatomic structure of matter in the high-density low-temperature portion of its phase diagram from observed pulsar data. To this aim Einstein's field equation of relativistic gravity, \\begin{eqnarray} G^{\\mu\\nu} \\equiv R^{\\mu\\nu} - {{1}\\over{2}} g^{\\mu\\nu} R = 8 \\, \\pi \\, T^{\\mu\\nu}(\\epsilon,P(\\epsilon)) \\, , \\label{eq:intro.1} \\end{eqnarray} is to be solved in combination with the latest theories of the subatomic structure of matter.\\cite{weber99:book} The latter follow according to the scheme \\begin{eqnarray} { {\\partial {\\cal L}(\\{\\phi\\}) }\\over{\\partial \\phi} } - \\partial_\\mu \\; { {\\partial {\\cal L}(\\{\\phi\\})}\\over{\\partial (\\partial_\\mu \\phi)} } = 0 \\quad \\Rightarrow \\quad P(\\epsilon) \\, , \\label{eq:intro.2} \\end{eqnarray} where ${\\cal L}(\\{\\phi\\})$ denotes a given stellar matter lagrangian.\\cite{weber99:book} In general, ${\\cal L}$ is a complicated functional of the numerous hadron and quark fields, collectively written as $\\{ \\phi \\}$, that acquire finite amplitudes up to the highest densities reached in the cores of neutron stars.  According to what has been said just above, plausible candidates for $\\phi$ are the charged states of the SU(3) baryon octet, $p, n, \\Sigma, \\Lambda, \\Xi$ \\cite{glen85:b}, the charged states of the $\\Delta$ \\cite{weber89:e,huber98:a}, $\\pi^-$ \\cite{umeda92:a} and $K^-$ \\cite{kaplan86:a,li97:a,li97:b,brown96:a,brown97:a} mesons, as well as $u, d, s$ quarks.\\cite{fritzsch73:a,baym78:a,kettner94:b} The conditions of chemical equilibrium and electric charge neutrality of stellar matter require the presence of leptons too, in which case $\\phi=e^-, \\mu^-$.  Theories of superdense matter enter Einstein's field equation (\\ref{eq:intro.1}) via the energy-momentum tensor $T^{\\mu\\nu}$, which contains the equation of state of the stellar matter, $P(\\epsilon)$.  Because of the rather uncertain behavior of the matter at supernuclear densities, the models derived for the equation of state differ considerably with from each other. This has its origin in various sources such as: (1) the many-body technique used for the determination of the equation of state, (2) the model adopted for the nucleon--nucleon force, (3) assumptions about the fundamental building blocks of neutron star matter, (4) the inclusion of boson condensates, and (5) the considerations of a possible phase transition of confined hadronic matter into deconfined quark matter.\\cite{weber99:book,heiselberg00:a}  In general Eqs.\\ (\\ref{eq:intro.1}) and (\\ref{eq:intro.2}) were to be solved simultaneously since the particles move in curved spacetime whose geometry, determined by Einstein's field equations, is coupled to the total mass energy $\\epsilon$ of the matter.  In the case of neutron stars, however, the long-range gravitational force can be cleanly separated from the short-range nuclear force so that Eqs.\\ (\\ref{eq:intro.1}) and (\\ref{eq:intro.2}) decouple from each other. This simplifies the study of compact stellar objects considerably. ",
        "conclusions": ""
    },
    "0207/astro-ph0207579_arXiv.txt": {
        "abstract": "High-resolution spectra covering the absorption features from interstellar C~I were recorded for four early-type stars with spectrographs on the {\\it Hubble Space Telescope}, in a program to measure the fine-structure excitation of this atom within neutral clouds inside or near the edge of the Local Bubble, a volume of hot ($T\\sim 10^6\\,$K) gas that emits soft x-rays and extends out to about 100~pc away from the Sun.  The excited levels of C~I are populated by collisions, and the ratio of excited atoms to those in the ground level give a measure of the local thermal pressure.  Absorptions from the two lowest levels of C~I were detected toward $\\alpha$~Del and $\\delta$~Cyg, while only marginal indications of excited C~I were obtained for $\\gamma$~Ori, and $\\lambda$~Lup.  Along with temperature limits derived by other means, the C~I fine-structure populations indicate that for the clouds in front of $\\gamma$~Ori, $\\delta$~Cyg and $\\alpha$~Del, $10^3<p/k<10^4\\,{\\rm cm}^{-3}\\,$K at about the $\\pm 1\\sigma$ confidence level in each case.  The results for $\\lambda$~Lup are not as well constrained, but still consistent with the other three stars.  The results indicate that the thermal pressures are below generally accepted estimates $p/k>10^4\\,{\\rm cm}^{-3}\\,$K for the Local Bubble, based on the strength of x-ray and EUV emission from the hot gas.  This inequality of pressure for these neutral clouds and their surroundings duplicates a condition that exists for the local, partly-ionized cloud that surrounds the Sun.  An appendix in the paper describes a direct method for determining and eliminating small spectral artifacts arising from variations of detector sensitivity with position. ",
        "introduction": "The Sun is immersed in two concentric volumes of interstellar material with very different properties.  In the immediate surroundings [out to distances ranging from 0.05 to several pc away  (Redfield \\& Linsky 2000)], the gas is partly ionized, with characteristic properties $0.10 < n({\\rm H~I}) < 0.24\\,{\\rm cm}^{-3}$, $0.04 < n(e) < 0.10\\,{\\rm cm}^{-3}$ and $T\\approx 7500\\,$K  (Lallement 1998; Puyoo \\& Ben Jaffel 1998; Ferlet 1999; Gry \\& Jenkins 2001).  This cloud, often called the Local Interstellar Cloud (LIC), is surrounded by a large cavity containing fully-ionized, very hot, low-density material with $n({\\rm H}^+)\\approx 0.008\\,{\\rm cm}^{-3}$ and $T\\sim 10^6\\,$K  (Bergh\\\"ofer et al. 1998; Burrows \\& Guo 1998) that is prominent in soft x-ray emission (Snowden et al. 1990, 1998).  This large volume of hot gas is believed to be the product of probably several supernova explosions  (Cox \\& Reynolds 1987; Breitschwerdt \\& Schmutzler 1994; Ma\\'iz-Apell\\'aniz 2001; Smith \\& Cox 2001) and is known as the Local Bubble (LB).  The LB is also conspicuous by showing a general absence of cold, neutral gas up to a well defined boundary  (Sfeir et al. 1999; Vergely et al. 2001), as is shown in Figure~\\ref{sfeirplt}.  A persistent puzzle has been the apparent mismatch in thermal pressures between the two media.  Using the parameters stated above, one finds that $1400 < p/k < 3600\\,{\\rm cm}^{-3}\\,$K for the LIC, which contrasts with an apparent representative value of $16,000\\,{\\rm cm}^{-3}\\,$K for the fully-ionized, hot gas in the LB.  The objective of the research presented here is to help answer the question, ``What thermal pressures are found for other neutral clouds within the Local Bubble?  That is, do they have values similar to the LIC, or are they closer to matching the apparent pressure of the LB?''  To gain insights on this question, one can observe stars located behind individual clouds inside the LB and measure their absorption features of C~I arising from different fine-structure levels of the ground state.  The ratios of populations of these states are governed by an equilibrium between collisional excitations (governed by local densities and temperatures) and spontaneous radiative decays.  Stars suitable for viewing the C~I features had to satisfy a number of conditions to yield useful results.  Their selection is described in \\S\\ref{plan}; ultimately four such stars were observed in a manner described in \\S\\ref{obs}.  The C~I absorption features are very weak, and in order to obtain useful measures of their strengths particular care was exercised to remove spurious signals arising from the detector, as outlined in \\S\\ref{artifacts}.  (Mathematical details about the correction method are presented separately in the Appendix.) Section~\\ref{analysis} describes how the equivalent widths of various absorption features were combined and corrected for saturation (very mild, except for one of the stars).  This section also discusses the derivations of fine-structure population ratios, which may be compared to the theoretically expected values for different conditions (\\S\\ref{expected_f1}).  Before one can derive useful limits for the thermal pressures, the allowable ranges of temperature must be constrained, and different methods of deriving these constraints are discussed in \\S\\ref{temp_limits}.  Ultimately, the limits for the population ratios and temperatures restrict the possible values for $p/k$ [and local density $n({\\rm H})$], as shown for the four cases in the diagrams presented in Fig.~\\ref{limit_panels}.  Three out of the four stars indicate internal thermal pressures for the foreground clouds that are below the generally accepted value for the LB.  Possible explanations for this imbalance, duplicating that seen for LIC, are presented in \\S\\ref{summary}. ",
        "conclusions": "\\label{summary}  A synthesis of the conclusions presented earlier is shown in Figure~\\ref{limit_panels}, again in the representation of $\\log (p/k)$ vs. $\\log n({\\rm H})$.  To constrain the allowed physical conditions, the outcomes presented in Table~\\ref{col_dens} for the fine-structure population ratios of C~I, represented by the parameter $f1$, must be supplemented by intersecting lines representing temperature limits derived in \\S\\ref{temp_limits}.  Of these limits, the upper bounds derived from the measurements of the Na~I line widths by other investigators are the most direct and reliable.  However these measurements are valuable only for the cases of $\\gamma$~Ori and $\\delta$~Cyg.  Alternative means for limiting $T$, such as using the thermal equilibrium curve to define a lower bound and the carbon ionization equilibrium for an upper bound, are indirect and subject to less certain theoretical assumptions.  \\placefigure{limit_panels} \\begin{figure} \\plotone{f7.eps} \\caption{Allowed combinations of $\\log (p/k)$ and $\\log[n({\\rm H})]$ permitted by the $\\pm 1\\sigma$ error limits for $f1$ for the C~I-absorbing regions in front of the four targets.  Intersecting constraints arise from the thermal equilibrium curve for an external absorbing column $N({\\rm H~I})=10^{19}\\,{\\rm cm}^{-2}$ calculated by Wolfire, et al  (1995), and temperature limits defined either from the Doppler width of Na~I absorption lines or, alternatively, the requirement that the expected yield of ${\\rm C~I_{total}}$ relative to C~II from the condition of ionization equilibrium (Eq.~\\protect\\ref{C_ioniz_equilib}) is as least as much as actually observed.  General layouts of the expected C~I fine-structure population ratios and ${\\rm C~I_{\\rm total}/C~II}$ ionization equilibria are shown in Figs.~\\protect\\ref{f1_conts} and \\protect\\ref{c_conts}, respectively.\\label{limit_panels}} \\end{figure}  It should be noted that for each case the lines that border the allowed (shaded) regions do not enclose the worst possible deviations.  For instance, the $f1$ limits represent only $\\pm 1\\sigma$ deviations.  The temperature constraints from the width of the Na~I absorption lines do not include possible errors in measurement, since they are difficult to assess.  The theoretical limits based on thermal equilibria or the expected fractional abundances of carbon in neutral form are only as good as the assumptions that were incorporated into their respective developments.  In spite of these shortcomings, when three of the stars (those other than $\\lambda$~Lup) are examined collectively, they indicate that clouds within the Local Bubble have thermal pressures in the range $10^3<p/k<10^4\\,{\\rm cm}^{-3}\\,$K.  From the introductory remarks presented in \\S\\ref{intro}, it is clear that the most interesting limit on pressure is the upper one.  Any single case is not compelling enough to support the proposal that $p/k<10^4\\,{\\rm cm}^{-3}\\,$K, since deviations in excess of $1\\sigma$ in a particular direction can be expected to occur 16\\% of the time. Nevertheless, if one supposes that all clouds within the Local Bubble have a single representative pressure, it is difficult to imagine that all of the random observational errors would conspire to mislead us to believe that $p/k<10^4\\,{\\rm cm}^{-3}\\,$K when in reality the pressure is much higher than this value.  The result for $\\lambda$~Lup is not helpful in constraining $p$ on the high side, but at least it shows an outcome on the low side that is consistent with those of the other stars.  As for systematic errors, the only easily identifiable ones are possible errors in the relative $f$-values of the C~I and C~I* transitions or errors in the parameters used to develop the predicted values of $f1$, such as the collisional rate constants or spontaneous radiative decay rates.  It is apparent from this study that a seemingly large disparity in pressure between neutral clouds embedded in the Local Bubble and the surrounding hot gas is a general phenomenon and not just confined to the LIC around the Sun.  Possible solutions to the dilemma posed by the apparent pressure imbalance may arise from the following considerations: \\begin{enumerate} \\item Perhaps the clouds have not had enough time to respond fully to a sudden increase in external pressure that happened during the event(s) that created the LB, perhaps some $10^5-10^7$~yr ago  (Cox \\& Reynolds 1987).  If this is the case, each cloud is in an intermediate condition where a shock is still progressing from the periphery of the cloud, perhaps well outside the portion containing the C~I, to its center. However, this explanation is not satisfactory for the LIC, since the distance from the Sun to the closest boundary divided by the sound speed is less than $10^4$~yr. \\item Perhaps magnetic fields within the clouds provide the support needed to counteract the high pressure from the outside.  An upper limit $B<3\\mu$G for the magnetic field just outside the heliosphere (Gloeckler, Fisk, \\& Geiss 1997) diminishes the attractiveness of this explanation in the case of the LIC. \\item Additional contributions to the clouds' internal pressures might arise from turbulent motions.  However, the narrowness of the Na~I lines observed for $\\gamma$~Ori and $\\delta$~Cyg lowers the prospects for this proposal.  For instance, if one uses the limit $n({\\rm H})<10^{1.7}\\,{\\rm cm}^{-3}$ shown in Fig.~\\ref{limit_panels} and assumes the Na~I lines are broadened only by turbulence, one finds that $p_{\\rm turb}/k=\\rho b^2/2k < 470$ and $800\\,{\\rm cm}^{-3}\\,$K for $\\gamma$~Ori and $\\delta$~Cyg, respectively. \\item Perhaps the measurements of the thermal pressure for the hot gas in the Local Bubble were incorrect because they relied on either inaccurate models for the line emission or the assumption that the radiation from the hot plasma is emitted under equilibrium conditions. For a plasma that is cooling very rapidly, perhaps through adiabatic expansion, the x-rays may come principally from delayed recombinations in cooler gas, rather than collisional excitations at higher temperatures  (Breitschwerdt \\& Schmutzler 1994).  Apart from this theoretical proposal, there have also been persistent difficulties in reconciling measurements of x-ray and EUV line emission with models for a hot plasma in collisional equilibrium and with cosmic abundance ratios of the heavier elements  (Jelinsky, Vallerga, \\& Edelstein 1995; McCammon et al. 2002).  In trying to understand the soft x-ray spectrum measured by the Diffuse X-ray Spectrometer (DXS), Sanders et al.  (2001) even explored models with depleted element abundances and nonequilibrium conditions arising from either rapid heating or cooling.  From their inability to achieve reasonable fits of these more complex models with their data, they suggested that the discrepancies might arise from inaccurate atomic data which were used to predict the expected emission line strengths.  If these models are in need of revision, then one might question the earlier estimates of hot gas emission measures from low resolution surveys of x-ray emission that led to the finding that $p/k > 10^4\\,{\\rm cm}^{-3}$K within the LB. \\end{enumerate}"
    },
    "0207/astro-ph0207323_arXiv.txt": {
        "abstract": "{ As part of a project to compute improved atomic data for the spectral modeling of iron K lines, we report extensive calculations and comparisons of atomic data for K-vacancy states in Fe~{\\sc xxiv}. The data sets include: (i) energy levels, line wavelengths, radiative and Auger rates; (ii) inner-shell electron impact excitation rates and (iii) fine structure inner-shell photoionization cross sections. The calculations of energy levels and radiative and Auger rates have involved a detailed study of orbital representations, core relaxation, configuration interaction, relativistic corrections, cancellation effects and semi-empirical corrections. It is shown that a formal treatment of the Breit interaction is essential to render the important magnetic correlations that take part in the decay pathways of this ion. As a result, the accuracy of the present $A$-values is firmly ranked at better than 10\\% while that of the Auger rates at only 15\\%. The calculations of collisional excitation and photoionization cross sections take into account the effects of radiation and spectator Auger dampings. In the former, these effects cause significant attenuation of resonances leading to a good agreement with a simpler method where resonances are excluded. In the latter, resonances converging to the K threshold display symmetric profiles of constant width that causes edge smearing. ",
        "introduction": "The iron K lines are among the most interesting features in astronomical X-ray spectra. These lines appear in emission in almost all natural X-ray sources, they are located in a relatively unconfused spectral region and have a well-known plasma diagnostics potential. They were first reported in the rocket observations of the supernova remnant Cas A \\citep{ser73}, in X-ray binaries \\citep{san75,bec77}, and in clusters of galaxies \\citep{ser77}, the latter thus manifesting the presence of extragalactic nuclear processed material. Observations of the galactic black-hole candidate Cyg X-1 showed that the line strength varied according to the spectral state \\citep{bar85,mar93}, and \\citet{tan95} found that the Fe~K lines from Seyfert galaxies were relativistically broadened and redshifted which suggested their formation within a few gravitational radii of a black hole.  Recent improvements in the spectral capabilities and sensitivity of satellite-borne X-ray telescopes ({\\em Chandra}, {\\em XMM--Newton}) have promoted the role of Fe~K lines in diagnostics, a trend that will continue to grow with the launch of future instruments such as {\\em Astro-E2} and {\\em Constellation-X}. Such plasma diagnostics ultimately rely on the knowledge of the microphysics of line formation and hence on the accuracy of the atomic data.  In spite of the line identifications by \\citet{see86} in solar flare spectra and the laboratory measurements of \\citet{bei89,bei93}, \\citet{dec93} and \\citet{dec95,dec97}, the K-vacancy level structures of Fe ions remain incomplete as can be concluded from the critical compilation of \\citet{shi00}. With regards to the radiative and Auger rates, the highly ionized members of the isonuclear sequence, namely Fe~{\\sc xxv}--Fe~{\\sc xxi}, have received much attention, and the comparisons by \\citet{chen86} and \\citet{kat97} have brought about some degree of data assurance. For Fe ions with electron occupancies greater than 9, \\citet{jac80} and \\citet{jac86} have carried out central field calculations on the structure and widths of various inner-shell transitions, but these have not been subject to independent checks and do not meet current requirements of level-to-level data.  The spectral modeling of K lines also requires accurate knowledge of inner-shell electron impact excitation rates and, in the case of photoionized plasmas, of partial photoionization cross sections leaving the ion in photoexcited K-vacancy states. In this respect, \\citet{pal02} have shown that the K-threshold resonance behavior is dominated by radiation and Auger dampings which induce a smeared edge. Spectator Auger decay, the main contributor of the K-resonance width, has been completely ignored in most previous close-coupling calculations of high-energy continuum processes in Fe ions \\citep{ber97,donnelly00,ber01,ball01}.  The exception is the recent $R$-matrix computation of electron excitation rates of Li-like systems by \\citet{whi02} where it is demonstrated that Auger damping is important for low-temperature effective collision strengths.  The present report is the first in a project to systematically compute improved atomic data sets for the modeling of the Fe K spectra. The emphasis is both on accuracy and completeness. For this purpose we make use of several state-of-the-art atomic physics codes to deliver for the Fe isonuclear sequence: energy levels; wavelengths, radiative and Auger rates, electron impact excitation and photoionization cross sections. Particular attention is given to the process of assigning reliable accuracy rankings to the data sets produced. Specifically, in the present report we have approached the radiative and Auger decay manifold of the $n=2$ K-vacancy states of Fe~{\\sc xxiv} as a test case of the numerical methods and the relevance of the different physical effects. By detailed comparisons with previous work, it has become evident that there is room for improvement, and that an efficient strategy can be prescribed for the treatment of the whole Fe sequence. Furthermore, we also compute inner-shell electron impact excitation rates of Fe~{\\sc xxiv}, the total photoionization cross sections of Fe~{\\sc xxiii} and the partial components of the latter into the K-vacancy levels of Fe~{\\sc xxiv} where the relevant effects of radiative and Auger dampings are fully established. ",
        "conclusions": "As a start in a project to compute improved atomic data for the spectral modeling of Fe K lines, we have carried out extensive calculations and comparisons of atomic data for modeling of the K spectrum of Li-like Fe~{\\sc xxiv}. The data set includes energy levels, radiative and Auger rates, collision strengths, and total and partial photoionization cross sections. Primary aims have been to select an applicable computational platform and an efficient strategy to generate accurate and complete data sets for other ions of the first row of the Fe isonuclear sequence.  We have studied several physical effects, namely orbital representations, core relaxation, CI, relativistic corrections, cancellation, semi-empirical corrections, and the damping of resonances by radiative and spectator Auger decay. For an $N$-electron ion, we have found that the most realistic representation is to have different orbital bases for the K-vacancy states, on the one hand, and for the valence states of the $N$- and $(N-1)$-electron systems on the other. This is available in {\\sc hfr}, but most other codes use orthogonal orbital bases for computational efficiency. In the case the {\\sc autostructure}, which uses a distorted-wave approach to compute Auger rates, orbitals of the $(N-1)$-electron system must then be used. Core relaxation leads to increases in the radiative and Auger widths no larger than 10\\%.  Level coupling within the $n=2$ complex has been found to be key, thus seriously questioning the reliability of the atomic model adopted by \\citet{lem84}. CI from higher complexes contributes negligibly. Contributions from the two-body relativistic operators, both fine structure and non-fine structure, play a conspicuous role in the decay of K-vacancy states of this ion, particularly in the Auger pathways. Electron correlation could be then interpreted as being highly magnetic: bound--free spin--spin effects have been shown to be important within the $n=2$ complex and specially critical for the Auger decay of the metastable ${\\rm 1s2s2p}\\ ^4{\\rm P}^{\\rm o}_{5/2}$ state. This state is also shown to decay radiatively through forbidden M1 and M2 transitions, the former requiring a relativistic corrected transition operator to avoid errors in the line strength of several orders of magnitude. In this highly ionized magnetic scenario, computer programs that do not include a formal numerical implementation of the Breit interaction, or neglect it, have limited applicability. Such is the case of {\\sc bprm} and {\\sc hfr}. Some of the large discrepancies found for the smallest rates have been attributed to strong cancellation effects. Fine tuning has been found to be a useful option to attain high numerical accuracy, particularly for line identification and to render intersystem couplings that can be very sensitive to level separations.  In the light of the problems discussed above, none of the codes seems to be the platform of choice for the calculation of radiative and Auger rates. We therefore employ several computational platforms to treat inner-shell processes which has proven to be key in elucidating the physics involved, and has been used previously by COR and SAF and more recently by \\citet{sav02}. This approach has therefore been retained in our current calculations of other members of the Fe isonuclear sequence.  The present {\\sc autostructure} calculations are an independent validation and refinement of that performed in COR; the level of agreement found at the different stages confirms this assertion. The excellent accord also obtained with the radiative rates by SAF allows us to establish a firm ranking of 10\\% for the present $A$-values. On the other hand, the fairly large discrepancies with the SAF Auger rates are believed to be caused by their approximate treatment of the Breit interaction in terms of screening constants. We therefore rank the present autoionization data at better than 15\\%. We can also conclude by comparing with SAF that the attained precision for the K-vacancy level energies of  $\\pm 4$ eV is a representative lower bound for current numerical methods. This however implies fine tuning that relies on spectroscopic measurements. Since complete experimental level structures are not available for most systems, further experiments would be welcome.   Radiative and spectator Auger dampings have been found to be of fundamental importance in the calculation of K-shell photoionization and electron excitation processes. In the former, resonances converging to the K threshold acquire a peculiar behavior that leads to edge smearing which, as discussed by \\citet{pal02}, has diagnostic potential in astrophysical plasmas. With regards to the latter, resonances are practically washed out, thus simplifying target modeling or the choice of a suitable numerical approach. This assertion is supported by the good agreement (10\\%) of the present excitation rates with the Coulomb--Born--Exchange results of \\citet{goett84} and with those in $R$-matrix calculation by \\citet{whi02} who used a more refined target. We have also found that the ground state of Fe~{\\sc xxiii} is mainly photoionized to the ${\\rm 1s2s}^2~^2{\\rm S}_{1/2}$ K level of Fe~{\\sc xxiv} which rapidly autoionizes rather than fluoresces. Thus  K$\\alpha$ emission from a Fe Li-like ion is mainly the result of electron impact excitation and dielectronic recombination."
    },
    "0207/astro-ph0207609_arXiv.txt": {
        "abstract": "Based on a refined generic dynamical model, we investigate afterglows from jetted gamma-ray burst (GRB) remnants numerically. In the relativistic phase, the light curve break could marginally be seen. However, an obvious break does exist at the transition from the relativistic phase to the non-relativistic phase, which typically occurs at time 10 to 30 days. It is very interesting that the break is affected by many parameters, especially by the electron energy fraction ($\\xi_{\\rm e}$), and the magnetic energy fraction ($\\xi_{\\rm B}^2$). Implication of orphan afterglow surveys on GRB beaming is investigated. The possible existence of a kind of cylindrical jets is also discussed. ",
        "introduction": "Researches on afterglows from long gamma-ray bursts (GRBs) have shown that they are of cosmological origin. The standard fireball model, which incorporates internal shocks to explain the main bursts and external shocks to account for afterglows, becomes the most popular model (for recent reviews, see van Paradijs et al. 2000). Some GRBs localized by BeppoSAX satellite have implied isotropic energy release of more than $10^{54}$ ergs, leading many theorists to deduce that GRB radiation must be highly collimated (Castro-Tirado et al. 1999; Huang 2000; Halpern et al. 2000; Dai \\& Gou 2001; Dai \\& Cheng 2001; Gou et al. 2001; Ramirez-Ruiz \\& Lloyd-Ronning 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2002).  To differentiate a jet from an isotropic fireball, we must resort to the afterglow light curves. When the bulk Lorentz factor of a jet drops to $\\gamma < 1 / \\theta$, with $\\theta$ the half opening angle, the edge of the jet becomes visible, the light curve will steepen by $t^{-3/4}$. This is called the edge effect (M\\'{e}sz\\'{a}ros \\& Rees 1999; Panaitescu \\& M\\'{e}sz\\'{a}ros 1999; Kulkarni et al. 1999). In addition, the lateral expansion of a relativistic jet will make the break even more precipitous. So it is generally believed that afterglows from jetted GRBs are characterized by an obvious break in the light curve at {\\em the relativistic stage}.  In this talk, we use our refined dynamical model to study the jet effect on the afterglow light curves. The possible existence of cylindrical jets is also discussed.   \\begin{figure}[htb] \\begin{center} \\leavevmode \\centerline{ \\hbox{ \\hbox{} \\hspace{0.5in} \\epsfig{figure=vr.eps,width=2.38in,height=1.2in,angle=-90, bbllx=200pt, bblly=270pt, bburx=450pt, bbury=680pt} \\hspace{0.5in} \\epsfig{figure=gamma.eps,width=2.38in,height=1.2in,angle=-90, bbllx=230pt, bblly=350pt, bburx=330pt, bbury=530pt} }} \\begin{flushright} \\parbox[t]{2.9in} { \\caption {  Velocity vs. radius for an isotropic adiabatic fireball (Huang 2000). The dashed line is the familiar Sedov solution in the Newtonian phase. The dash-dotted line is drawn according to Eq.~(1), which differs from the dashed line markedly. The solid line corresponds to our refined model (i.e., Eq.~(2)), which is consistent with the Sedov solution (Huang 2000).   }} \\ \\hspace{.2in} \\ \\parbox[t]{2.9in} { \\caption { Evolution of $\\gamma$. The solid line corresponds to a jet with ``standard'' parameters. Other lines are drawn with only one parameter altered: the dashed line corresponds to $\\theta_0 = 0.1$, and the dash-dotted line corresponds to $n = 10^6$ cm$^{-3}$ (Huang et al. 2000c). }} \\end{flushright} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207115_arXiv.txt": {
        "abstract": "The technique of imaging atmospheric Cherenkov telescopes (IACTs) has proved to be an effective tool to register cosmic $\\gamma$-rays in the very high energy region. The high detection rate of the IACT technique and its capability to reconstruct accurately the air shower parameters make it attractive to use this technique for the study of the mass composition of cosmic rays. In this article we suggest a new approach to study the CR mass composition  in the energy region from 30 TeV/nucleus up to the \"knee\" region, i.e. up to a few PeV/nucleus, using an array of imaging telescopes  of a special architecture. This array consists of telescopes with a relatively small mirror size ($\\sim 10$~m$^2$) separated from each other by large distances ($\\sim 500$~m) and equipped by multichannel cameras with a modest pixel size ($0.3-0.5^{\\circ}$) and a  sufficiently large viewing angle ($6-7^{\\circ}$).  Compared to traditional IACT systems (like HEGRA, HESS or VERITAS) the IACT array considered in this study could provide a very large detection area (several km$^2$ or more). At the same time, it allows an accurate measurement of the energy of CR induced air showers (the energy resolution ranges within 25-35\\%) and an effective separation of air showers created by different nuclei. Particularly, it is possible to enrich air showers belonging to the nucleus group assigned for selection up to $\\sim 90\\%$ purity at a detection efficiency of 15-20\\% of such showers. ",
        "introduction": "Last years the technique of imaging atmospheric Cherenkov telescopes (IACTs) has proved to be an effective tool to register cosmic $\\gamma$-rays in the energy region above 100~GeV \\cite{Cawley,Hillas,Aharonian1}. In comparison to the satellite experiments this technique provides as much as a few magnitudes larger detection area. Compared to the particle detector arrays it comprises a low energy threshold of detected  air showers. In addition, IACT arrays  (like HEGRA \\cite{HEGRA}, HESS \\cite{HESS} or VERITAS \\cite{VERITAS}) provide the ability  of accurate reconstruction of air shower parameters such as the arrival direction, the core location and the primary energy {\\it both} for $\\gamma$-ray and cosmic ray (CR) induced air showers. All this makes it  attractive  to apply the IACT technique for the study of the CR mass composition.  In \\cite{Plyasheshnikov2} an approach was developed allowing to use the IACTs for study of the chemical composition of CRs in the primary energy region $\\ge$1~TeV/nucleus. Later this approach received an application in the analysis of the observational data of the HEGRA collaboration (see Ref.~\\cite{PhysRev}). In this study we also discuss possibilities of application of the IACT technique for an analysis of the CR mass composition. However, on the contrary to \\cite{Plyasheshnikov2}, we analyse here higher primary energies, i.e. the energy interval between a few dozen TeV/nucleus and the \"knee\" region where the energy spectrum steepens from differential spectral index $\\sim 2.7$ below $\\sim 3$ PeV to $\\sim 3.0$ above (see Fig.~1). This energy region becomes important for ground based observations hence the balloon and satellite born experiments run out of statistics here. The region around the \"knee\" is of special interest, because in all likelihood it contains contributions of different types of CR sources. Presumably this could be the case when the contribution to CR population made by the supernovae remnants of the Galaxy changes for sources of other origin (see e.g. Ref.~\\cite{Shibata}). Different types of sources and a rather smooth energy dependence in the \"knee\" region combined with the varying properties of interstellar medium permits sufficient freedom in selection of alternative acceleration models. In such a circumstance it is important to develop new methods enabling an accurate determination of CR chemical composition and capable of reconstructing the structure of energy spectra in the \"knee\" region for individual CR nucleus groups.  A new approach to the analysis of the CR mass composition developed in this study is based on a special kind of IACT array consisting of telescopes separated by distances considerably larger than for traditional IACT systems and equipped  by multichannel cameras with a wide field of view\\footnote{An idea to apply an IACT array with such an architecture for observations of $\\gamma$-ray sources in the multi-TeV energy region was first suggested in \\cite{Plyasheshnikov1}.}. We show that such an array could provide a huge detection area (up to several km$^2$) and, thus, could extend an interval of observed primary energies up to the \"knee\" region. A detailed study is undertaken to estimate the possibility of the array to measure the energy of individual air showers and to separate showers induced by different primary nuclei in the \"event by event\" mode, which is the natural way of data analysis for stereoscopic IACT observations. Besides, such an approach, being a realization of the so-called non parametric analysis \\cite{Chilingarian}, yields not only an estimate of the primary energy and mass composition, but also allows to specify the uncertainty of the results in a quantitative way. ",
        "conclusions": "We suggest a new approach to study the mass composition of cosmic rays in the energy region from several dozen TeV/nucleus up to the knee region, i.e. up to a few PeV/nucleus, using an array of imaging atmospheric Cherenkov telescopes with an architecture sufficiently different from traditional IACT systems (like HEGRA, HESS or VERITAS). This array consists of imaging  telescopes with a relatively small mirror size ($\\sim$ 10 m$^2$) separated from each other by large distances  ($\\sim$500~m) and equipped  by multichannel cameras with a  modest pixel size ($0.3-0.5^{\\circ}$) and a  rather large viewing angle ($6-7^{\\circ}$). A square cell consisting of four telescopes  is used  to investigate basic properties of the array.  It is shown that the cell provides considerably more effective detection of air showers initiated by heavy nuclei compared  with the traditional IACT systems. For example, the proportion  of nuclei  with the charge number $Z\\ge 10$ in the cell detection rate is ($\\sim$ 15\\%), whereas for traditional systems this proportion is less than 5\\%.  It is found that the approach considered here could provide a precise determination of primary energy of air showers. Its energy resolution is within 25-35\\% and depends only weakly on the air shower energy and the atomic number of primary nucleus.   To separate air showers created by  different primary nuclei with the IACT technique one needs image parameters sensitive to the variation of the nucleus atomic number. We show in this study that a number of  image parameters exhibits such a sensitivity. The $Width$,  $Conc$  and $H_{\\mathrm{max}}$ parameters are the most effective ones.  An effective technique exploiting a criterion similar to the distance in multiparameter space has been developed to classify air showers detected by the cell on the basis of differences in the image parameters. This technique allows to enrich the residual detection rate of the cell (which includes events satisfying to the selection criteria) up to 90\\% of air showers belonging to the nucleus group assigned for selection saving about 15-20\\% of these showers.   The technique considered here could provide us with a sufficiently good statistics of observations everywhere in the energy region 30-1000~TeV. For example, an IACT array with a modest number of cells (16) could provide detection in the energy region above $\\sim 200$~TeV about 2000 nuclei with $Z\\ge 10$ for a modest observation time  (100~h). This number is considerably greater than the total number of such nuclei have been accumulated until now in the balloon and satellite born experiments.    Not having the ability of precise measurement of the primary particle mass as in satellite or balloon borne experiments, the proposed system nevertheless comprises capability of separation of primaries owing to which it could be an instrument complementary to ground-based particle detectors. To have confident results the number of predefined nuclei groups in ground-based experiments is typically limited to 2-3 for the cases of event by event analysis (HEGRA~\\cite{HEGRA_light}, KASCADE \\cite{Vardanyan}, \\cite{Roth}). The higher numbers are available only for techniques utilizing the information on the registered spectra of secondaries (four groups reported in~\\cite{Ulrich}). According to \\cite{HEGRA_light} and \\cite{Vardanyan} the proposed system would deliver the rather higher efficiency of proton selection than that for the HEGRA and the KASCADE for the time being. For the first project the purity $\\eta_{\\mathrm{(P, He)}}=0.97$ for (P, He) nuclei group at $\\kappa_\\mathrm{(P, He)}=0.3$ is reported. In our case such a value of $\\kappa$ yields $\\eta \\simeq 0.85$ for protons only, thus successfully rejecting group of He nucleus which comprises $\\simeq 40\\%$ of the (P, He) group. For the KASCADE $\\kappa_\\mathrm{(P, He)}=0.7$ corresponds to $\\eta_\\mathrm{(P, He)} \\simeq 0.9$. Only for protons our approach yields $\\eta_\\mathrm{P} \\simeq 0.65$. Comparison of our data on selection of HVH group with that for the KASCADE (\\cite{Vardanyan} and \\cite{Roth}) shows that the proposed system is inferior to the KASCADE unless $\\kappa_\\mathrm{HVH} \\lesssim 0.5$. At the same time for HVH and LM groups the selection capabilities of the KASCADE deteriorates with the decrease of primary energy \\cite{Roth} so that our approach would benefit in the sub-knee energy region.  \\begin{ack} We are grateful to A. Akhperjanian for supplying us with numerical results concerning the mirror aberrations of imaging telescopes. \\end{ack}"
    },
    "0207/astro-ph0207359_arXiv.txt": {
        "abstract": "{\\small We discuss the possibility that microquasar jets may be powerful emitters of TeV neutrinos.  We estimate the neutrino fluxes produced by photopion production in the jets of a sample of identified microquasars and microquasar candidates, for which available data enables rough determination of the jet parameters. We demonstrate that in several of the sources considered, the neutrino flux at Earth, produced in events similar to those observed, can exceed the detection threshold of a km$^2$ neutrino detector. Sources with bulk Lorentz factors larger than those characteristic of the sample considered here, directed along our line of sight may be very difficult to resolve at radio wavelengths and hence may be difficult to identify as microqusar candidates. However these sources can  be identified through their neutrino and gamma-ray emission.} ",
        "introduction": "The composition of microquasars jets is yet an open issue. The synchrotron emission both in the radio and in the IR is consistent with near equipartition between electrons and magnetic field, which is also implied by minimum energy considerations \\cite{Levinson96}. However, the dominant energy carrier in the jet is presently unknown (with the exception of the jet in SS433). A possible diagnostic of hadronic jets is emission of TeV neutrinos \\cite{Levinson01}.  As shown in reference \\cite{Levinson01}, for typical microquasar jet parameters, protons may be accelerated in the jet to energies in excess of $\\sim10^{16}$~eV.  The interaction of these protons with synchrotron photons emitted by thermal electrons is expected to lead to 1--100~TeV neutrino emission. The predicted fluxes are detectable by large, km$^2$-scale effective area, high-energy neutrino telescopes, such as the operating south pole detector AMANDA \\cite{Andres00} and its planned 1~km$^2$ extension IceCube \\cite{icecube}, or the Mediterranean sea detectors under construction (ANTARES \\cite{antares};  NESTOR \\cite{nestor}) and planning (NEMO  \\cite{nemo} and \\cite{Halzen01} for a recent review).  In this paper we consider a class of identified Galactic microquasars with either persistent jets or documented outbursts. For each source we provide, for illustrative purposes, our model prediction for the neutrino flux that should have been emitted during particular events, using radio data available in the literature. Although the temporal behavior of many of these sources may be unpredictable, we demonstrate that some of the sources could have been detected by a neutrino telescope with effective area larger than km$^2$ (in some cases even 0.1 km$^2$) had such a detector been in operation during the time of the recorded events and, therefore, propose that they should be potential targets for the planned neutrino telescopes.  In addition we consider a list of XRBs thus far unresolved at radio wavelengths, that are believed to be microquasar candidates. In \\S\\ref{sec:model} we briefly discuss the neutrino production mechanism in microquasars. In \\S\\ref{sec:L_jet} we use observational data available for each source to estimate the jet parameters, and then use these parameters to derive the expected neutrino flux. The number of neutrino induced muon events in km$^2$-scale neutrino telescopes is derived in \\S\\ref{sec:N_mu}. The implications of our results are briefly discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  There are large uncertainties involved in the derivation of the jet parameters for most of the sources listed in table 4. The best studied cases are perhaps GRS 1915+105, GRO J1655-40, and SS433. Nonetheless, we have demonstrated that if the jets in microquasars are protonic, and if a fraction of a few percent of the jet energy is dissipated on sufficiently small scales, then emission of TeV neutrinos with fluxes in excess of detection limit of the forthcoming, km$^2$ scale, neutrino telescopes is anticipated.  The present identification of microquasars, and the inferred distribution of their jet Lorentz factors, may be strongly influenced by selection effects. It is quite likely that the class of Galactic microquasars contains also sources with larger bulk Lorentz factors and smaller viewing angles, which should emit neutrinos with fluxes considerably larger than the extended microquasars discussed in this paper.  Such sources may be identified via their gamma-ray emission with, e.g. AGILE and GLAST, or by their neutrino emission. The gamma-rays should originate from larger scales where the pair production opacity is sufficiently reduced. Predictions for AGILE and GLAST will be discussed in a forthcoming paper \\cite{Guetta02}. There are currently about 280 known XRBs \\cite{Liu00,Liu01}, of which $\\sim 50$ are radio loud. These may also be potential targets for the planned neutrino detectors. Our results quoted in Table \\ref{tab:tab} are consistent with experimental upper limits on neutrino fluxes from point sources set by MACRO \\cite{Ambrosio01}, AMANDA \\cite{Biron} and SuperKamiokande \\cite{Okada00}."
    },
    "0207/astro-ph0207673_arXiv.txt": {
        "abstract": "The new measurement of the anomalous magnetic moment of the muon by the Brookhaven AGS experiment 821 again shows a discrepancy with the Standard Model value.  We investigate the consequences of these new data for neutralino dark matter, updating and extending our previous work [E.~A.~Baltz and P.~Gondolo, Phys.~Rev.~Lett.~{\\bf 86}, 5004 (2001)].  The measurement excludes the Standard Model value at $2.6\\sigma$ confidence.  Taking the discrepancy as a sign of supersymmetry, we find that the lightest superpartner must be relatively light and it must have a relatively high elastic scattering cross section with nucleons, which brings it almost within reach of proposed direct dark matter searches.  The SUSY signal from neutrino telescopes correlates fairly well with the elastic scattering cross section.  The rate of cosmic ray antideuterons tends to be large in the allowed models, but the constraint has little effect on the rate of gamma ray lines.  We stress that being more conservative may eliminate the discrepancy, but it does not eliminate the possibility of high astrophysical detection rates. ",
        "introduction": "In early 2001, the Brookhaven AGS experiment 821 measured the anomalous magnetic moment of the muon $a_\\mu=(g-2)/2$ with three times higher accuracy than it was previously known \\cite{DATA}.  Their result disagreed with the Standard Model prediction at greater than 2.6$\\sigma$.  However, a sign error in the calculation of the hadronic light-by-light contribution to $a_\\mu$ was discovered, reducing the discrepancy to 1.6$\\sigma$ \\cite{LBLerror}.  Recently, the same collaboration has released a result with much improved statistics \\cite{newdata}, and there is again a discrepancy at the $2.6\\sigma$ level. Supersymmetric particles can give significant corrections to $a_\\mu$ \\cite{a_mu_old,moroi,SM}, thus the Brookhaven measurement is an important constraint on supersymmetric models.  There has been a substantial literature on this topic since the announcement of the discrepancy \\cite{bg2001,otherpapers}, discussing various consequences of the older measurement.  In this paper, we update the results of \\cite{bg2001} concerning the implications of the Brookhaven data for supersymmetric cold dark matter, assuming that supersymmetry is the only relevant physics outside of the Standard Model.  There are two significant assumptions in our discussion.  The first is that the Standard Model prediction for the muon anomalous magnetic moment is somewhat disputed, primarily in the hadronic contribution.  This was clearly demonstrated in the sign error discovered in the last year.  The hadronic error is a very significant part of the error budget when comparing the Brookhaven results to the Standard Model.  In fact it has been claimed that the Standard Model errors have been significantly underestimated \\cite{nodiscrepancy}, but this claim has been refuted \\cite{newSM}.  Furthermore, there are new evaluations of the hadronic vacuum polarization from firstly $e^+e^-\\rightarrow{\\rm hadrons}$ indicating a larger discrepancy $(3.6\\sigma)$ and secondly hadronic tau lepton decays indicating a smaller discrepancy $(1.3\\sigma)$ \\cite{newdata}.  The second caveat is that supersymmetry is only one of many possible scenarios providing corrections to $a_\\mu$ at the weak scale.  Theoretical prejudice tends to favor supersymmetry, but other possibilities exist, summarized in Ref.~\\cite{SM}. ",
        "conclusions": "In this paper we have discussed the recent confirmation of a discrepancy with the Standard Model of the anomalous magnetic moment of the muon \\cite{newdata}, updating and expanding the results of Ref.~\\cite{bg2001}.  Assuming that supersymmetry is responsible for the discrepancy, we have investigated the consequences for astrophysical dark matter searches.  We have confirmed that the constraint significantly improves the prospects for direct detection experiments trying to measure the rare scatterings of galactic neutralinos. Neutrino telescopes are also helped by this result.  The prospects for the detection of gamma ray lines from neutralino annihilations at the galactic center are not much affected.  The prospects for detecting cosmic ray antideuterons as neutralino annihilation products are also significantly improved.  In all cases, if the discrepancy disappears, there remain supersymmetric models with detectable rates for all of these experiments."
    },
    "0207/astro-ph0207390_arXiv.txt": {
        "abstract": "We review the observational information on the constancy of the fine structure constant $\\alpha$. We find that small improvements on the measurement of $^{187}Re$ lifetime can provide significant progress in exploring the range of variability  suggested by QSO data. We also discuss  the effects of a time varying $\\alpha$ on stellar structure and evolution. We find that radioactive dating of ancient stars can offer a new observational window. ",
        "introduction": "The possibility that some of the ``fundamental constants'' may depend on time was first discussed by Dirac \\cite{Dirac1}. He remarked  that the huge  ratio of electric to gravitational forces between a proton and an  electron, about $10^{39}$, was of the same order of magnitude as the age of the universe  in units provided by the atomic constants, $e^2/m_e c^3$. If this coincidence is not casual, then one must have varying constants, their values changing as the age of the universe changes: \\emph{``This suggests that the above mentioned large numbers are to be regarded not as constants, but as simple functions of our present epoch, expressed in atomic units .... In this way we avoid the need of a theory to determine numbers of the order $10^{39}$.''} The approach of Dirac to what is now called the  hierarchy problem opened a rich field of investigation. The variability of fundamental constants was analysed by Gamow, Dyson and others and then it was forgotten for a while.  Interest in this topic has been revived in the context of string theories, where all the coupling constants and parameters, except the string tension, are actually derived quantities, which are determined by the vacuum expectation values of the dilaton and moduli. Since all these fields evolve on cosmological scales the time variation of the constants of nature during the evolution of the universe arises as a natural possibility, see e.g. \\cite{DP,Witten}.  On the observational side, Webb et al. \\cite{Webb} have presented evidence for a cosmological evolution of the fine structure constant $\\alpha=e^2/\\hbar c$. The absorption spectra of diffuse clouds illuminated by quasars suggest that   ten billion years ago  $\\alpha$ was slightly smaller, by about ten part per million. Of course this indication, if confirmed, would have enormous importance.  This short review attempts to provide an answer to some natural questions following the  claim of ref.\\cite{Webb}:  i)What are the observational constraints on the variability of $\\alpha$ and how do they compare with the result of ref. \\cite{Webb}?  ii)What  are the prospects for improvements?  iii)What  are the effects of a time  varying $\\alpha$ on stellar structure and evolution?      \\begin{figure}[htb] \\begin{center} \\includegraphics[width=0.5\\textwidth, angle=270]{fig_qso.ps} \\end{center} \\caption[]{ $\\Delta \\alpha/\\alpha$ vs. fractional look-back time to the Big Bang, from \\cite{Webb}. } \\label{Figqso} \\end{figure}    \\section {What do quasars tell us?}  The measurement of  the spectra of distant quasars as a mean to study possible  variations of $\\alpha$ was first suggested  by Savedoff \\cite{Savedoff}. Narrow lines in quasar spectra are produced by absorption of radiation in intervening clouds of gases. Essentially one needs to identify two (sets of) lines, which depend differently on $\\alpha$, so as to extract the value of the redshift factor $z$ together with the value of $\\alpha$ at that epoch. The fine structure doublets of ``alkali atoms''--  a term used to denote   atoms and atomic ions with just one electron in the outer shell  -- are well suited for this study.  This method has been  used by several authors and it has been  recently applied to a selection of high resolution observations, see \\cite{VPI}. No indication of a variable $\\alpha$ has been found and the constraint $\\Delta \\alpha / \\alpha = (-4.6 \\pm 4.3[\\mbox{stat}] \\pm 1.4[\\mbox{sys}] ) 10^{-5}$ has been obtained \\cite{VPI} on the possible deviation at $z=2\\div4$ from the present ($z=0$)value.   On the other hand, Webb et al. \\cite{Webb} have used a ``many multiplet'' method, where $\\alpha$ is estimated from comparison of the lines of \\emph{different} atomic species, so as to obtain a sensitivity gain. The data are summarized in Fig. \\ref{Figqso}. In this way they claim to have found  a deviation from the present $\\alpha$ value over the redshift range  $z=1\\div3$ : \\begin{equation} \\Delta \\alpha / \\alpha = (-0.72 \\pm 0.18 ) \\cdot 10^{-5} \\end{equation}  This result has been criticized in ref. \\cite{VPI} on the grounds that some systematic effect could mimic the variation of $\\alpha$. For example, the lines of the two atomic species considered in \\cite{Webb} are situated in different regions, so that calibration errors could simulate the effect of $\\alpha$ variation. In contrast, the method based on the fine splitting of a line of the same species is not affected by these uncertainty sources.    \\begin{table}[htb] \\caption{Summary on the variation of $\\alpha$ } \\begin{center} \\renewcommand{\\arraystretch}{1.4} \\setlength\\tabcolsep{9pt} \\begin{tabular}{p{2.2cm}lllll} \\hline\\noalign{\\smallskip} Source &$\\Delta \\alpha /\\alpha$ & Look back  & $z^*$  & $\\dot\\alpha/\\alpha$ & ref.\\\\ &                        & time (Gyr)           &    &     (yr$^{-1}$)&        \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Laboratory & $ \\leq 1.6 \\cdot 10^{-14}$ & $ 4 \\cdot 10^{-10}$ & 0     & $ \\leq 4 \\cdot 10^{-14}$ &\\cite{lab}\\\\ Oklo       & $ \\leq 1 \\cdot 10^{-7}$    &  1.8       &  $\\simeq 0.1 $ & $ \\leq 6 \\cdot 10^{-17}$ &\\cite{oklo3}\\\\ Meteorites & $\\leq 4  \\cdot 10^{-6}$   &   4.5       &  $\\simeq 0.4 $ & $ \\leq 2 \\cdot 10^{-15}$ & --\\\\ $^{12}$C & $\\leq  10^{-2}$   &   $\\simeq 10$       &  $\\simeq 1.5 $ & $ \\leq 10^{-12}$ &--\\\\ stellar dating      & $ \\leq 10^{-3}$      & $\\simeq 10 $& $\\simeq 1.5$        & $  \\leq 10^{-13}$    & --\\\\ QSO(doublet)  & $ \\leq 10^{-4}$      & $\\simeq 11 - 13 $&    2--4        & $  \\leq 10^{-14}$    &\\cite{VPI}\\\\ \\textbf{QSO(multiplet)}        & $\\mathbf{ + 1 \\cdot 10^{-5}}$      & $\\mathbf{\\simeq 8 - 12 }$& \\textbf{1--3}        & $\\mathbf{ + 1 \\cdot 10^{-15}  }$  &\\textbf{\\cite{Webb}}\\\\ CMB        & $ \\leq 5 \\cdot 10^{-2}$   & $\\simeq 14 $&  $\\simeq 10^3 $& $ \\leq 3 \\cdot 10^{-12}$ &\\cite{Avelino}\\\\ BBN        & $ \\leq 1 \\cdot 10^{-2}$   & $\\simeq 14 $ & $\\simeq 10^9 $& $ \\leq 7 \\cdot 10^{-13}$ &\\cite{Avelino}\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^*$ The red shift -- time connection is estimated for $H_o=68$ Km/s/Mpc, $\\Omega_M=0.3$ and $\\Omega_\\Lambda =0.7$ ($t_u\\simeq14$ Gyr). \\\\ \\label{Tabalfa} \\end{table} ",
        "conclusions": ""
    },
    "0207/astro-ph0207445_arXiv.txt": {
        "abstract": "We derive correlations between X-ray temperature, luminosity, and gas mass for a sample of 22 distant, $z>0.4$, galaxy clusters observed with Chandra. We detect evolution in all three correlations between $z>0.4$ and the present epoch. In particular, in the $\\Omega=0.3$, $\\Lambda=0.7$ cosmology, the luminosity corresponding to a fixed temperature scales approximately as $(1+z)^{1.5\\pm0.3}$; the gas mass for a fixed luminosity scales as $(1+z)^{-1.8\\pm0.4}$; and the gas mass for a fixed temperature scales as $(1+z)^{-0.5\\pm0.4}$ (all uncertainties are 90\\% confidence). We briefly discuss the implication of these results for cluster evolution models. ",
        "introduction": "Correlations between the X-ray properties of galaxy clusters are useful statistical tools which allow to study the cluster physics. The best-studied correlation for the low-redshift clusters is that of the X-ray luminosity with the temperature (e.g., Mushotzky 1984, David et al.\\ 1993, Markevitch 1998, Arnaud \\& Evrard 1999).  The mass of the intracluster gas correlates with both temperature (Mohr, Mathiesen \\& Evrard 1999, Vikhlinin, Forman \\& Jones 1999) and X-ray luminosity (Voevodkin, Vikhlinin \\& Pavlinsky 2002b). The observed tightness of these correlations indicates a similar formation history for all clusters and is consistent with the predictions of the self-similar models of the cluster formation.  However, the details of these scaling relations are different from the predictions of the self-similar theory. The best-known example is the slope of the $L-T$ correlation: it is observed that $L\\propto T^{2.7}$ for hot clusters (e.g., Markevitch 1998), while theory predicts $L\\propto T^{2}$ (e.g., Kaiser 1981).  Such deviations may be due to non-gravitational processes such as preheating (e.g., Cavaliere, Menci \\& Tozzi 1997) or radiative cooling and feedback from star formation (Voit \\& Bryan 2001). Observing the evolution of scaling relations can provide useful constraints on such models. The scaling relations at high redshift also are of great value for cosmological studies based on cluster evolution. They provide the means to convert an easily observed X-ray luminosity function into the more cosmologically useful temperature or mass functions (e.g., Borgani et al.\\ 2001).   Most of the previous high redshift studies have focused on the $L-T$ relation. Mushotzky \\& Scharf (1997) analyzed a large sample of distant clusters observed with ASCA (most at $z\\sim 0.3$, with a few at $z>0.4$) and found no evidence for evolution in the $L-T$ relation. Several results for a small number of distant cluster observed by \\emph{Chandra} have been published recently. Borgani et al.\\ (2001) analyzed a sample of 7 clusters at $z>0.5$ and concluded that the data allow at most a very mild evolution --- if the luminosity for the given temperature is $L(z)\\propto(1+z)^A$, then $A<1$. A similar conclusion has been reached by Holden et al. (2002) from an analysis of 12 clusters at $z>0.7$.  The $L-T$ relation at both low and high redshifts has a large intrinsic scatter, comparable to the expected evolutionary effects. The scatter in the low-redshift $L-T$ relation is significantly reduced when the central cooling regions of the clusters are excised from both the luminosity and temperature measurements (Fabian et al.\\ 1994, Markevitch 1998). Therefore, it is desirable to exclude the cooling cores in the distant clusters because this too may reduce the scatter and thus more easily expose any evolution. This task is feasible only with the \\emph{Chandra}'s arcsecond angular resolution.  As of Spring 2002, \\emph{Chandra} had observed 22 clusters at $z>0.4$ with sufficient exposure for accurate temperature measurements. Most of this sample is derived from flux-limited X-ray surveys: 7 clusters from the EMSS (Henry et al.\\ 1992), 11 from the \\emph{ROSAT} serendipitous surveys, 160 deg$^2$ (Vikhlinin et al.\\ 1998), RDCS (Rosati et al.\\ 1998), WARPS (Ebeling et al. 2000), and 2 from the \\emph{ROSAT} All-Sky Survey. We use these \\emph{Chandra} observations to derive correlations between the X-ray luminosity, temperature, and gas mass at $z>0.4$.  We use $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ throughout.    \\begin{table*} \\def\\j{\\phantom{1}} \\caption{Cluster sample} {\\centering \\footnotesize  \\begin{tabular}{lccrrc} \\hline \\hline Name            & $z$   & $T$  & $L_{0.5-2}^a$ & \\multicolumn{1}{c}{$L_{\\rm bol}^b$}  & $\\M324$ \\\\ &       & (keV)&  &  &  $(10^{14}\\, M_\\odot)$ \\\\  \\hline MS 0016+1609     & 0.541 &  $\\j 9.9\\pm0.5$ &  22.8 &  113.3&  $ 6.43\\pm 0.65$  \\\\ MS 0302+1658     & 0.424 &  $\\j 3.6\\pm0.5$ &   4.7 &   10.6&  $ 1.07\\pm 0.40$  \\\\ MS 0451--0305    & 0.537 &  $\\j 8.1\\pm0.8$ &  20.7 &   91.7&  $ 3.68\\pm 0.77$  \\\\ MS 1054--0321    & 0.823 &  $\\j 7.8\\pm0.6$ &  16.5 &   70.9&  $ 2.58\\pm 0.37$  \\\\ MS 1137+6625     & 0.782 &  $\\j 6.3\\pm0.4$ &   8.4 &   32.4&  $ 1.41\\pm 0.28$  \\\\ MS 1621+2640     & 0.426 &  $\\j 7.6\\pm0.9$ &   6.3 &   27.0&  $ 2.89\\pm 0.62$  \\\\ MS 2053-0449     & 0.583 &  $\\j 5.2\\pm0.7$ &   3.5 &   10.8&  $ 0.95\\pm 0.32$  \\\\ CL 1120+2326     & 0.562 &  $\\j 4.8\\pm0.5$ &   3.7 &   12.5&  $ 1.19\\pm 0.27$  \\\\ CL 1221+4918     & 0.700 &  $\\j 7.2\\pm0.6$ &   7.0 &   28.7&  $ 2.01\\pm 0.36$  \\\\ CL 1416+4446     & 0.400 &  $\\j 3.7\\pm0.3$ &   4.2 &    8.9&  $ 1.42\\pm 0.38$  \\\\ CL 1524+0957     & 0.516 &  $\\j 5.1\\pm0.6$ &   4.5 &   15.7&  $ 1.67\\pm 0.40$  \\\\ CL 1701+6421     & 0.453 &  $\\j 5.8\\pm0.5$ &   4.9 &   15.9&  $ 1.81\\pm 0.47$  \\\\ CL 0848+4456     & 0.574 &  $\\j 2.7\\pm0.3$ &  10.6 &   38.8&  $ 0.36\\pm 0.13$  \\\\ WARPS 0152--1357 & 0.833 &  $\\j 5.8\\pm0.6$ &   1.2 &    3.1&  $ 2.88\\pm 0.55$  \\\\ RDCS 0848+4452   & 1.261 &  $\\j 4.7\\pm1.0$ &   1.8 &    6.0&  $ 0.20\\pm 0.08$  \\\\ RDCS 0910+5422   & 1.100 &  $\\j 3.5\\pm0.7$ &   2.0 &    5.9&  $ 0.26\\pm 0.11$  \\\\ RDCS 1317+2911   & 0.805 &  $\\j 2.2\\pm0.5$ &   0.8 &    2.0&  $ 0.21\\pm 0.09$  \\\\ RDCS 1350.0+6007 & 0.805 &  $\\j 4.3\\pm0.6$ &   4.2 &   13.2&  $ 1.04\\pm 0.33$  \\\\ RASS 1347--114   & 0.451 &  $  14.1\\pm0.9$ &  60.1 &  260.4&  $ 8.77\\pm 1.60$  \\\\ RASS 1716+6708   & 0.813 &  $\\j 6.6\\pm0.8$ &   7.2 &   28.8&  $ 1.25\\pm 0.33$  \\\\ 3C295            & 0.460 &  $\\j 5.3\\pm0.5$ &   9.1 &   16.3&  $ 1.51\\pm 0.48$  \\\\ CL0024+17        & 0.394 &  $\\j 4.8\\pm0.6$ &   3.1 &    9.2&  $ 1.24\\pm 0.37$  \\\\ \\hline \\end{tabular} \\par  \\medskip  \\begin{minipage}{0.8\\linewidth} \\footnotesize  $^a$ --- Total X-ray luminosity in the 0.5--2 keV band within the 2~Mpc radius, $10^{44}\\,$erg~s$^{-1}$.  $^b$ --- Bolometric luminosity within the 2~Mpc radius excluding the central cooling regions, $10^{44}\\,$erg~s$^{-1}$. All quantities are computed for the $\\Omega=0.3$, $\\Lambda=0.7$ cosmology.  \\end{minipage} \\par }  \\end{table*} ",
        "conclusions": "We have used \\emph{Chandra} observations of a sample of 22 distant, $z>0.4$, clusters to show that the correlations between the cluster temperature, luminosity, and gas mass evolve significantly with respect to the low-redshift relations.  Our detection of significant evolution in the $L-T$ relation appears to contradict some other recent \\emph{Chandra} studies (Borgani et al.\\ 2001, Holden et al.\\ 2002). The difference between their and our results can be traced mostly to a more consistent comparison of the high- and low-redshift samples, such as exclusion of the cool cores and extraction of the luminosities in the 2~Mpc aperture. Note that the results of Novicki et al.\\ (2002) who self-consistently use \\emph{ASCA} data for the nearby and distant clusters, agree well with our $L-T$ evolution.  It is theoretically expected that within clusters, the baryon contribution to the total mass, $f_b$, should be close to the average value in the Universe (e.g., White et al.\\ 1993).  This notion continues to gain observational support (most recently, Allen, Schmidt \\& Fabian 2002). If $f_b$ is indeed constant, the evolutions in the $M-L$ and $M-T$ relations involving the gas mass or total mass should be identical.  The observed evolution of the cluster $M_{g}-T-L$ correlations indicates that clusters at high redshift were systematically denser than at present --- hotter and more luminous for a given mass, as expected in a theory of the hierarchical self-similar formation. However, the details of the observed evolution contradict the self-similar predictions. For example, the standard theory (e.g., Bryan \\& Norman 1998) predicts that for a given temperature, the product $H(z) M_\\Delta\\,(T,z)$ should be constant, where $M_\\Delta$ is the mass measured within the radius of the overdensity $\\Delta$ with respect to the critical density at redshift $z$, and $H(z)$ is the Hubble constant. For the realistic cluster density profiles, this implies that approximately $M_\\delta(z,T) \\propto (1+z)^{-3/2}$ in almost any cosmology, where $M_\\delta$ corresponds to the mean overdensity $\\delta$ relative to the average density at redshift $z$. However, we observe little evolution in the $M_{g,\\delta}-T$ relation for the currently favored $\\Omega=0.3$, $\\Lambda=0.7$ cosmology, seemingly at odds with the theoretical prediction. This possibly indicates the importance of non-gravitational processes for heating the intracluster gas that would change $T$ but are unlikely to modify $M_g$, just as is inferred from the scaling relations for low-redshift clusters."
    },
    "0207/astro-ph0207029_arXiv.txt": {
        "abstract": "We present an analysis of different sets of gravitational N-body simulations, all describing the dynamics of discrete particles with a small initial velocity dispersion. They encompass very different initial particle configurations, different numerical algorithms for the computation of the force, with or without the space expansion of cosmological models. Despite these differences we find in all cases that the non-linear clustering which results is essentially the same, with a well-defined simple power-law behaviour in the two-point correlations in the range from a few times the lower cut-off in the gravitational force to the scale at which fluctuations are of order one. We argue, presenting quantitative evidence, that this apparently universal behaviour can be understood by the domination of the small scale contribution to the gravitational force, coming initially from nearest neighbor particles. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207503_arXiv.txt": {
        "abstract": "The growth and saturation of magnetic field in conducting turbulent media with large magnetic Prandtl numbers are investigated. This regime is very common in low-density hot astrophysical plasmas. During the early (kinematic) stage, weak magnetic fluctuations grow exponentially and concentrate at the resistive scale, which lies far below the hydrodynamic viscous scale. The evolution becomes nonlinear when the magnetic energy is comparable to the kinetic energy of the viscous-scale eddies. A physical picture of the ensuing nonlinear evolution of the MHD dynamo is proposed. Phenomenological considerations are supplemented with a simple Fokker--Planck model of the nonlinear evolution of the magnetic-energy spectrum. It is found that, while the shift of the bulk of the magnetic energy from the subviscous scales to the velocity scales may be possible, it occurs very slowly --- at the resistive, rather than dynamical, time scale (for galaxies, this means that generation of large-scale magnetic fields cannot be explained by this mechanism). The role of Alfv\\'enic motions and the implications for the fully developed isotropic MHD turbulence are discussed.  \\vskip1cm  \\noindent 24 July 2002\\\\\\\\ \\fbox{{\\em New Journal of Physics}~{\\bf 4}~84~(2002) ({\\tt http://www.njp.org})}\\\\\\\\ E-print {\\tt astro-ph/0207503} ",
        "introduction": "\\label{sec_intro}  It has long been appreciated~\\cite{Batchelor_dynamo} that a generic three-dimensional turbulent flow of a conducting fluid is likely to be a {\\em dynamo,} i.e., it amplifies an initial weak seed magnetic field provided the magnetic Reynolds number is above a certain instability threshold (usually between~10 and~100). The amplification is exponentially fast and occurs due to stretching of the magnetic field lines by the random velocity shear associated with the turbulent eddies. The same mechanism leads to an equally rapid decrease of the field's spatial coherence scale, as the stretching and folding by the ambient flow brings the field lines with oppositely directed fields ever closer together~\\cite{Ott_review,SCMM_folding,SMCM_structure}.  If the magnetic Prandtl number ($\\Pr=\\nu/\\eta$, the ratio of the viscosity and the magnetic diffusivity of the fluid) is large and the hydrodynamic turbulence has Kolmogorov form, the scale of the magnetic field can decrease by a factor of~$\\Pr^{1/2}$ below the viscous scale of the fluid. This situation is known as {\\em the Batchelor regime}~\\cite{Batchelor_regime} and is characterized by the magnetic-energy concentration at subviscous scales. This MHD regime is very common in astrophysical plasmas, which tend to have very low density and high temperature. Examples include warm interstellar medium, some accretion discs, jets, protogalaxies, intracluster gas in galaxy clusters, early Universe,~etc. (a representative set of recent references is~\\cite{KA,Kulsrud_review,SBK_review,Balbus_Hawley_review,Heinz_Begelman,Kulsrud_etal_proto,Malyshkin_clusters,Narayan_Medvedev,Son,Christensson_Hindmarsh_Brandenburg}). The object of our principal interest are magnetic fields of galaxies, often thought to have been generated by the turbulent dynamo in the interstellar medium. The challenge is to reconcile the preponderance of small-scale magnetic fluctuations resulting from the weak-field (kinematic) stage of the dynamo with the observed galactic fields coherent at scales of approximately 1~kpc, or 10~times larger than the outer scale of the interstellar turbulence, which is forced by supernova explosions at scales of~$\\sim100$~pc (\\cite{Beck_review,Widrow_review,Han_Wielebinski} are some of the recent reviews of relevant observations).  The galactic magnetic fields generally have energies comparable to the energy of the turbulent motions of the interstellar medium, and, therefore, cannot be considered weak. Indeed, the growth of the magnetic energy established for the kinematic regime inevitably leads to the field becoming strong enough to resist further amplification and to modify the properties of the ambient turbulence by exerting a force (Lorentz tension) on the fluid. This marks the transition from the kinematic (linear) stage of the dynamo to the nonlinear regime. In the astrophysical context, the main issue is what happens to the coherence scale of the field during the nonlinear stage and, specifically, whether a mechanism could be envisioned that would effectively transfer the magnetic energy from small (subviscous) to velocity scales: perhaps as large as the outer scale of the turbulence and above.  A large amount of work has been devoted to this issue during the last 50~years (a small subset of the relevant references is~\\cite{Batchelor_dynamo,Pouquet_Frisch_Leorat,Meneguzzi_Frisch_Pouquet,Kida_Yanase_Mizushima,Chandran_closure,Cho_Vishniac,Brandenburg,Chou,MCM_dynamo}). Since most phenomenological theories of the {\\em steady-state} MHD turbulence envision a state of eventual detailed scale-by-scale equipartition between magnetic and velocity energies~\\cite{Iroshnikov,Kraichnan_IK,Goldreich_Sridhar_strong}, it has been widely expected that the dynamo would converge to such a state via some form of inverse cascade of magnetic energy. The implication would be an eventual dominant magnetic-energy concentration at the outer scale of the turbulence --- a convenient starting point for the operation of the helical mean-field dynamo~\\cite{Moffatt,Blackman_review} or of the helical inverse cascade~\\cite{Frisch_etal,Pouquet_Frisch_Leorat}. Note that we believe the inverse cascade of magnetic helicity that takes place in helical MHD turbulence~\\cite{Frisch_etal,Pouquet_Frisch_Leorat} to be unrelated to the small-scale processes we discuss here: we are interested in ways to shift the magnetic energy from subviscous scales to the outer (forcing) scale of the turbulence, while the helical inverse cascade transfers magnetic helicity from the forcing scale to even larger scales. In the galactic dynamo theory, helicity-related effects should not be relevant for the small-scale turbulence because they operate at the time scale associated with the overall rotation of the galaxy, which is 10 times larger than the turnover time of the largest turbulent eddies ($10^8$ and $10^7$~years, respectively~\\cite{KA}).  To prevent confusion, let us spell out that by {\\em turbulent dynamo} we mean the turbulent amplification of the energy of magnetic {\\em fluctuations}, not of the mean field. Fluctuation-dynamo theories describe magnetic fields at the scales of the turbulence, mean-field-dynamo theories refer to much larger scales (the mean field is usually the field averaged over the flow scales). While fast mean-field dynamo is not possible in an isotropic system without net helicity (see, e.g.,~\\cite{Moffatt}), the fluctuation dynamo is~\\cite{Batchelor_dynamo,Kazantsev,KA}. Moreover, fluctuation dynamos in three-dimensional turbulent flows are usually fast, i.e., proceed at a finite rate in the limit of arbitrarily small magnetic diffusivity.  Thus, in this work, we study the possibility of the nonhelical, nonlinear, large-$\\Pr$ fluctuation dynamo saturating with magnetic energy concentrated at the outer scale of the turbulence. Using a simple spectral Fokker--Planck model inspired by the quantitative theory of the kinematic stage of the dynamo and by a physical model of the nonlinear stage, we establish a possible mechanism for shifting the bulk of the magnetic energy from the subviscous scales to the velocity scales. The energy transfer mechanism proposed here involves selective resistive decay of small-scale fields combined with continued amplification of the field at larger scales. Therefore, the time needed for this process to complete itself turns out to be the resistive time associated with the velocity scales of the turbulence, {\\it not} a dynamical time. In large-$\\Pr$ systems, this time is enormously long and, in fact, can easily exceed the age of the Universe. Thus, a steady state with magnetic energy at velocity scales, even if attainable in principle, is not relevant for such systems. However, the scenario of the evolution and saturation of the nonlinear dynamo proposed here has fundamental implications for our understanding of the structure of the forced isotropic MHD turbulence and of the results of the recent numerical experiments.  The plan of the rest of this paper is as follows. In \\secref{sec_kinematic}, we review the kinematic theory. It is worked out in some detail because it constitutes the mathematical framework of the subsequent nonlinear extension presented in \\secref{sec_FP}. In \\secref{sec_phenom}, a phenomenological theory of the nonlinear dynamo is proposed. This section contains our main physical argument and is, therefore, the part of this paper that a reader not interested in details should peruse before jumping to the discussion section at the end. In \\secref{sec_FP} we proceed to formulate a Fokker--Planck model of the evolution of the magnetic-energy spectrum. The model is then solved numerically and analytically. Self-similar solutions are obtained. Both the qualitative physics of \\secref{sec_phenom} and the analytical solutions of \\secref{sec_FP} exhibit a resistively controlled effective transfer of the magnetic energy to the velocity scales. The resulting long-time asymptotic state of the fully developed isotropic MHD turbulence is discussed in \\ssecref{ssec_MHD}. \\Secref{sec_conclusions} contains a discussion of our findings and of the issues left open in our approach. ",
        "conclusions": "\\label{sec_conclusions}  We have presented a physical scenario of the evolution of the nonlinear large-$\\Pr$ dynamo according to which a magnetic-energy spectrum concentrated at velocity scales does emerge, but the time scale for it is resistive, not dynamical: the resistive scale approaches the viscous scale after~$\\tres\\sim\\(\\eta\\kd^2\\)^{-1}$. In the interstellar medium, the Spitzer~\\cite{Spitzer} value of the magnetic diffusivity is~$\\eta\\sim10^{7}~{\\rm cm}^2/{\\rm s}$, $\\kd\\sim10^{-16}~{\\rm cm}^{-1}$~\\cite{SBK_review}, so $\\tres\\sim 10^{17}$~years, which exceeds the age of the Universe by 7~orders of magnitude! The conclusion is that this mechanism cannot be invoked as a workable feature of the galactic dynamo --- at least not if the dissipation of the magnetic field is controlled by the Spitzer magnetic diffusivity.      The statistically steady solutions discussed in \\ssecref{ssec_MHD} represent the possible long-time asymptotic states of the fully developed isotropic large-$\\Pr$ MHD turbulence. Since the approach to these states is controlled by the resistive time scale associated with the velocity spatial scales, they have practically no relevance for astrophysical plasmas, which have large~$\\Re$ and huge~$\\Pr$. Studying these states may, therefore, appear to be a purely academic exercise in fundamental turbulence theory. While the importance of such pursuits ought never to be underestimated, we should like to point out another, more practical angle from which our results could be viewed. The enormous scale separations so often encountered in astrophysical objects are not accessible in the turbulence simulations that constitute the state of the art on Earth, so one usually has to be satisfied with only very modest $\\Re$ and~$\\Pr$. Most of the numerical studies of the small-scale dynamo have so far adopted the strategy of resolving a reasonable hydrodynamic inertial range while only allowing for Prandtl numbers of at most order~10~\\cite{Meneguzzi_Frisch_Pouquet,Kida_Yanase_Mizushima,Brandenburg,Chou,MCM_dynamo,Brummell_Cattaneo_Tobias}. In this context, our steady-state solution and formula~\\eref{W_curve} point to an important possibility. When the scale separations in the problem are not very large, the system may converge to a steady state with a subequipartition magnetic energy $W(\\infty)<\\Weq$. Subequipartion saturated states have indeed been reported in the numerical experiments cited above. Since the condition for the true asymptotic regime is $\\Pr\\gg\\Re^{1/2}$ (\\ssecref{ssec_MHD}), numerical experiments with very high resolution are required to adequately study the large-$\\Pr$ MHD turbulence.  In conclusion, we list some of the unresolved issues that require further study.  \\begin{itemize}  \\item The detailed mechanism of the flow suppression by the small-scale magnetic fields is a long-standing physical problem~\\cite{Cattaneo_Hughes_Kim,Zienicke_Politano_Pouquet,Kim,Nazarenko_Falkovich_Galtier}. In our model, we have conjectured the suppression of the shearing flows but allowed for oscillatory Alfv\\'enic motions, that do not stretch the field (\\ssecref{ssec_nlin_growth}). Strictly speaking, only the parallel component of the shear tensor~$\\bnabla\\vu$ must be suppressed in order to block the growth of the magnetic energy. The perpendicular components, if unaffected, could perhaps mix the field (without amplifying it) via a quasi-two-dimensional field-line-interchange mechanism and thus prevent any increase in the field's characteristic scale. In fact, this seems to be the physics behind the recent numerical results on the large-$\\Pr$ MHD turbulence {\\em with a fixed uniform mean field}~\\cite{Cho_Lazarian_Vishniac_new_regime}. However, it is unclear to what extent such motions can remain quasi-two-dimensional and avoid twisting up the magnetic field, which would then resist further mixing --- particularly, in our case of a folded field-line structure without an externally imposed mean field.  \\item Our physical picture of the final approach to saturation (\\ssecref{ssec_approach}) is not based on a solid phenomenological argument such as that for the nonlinear-growth stage, and remains incompletely understood. Indeed, it remains to be proven definitively that the selective decay will continue after the small-scale magnetic energy reaches the energy of the outer-scale eddies. Due to resolution constraints discussed above, no numerical evidence of this regime is available as yet.  \\item Two possibilities for the long-time behaviour of the isotropic MHD turbulence were identified in \\ssecref{ssec_MHD}: saturation in the generic Alfv\\'enic MHD turbulent state and saturation in a state where a significant amount of the magnetic energy remains tied up in the viscous-scale fields. Which one is realized depends on the way the small-scale folded fields interact with the inertial-range Alfv\\'en-wave turbulence. We should like to observe that there is no numerical evidence available at present that would confirm that, the isotropic forced MHD turbulence {\\em without externally imposed mean field} --- at any Prandtl number --- attains the Alfv\\'enic state of scale-by-scale equipartition envisioned in Iroshnikov--Kraichnan~\\cite{Iroshnikov,Kraichnan_IK} or Goldreich--Sridhar~\\cite{Goldreich_Sridhar_strong} phenomenologies. In fact, medium-resolution numerical investigations~\\cite{MCM_dynamo} rather seem to support the final states with small-scale energy concentration even for~$\\Pr=1$. This does not mean that the Alfv\\'enic picture is incorrect {\\em per se}. However, all existing phenomenologies of the Alfv\\'enic turbulence depend on Kraichnan's~\\cite{Kraichnan_IK} assumption that the most energetic magnetic fields are those at the outer scale. This is automatically satisfied if a finite mean field is imposed externally~\\cite{Maron_Goldreich,Cho_Lazarian_Vishniac_GS,Cho_Lazarian_Vishniac_new_regime}. However, it remains to be seen if such a distribution of energy is set up self-consistently when the turbulence is isotropic. The alternative identified here involves Alfv\\'enic motions superimposed on small-scale folded fields.  \\item An extension of the MHD description itself that may change the properties of the small-scale magnetic turbulence is the incorporation of the Braginskii~\\cite{Braginskii} tensor viscosity~\\cite{Montgomery,MCM_dynamo,Malyshkin_Kulsrud}. Even at magnetic energies small enough for the kinematic approximation to hold, the plasma is already well magnetized and the anisotropy of the viscous stress tensor leads to suppression of the velocity diffusion perpendicular to the local magnetic field~\\cite{Malyshkin_Kulsrud}. This anisotropy is all the more important in view of the locally anisotropic (folding) structure of the magnetic field itself~\\cite{SCMM_folding,Malyshkin_Kulsrud}.  \\item Another important extension is to allow compressible motions. The interstellar turbulence is sonic at the outer (supernova) scale but becomes subsonic and predominantly vortical in the inertial range (cf.~\\cite{Balsara_Pouquet,Lithwick_Goldreich_compr}). This is the commonly accepted justification for the use of incompressible MHD in the models of galactic dynamo. Of course, a certain amount of compressive motion is always present in the interstellar medium (as well as in other astrophysical plasmas), and it has been suggested in the literature that the compressibility of the motions and interactions between Alfv\\'en waves and density inhomogeneities can be important in the description of turbulence (see, e.g.,~\\cite{Tsiklauri_Nakariakov} and many recent references cited therein). The treatment of the kinematic-dynamo problem for a model compressible case can be found in~\\cite{SBK_review,SCMM_folding}. However, it remains to be understood whether and how weak compressibility affects the basic properties of the nonlinear dynamo.  \\end{itemize}  It is a very old observation that any extension of the domain of one's knowledge lengthens the perimeter along which this domain borders on the unknown. The perceptive reader must have realized that this work raises more questions than it gives definitive answers. We shall address these questions in our future investigations.  \\ack  The authors wish to thank B.~Chandran for several stimulating discussions. This work was supported by the UKAEA Agreement No.~QS06992, the EPSRC Grant No.~GR/R55344/01, and by the US-DOE Contract No.~DE--AC02--76CHO3073.  \\appendix"
    },
    "0207/astro-ph0207509_arXiv.txt": {
        "abstract": "Time evolution of an advection-dominated accretion flow is explored in terms of relativistic fluid equations. An axially symmetric, vertically averaged mean flow is constructed and then perturbed. An axisymmetric pattern is followed as it propagates in the form of a wave towards the horizon of a rotating (Kerr) black hole. Several assumptions are relaxed in comparison with previous works (Manmoto et al.\\ 1996). The wave reflection and steepening are examined and the influence of black hole angular momentum is discussed. ",
        "introduction": "In this paper we report on computations in which relativistic hydrodynamics has been employed to describe propagation of a wave pattern in accreted fluid. An axisymmetric perturbation is imposed on a background model of a transonic, advection-dominated flow. The fluid has negligible self-gravity; it moves in spacetime of a Kerr black hole that is described by parameters $M$ and $a$.\\footnote{Hereafter, equations will be written in geometrized units and scaled with the central mass $M$. Hence, dimensionless angular momentum acquires values $0{\\leq}a\\leq1$. Corresponding quantities in physical units can be obtained by straightforward conversions (Misner et al.\\ 1993, p.~36). Gravitational radius, $R_{\\rm{g}}=GM/c^2$, will be adopted as a convenient unit of length.} While our approach is essentially one-dimensional, it is able to treat non-linear regime of the growing perturbation. The model can provide some insight into phenomena expected to occur in astrophysically realistic models.  We build our computations on the results from several earlier papers. Particularly relevant have been computations of Manmoto et al.\\ (1996; cited as P1 hereafter) who explored wave propagation and reflection in a transonic flow (see also Kato et al.\\ 1998, chapt.~11). These authors found clear evidence for reflecting waves near the horizon. The perturbation then steepens into a shock wave. Observational motivation for these studies comes from the evidence for X-ray fluctuations, origin of which is often linked with disturbances in the innermost regions of accretion flows. In P1, the background accretion flow was constructed in the pseudo-Newtonian framework, while we employ general relativistic description and relax several other restrictions.  The main aim of this note is to show the results of our computations, where black hole rotation stands as one of the parameters. We refer to other, more extended papers for a thorough discussion of various connections and the place that these computations may have in a general scheme of the black-hole accretion. For instance, a lot of interest has been focused on the possibility of limit-cycle oscillations in accretion disks (Honma et al.\\ 1991). Indeed, non-linear time-dependent calculations of vertically integrated transonic accretion disks exhibit limit-cycle behaviour associated with thermal instability (Szuszkiewicz \\& Miller 1998). Recently, Bate et al.\\ (2002) analyzed axisymmetric waves with a two-dimensional hydrodynamical scheme.  This paper is organized as follows. In Sec.~\\ref{sec:background} the background stationary flow is constructed. In Sec.~\\ref{sec:perturbed} a perturbation is imposed on the background flow and its time evolution is examined. Finally, the properties of the temporal behaviour are summarized in Sec.~\\ref{sec:results}. ",
        "conclusions": "We examined temporal behaviour of a simple relativistic model of the advection-dominated flow. To this aim, we observed the evolution of disturbances originating in outer parts of the disk. Reflection of the waves was discussed previously in various contexts. We developed a robust code which is able to deal with the problem across the large range of disk parameters without introducing numerical viscosity. The scheme appears useful as a test bed for more complicated investigations with unstable behaviour, treated in more dimensions (work in progress).  In this note we described a stable system which after a transition period eventually relaxes to the initial steady state. Our scheme, thanks to its rather general formulation, exhibits more complex behaviour of reflected waves than the pseudo-Newtonian model described in P1. The main advantages are general relativistic treatment of hydrodynamic equations including finite propagation speed of viscous effects. We are also able to follow the evolution for longer intervals, till the shock escapes from the observed region and the model settles back to its initial equilibrium state.  VK acknowledges helpful discussion with Professor Marek Abramowicz. Support from the grants GACR 205/00/1685 and 202/02/0735 is also acknowledged.  \\appendix"
    },
    "0207/astro-ph0207023_arXiv.txt": {
        "abstract": "{ We have combined new high resolution UVES-VLT observations of a sample of four damped Ly$\\alpha$ systems (DLAs) at redshifts between $z_{\\rm abs} = 1.7-2.5$ with the existing HIRES-Keck spectra to undertake a comprehensive study of their physical conditions and their abundances of up to 15 elements. In this paper, we present abundance measurements for Mn, Ti and Mg which are among the first presented in the literature at these redshifts. We confirm the underabundance of Mn with respect to Fe as observed in lower redshift DLAs and a trend with decreasing metallicity very similar to the one observed in Galactic metal-poor stars. This agreement between the Mn/Fe ratios in DLAs and in Galactic stars suggests that the DLAs and the Milky Way share some similarities in their star formation histories. However, these similarities must be cautiously interpreted and investigated in light of all of the elements observed in the DLAs. We have obtained a first measurement and a significant upper limit of the Ti abundance at $z_{\\rm abs} \\sim 2$ from the \\ion{Ti}{ii} lines at $\\lambda_{\\rm rest} > 3000$ \\AA, and we discuss how the relative abundance of this highly depleted element can be used to prove unambiguously any enhancement of the abundance of the $\\alpha$-elements relative to Fe-peak elements. We present the abundances of Mg for two DLAs in addition to the single Mg measurement existing in the literature. Contrary to the trend expected from differential depletion, the [Mg/Si,S] ratios tend to be over-solar. The effect is at the level of the measurement errors, but worth investigating in a larger sample because it could be suggestive of a peculiar nucleosynthesis effect. ",
        "introduction": "A measure of chemical enrichment in high redshift galaxies can be obtained through the study of absorption line systems in quasars, specifically via damped Ly$\\alpha$ systems (DLAs). These systems with $N$(\\ion{H}{i}) $> 2\\times 10^{20}$ cm$^{-2}$ dominate the neutral hydrogen content of the Universe and are likely the protogalactic gas reservoirs for the formation of the majority of stars today (e.g. Wolfe et al. 1995). They provide the best opportunity to measure accurately the chemical abundances of many elements for a variety of galactic systems spanning a wide redshift interval. By comparing these abundance measurements with the chemical evolution models and abundance patterns of our Galaxy and nearby galaxies, one can infer information on the star formation history and galaxy evolution of these distant objects.  In particular, abundances of the $\\alpha$-elements (e.g. Si, O, S, Ar) relative to the Fe-peak elements (e.g. Fe, Cr, Ni) are of great importance for defining the chemical evolution history. Being produced mainly by Type II and Type Ia supernovae (SNe) with relatively different timescales -- $< 2\\times 10^7$ ($\\alpha$-elements) and $10^8-10^9$ yrs (Fe-peak elements) -- respectively, the $\\alpha$/Fe ratios are strongly dependent on the lifetimes of the element progenitor, whereas [Fe/H]\\footnote{[X/H] $\\equiv \\log$[$N$(X)/$N$(H)]$_{\\scriptsize{\\textrm{DLA}}}$ $- \\log$[$N$(X)/$N$(H)]$_{\\odot}$.} depends on the star formation rate (Tinsley 1979; Matteucci 2001). Therefore, the [$\\alpha$/Fe] versus [Fe/H] relation is a strong function of the star formation history.  \\begin{table*}[t] \\begin{center} \\caption{Metal abundances} \\label{} \\begin{tabular}{l c c c c c c c} \\hline Quasar  & $z_{\\rm abs}$  & [Fe/H] & [Si/Fe] & [Zn/Fe] & [Mn/Fe] & [Ti/Fe] & [Mg/Fe] \\smallskip \\\\ \\hline Q0100+13   & 2.309 & $-1.78\\pm 0.08$ & $+0.35\\pm 0.07$$^{\\star}$ & $+0.25\\pm 0.04$ & $-$   & $-$             & $+0.44\\pm 0.19$ \\\\ Q1331+17   & 1.776 & $-2.01\\pm 0.09$ & $+0.61\\pm 0.03$ & $+0.75\\pm 0.05$ & $-0.15\\pm 0.04$ & $< -0.45$       & $+0.83\\pm 0.22$ \\\\ Q2231$-$00 & 2.066 & $-1.20\\pm 0.09$ & $+0.40\\pm 0.05$ & $+0.41\\pm 0.07$ & $-0.16\\pm 0.10$ & $+0.70\\pm 0.09$ & $-$ \\\\ Q2343+12   & 2.431 & $-1.19\\pm 0.06$ & $+0.54\\pm 0.05$ & $+0.62\\pm 0.06$ & $-0.21\\pm 0.05$ & $-$             & $-$ \\\\ \\hline \\end{tabular} \\begin{minipage}{160mm} \\smallskip $^{\\star}$ The detected \\ion{Si}{ii} lines are all saturated, so we report the \\ion{S}{ii} column density. \\\\ Abundances relative to the solar values of Grevesse et al. (1996). \\\\ The [Fe/H] ratios correspond to the total metallicities, whereas the [X/Fe] ratios are computed by summing only the column densities of the components detected both in the \\ion{X}{ii} and \\ion{Fe}{ii} profiles. \\end{minipage} \\end{center} \\vspace{-0.4cm} \\end{table*}  Unfortunately this scheme to discriminate between different star formation histories has led to contradictory conclusions and interpretations of the observed relative abundances of the damped Ly$\\alpha$ systems (e.g. Lu et al. 1996; Centuri\\'on et al. 2000). The principle difficulty is to disentangle nucleosynthetic contributions from dust depletion effects. Because we are studying gas-phase abundances in DLAs, the measured abundances may not represent the intrinsic composition of the system if part of the elements is removed from the gas to the solid phase (Savage \\& Sembach 1996). Several pieces of evidence show that some dust is indeed present in DLAs with a dust-to-gas ratio between 2 to 25\\% of the Galactic value (Pei et al. 1991; Vladilo 1998). The presence of dust hence implies a depletion of refractory elements (e.g. Si, Fe, Cr, Ni) preferentially incorporated into dust grains.  In such a mixed situation (nucleosynthesis plus dust), the analysis of DLA chemical histories is severely limited by the low number of routinely observed elements: {\\em $\\alpha$-element} Si and {\\em Fe-peak element} Fe. Occasionally the very important non-refractory {\\em $\\alpha$-elements} O, S and Ar (located in the Ly$\\alpha$ forest) and more frequently the non-refractory {\\em Fe-peak element} Zn are detected. We show in this paper that Mn, Ti and Mg -- Fe-peak and $\\alpha$-elements respectively -- are additional elements which can provide important clues to the nature of the DLA star formation history. Well studied in Galactic stars (Nissen et al. 2000; Prochaska \\& McWilliam 2000; Chen et al. 2000), relatively little attention has been granted to these elements in DLAs due to their difficult detections (for Mn see Lu et al. 1996; Pettini et al. 2000; Ledoux et al. 2002 and for Ti see Prochaska \\& Wolfe 1999; 2002; Ledoux et al. 2002). Ratios of Ti and Mn relative to Fe are of special importance because differential depletion and nucleosynthesis tend to work in the opposite sense. Therefore, these ratios help alleviate the dust/nucleosynthesis degeneracy inherent to many other ratios (Prochaska \\& Wolfe 2002). Meanwhile, Mg is a refractory $\\alpha$-element and a comparison of its abundance with Si, S and O provides greater insight into the dust depletion pattern inherent to DLAs.  We report here three new Mn measurements in DLAs at $z_{\\rm abs} \\sim 2$ in addition to the two previous measurements of Lu et al. (1996) and Ledoux et al. (2002) at similar redshifts, two new Mg measurements representing the only Mg measurements in addition to the single previous measurement of Srianand et al. (2000), and the first Ti measurements from \\ion{Ti}{ii} lines at $\\lambda_{\\rm rest} > 3000$ \\AA\\ in DLAs at $z_{\\rm abs} \\sim 2$.  In Sect.~2 we briefly review our observations. In Sect.~3 we describe how we have succeeded to measure the column densities of the weak lines we are interested in and we show the results. In Sect.~4, we highlight the importance of our measurements by discussing individually element by element. ",
        "conclusions": "In the following discussion we interpret our results on the Mn, Ti and Mg abundances and underline their complementarity to the measurements of other elements.  {\\em Manganese}. Major recent studies of Mn abundances in the Milky Way have been completed by Nissen et al. (2000) and Prochaska \\& McWilliam (2000). About 120 [Mn/Fe] measurements from F and G stars in the metallicity interval $-1.4<$ [Fe/H] $<0.1$ have been obtained. Their results confirmed that Mn behaves in an opposite sense to the $\\alpha$-elements: a steady decline to [Mn/Fe] $\\sim -0.15$ from solar metallicity to [Fe/H] $\\sim -0.7$ and a possible drop in [Mn/Fe] below [Fe/H] $\\sim -0.7$. Mn is an Fe-peak element produced by massive stars and mostly by SNIa. Its behavior is not completely understood, although it is relatively well accounted for by metallicity-dependent yields for massive stars and for SNIa, illustrating a nice example of the ``odd-even'' effect for the Fe-nuclei (Goswami \\& Prantzos 2000).  Although Mn is a refractory element, the Mn/Fe ratio is an important and useful diagnostic ratio. Indeed, from Savage \\& Sembach (1996) we can see that Mn and Fe have nearly identical depletion levels in the warm halo ISM clouds (the physical environment whose depletion level is most likely similar to the DLAs) and Mn is less depleted than Fe in the cold ISM clouds. Therefore, as stressed by Lu et al. (1996), sub-solar [Mn/Fe] ratios in gas-phase abundances must be interpreted in terms of nucleosynthesis while super-solar [Mn/Fe] ratios would require significant dust depletion.  \\begin{figure}[t] \\centering \\includegraphics[width=9cm]{H3718F6a.ps} \\includegraphics[width=9cm]{H3718F6b.ps} \\caption{Relative abundances of [Mn/Fe] versus $z_{\\rm abs}$ of the DLA systems (upper panel) and versus the abundances of Zn, i.e. metallicities (lower panel). The open circles are literature values (Lu et al. 1996; Lopez et al. 1999; Pettini et al. 1999, 2000; Ledoux et al. 2002) which have been corrected from a possible underestimation due to a larger number of detected \\ion{Fe}{ii} components relative to the \\ion{Mn}{ii} components (see Section 3) when a detailed profile modeling was available. The filled circles are our measurements presented here. We have included to this sample only the DLAs for which a direct \\ion{H}{i} column density measurement is known and we assume that in DLAs the metallicity [Zn/H] $\\approx$ [Fe/H] in absence of dust. The dots are [Mn/Fe] versus [Fe/H] values measured in Galactic stars from Prochaska \\& McWilliam (2000).} \\label{} \\end{figure}  The three new [Mn/Fe] measurements in the DLAs at $z_{\\rm abs} \\sim 2$ (see Table~1) are in agreement with the [Mn/Fe] measurements in DLAs at $z_{\\rm abs} < 2$ (Lu et al. 1996; Lopez et al. 1999; Pettini et al. 1999, 2000; Ledoux et al. 2002), and confirm that Mn is underabundant in galaxies associated with DLAs in a very similar way to Galactic stars (see Fig.~6). We also note that the DLA system toward Q1331+17 exhibits a much larger Mn/Fe ratio -- [Mn/Fe] $= -0.15\\pm 0.04$ at a metallicity of [Zn/H] = $-1.27\\pm 0.09$ ([Fe/H] $= -2.01\\pm 0.09$) -- than the observed values in the Galactic stars at the same metallicity consistent with the fact that this DLA has a large depletion level.  The nice agreement between the Mn/Fe ratios in DLAs and in Galactic stars seen in Fig.~6 with a very comparable trend with decreasing metallicity, but a greater scatter in DLAs, suggests that the DLAs and the Milky Way share some similarities in their star formation histories. However, although the greater scatter may be explained by measurement uncertainties and the effects of differential depletion, these similarities must be cautiously interpreted and investigated in light of all of the elements observed in the DLAs. For example, it remains unclear what fraction of the damped systems exhibit $\\alpha$-enhancements relative to Fe-peak elements consistent with the observations in Galactic metal-poor stars (e.g. Vladilo 1998; Centuri\\'on et al. 2000; Prochaska \\& Wolfe 2002). These similarities/discrepancies between the abundance patterns in DLAs and in Galactic stars have to be further carefully analyzed. They emphasize the importance of studying the Mn/Fe ratios in DLAs in addition to the $\\alpha$/Fe ratios to better define the differences in the star formation histories of the DLAs and our Galaxy. Moreover, since direct observations of DLAs show a wide range of galactic morphological types (e.g. Le Brun et al. 1987), it is important to examine variations within their abundance patterns to search for variations in their star formation histories. Finally, by obtaining Mn measurements in low metallicity DLAs one can test nucleosynthetic models of Mn production -- major source of Mn: SNIa or SNII with strong metallicity dependence of the yield.  {\\em Titanium}. Many Ti stellar abundance measurements exist in our Galaxy over a wide metallicity interval $-4.0<$ [Fe/H] $<0.5$ (e.g. Chen et al. 2000). Ti is generally accepted as an $\\alpha$-element, because it exhibits abundance trends similar to other $\\alpha$-elements, although this behavior has not been reproduced by chemical evolution models (Timmes et al. 1995; Goswami \\& Prantzos 2000).  Thus far, only a few Ti measurements exist in DLAs at $z_{\\rm abs} < 1.5$ (Ledoux et al. 2002), and efforts to measure Ti in DLAs at $z_{\\rm abs} > 1.5$ have focused on the pair of very weak lines at $\\lambda_{\\rm rest} \\approx 1910$ \\AA\\ and have yielded only a few tenuous detections (e.g. Prochaska \\& Wolfe 1997, 1999; Prochaska et al. 2001). Our observations of the stronger \\ion{Ti}{ii} $\\lambda$3073,3242,3384 lines resulted in one robust detection and one significant upper limit.  Ti is a refractory element and exhibits an equivalent or higher depletion level in the ISM clouds than Fe (Howk et al. 1999). This runs contrary to the expectation for an $\\alpha$-enhancement and, therefore, positive departures of [Ti/Fe] from the solar ratio is evidence for an $\\alpha$-enhancement in the system independently of the presence of dust. Negative [Ti/Fe] ratios on the contrary should provide evidence for dust depletion. We found in the DLA toward Q2231$-$00, [Ti/Fe] $= +0.70\\pm 0.09$ which suggests a dominance of Type II SNe relatively well consistent with the ratio of [S/Zn] $= 0.17\\pm 0.09$ (Dessauges-Zavadsky et al. in prep.). In the DLA toward Q1331+17, on the other hand, [Ti/Fe] $< -0.45$ suggests significant dust depletion in agreement with [Zn/Fe] $= +0.75\\pm 0.05$.  {\\em Magnesium}. Mg is an $\\alpha$-element with a refractory nature whose depletion level is slightly higher than Si but lower than Fe (Savage \\& Sembach 1996). Mg/Fe ratios are useful complements to other $\\alpha$/Fe ratios measured in the same DLA to deduce information on the star formation history of the galaxy associated with the DLA. Moreover, the study of $\\alpha$-element ratios with different depletion levels, like Si/Mg, Mg/Ti and Si/Ti, is another very interesting way to highlight the presence of dust in DLAs. But perhaps the most valuable aspect of the Mg measurements relates to the use of Mg equivalent widths to identify low redshift absorption galaxies when the Ly$\\alpha$ lines are not detected (e.g. Bergeron \\& Boiss\\'e 1991; Rao \\& Turnshek 2000). It is therefore important to know the range of Mg column densities in DLAs to confirm the empirical method to recognize them at low redshift.  While interesting for the reasons quoted above, the number of Mg measurements in DLAs remains tiny. With our two measurements a total of only three Mg abundance measurements are known in DLAs. They all show positive [Mg/S,Si] ratios of $0.10-0.20$~dex (see Table~1 and the DLA system in Srianand et al. 2000), in contradiction to the negative values expected from the differential depletion. But, since the detection of positive Mg/S,Si ratios is at the level of the measurement errors of $\\sim 0.2$~dex and differs from the maximum over-depletion of Mg relatively to Si (of 0.2~dex) observed in the ISM warm gas by only 2$\\sigma$ (Welty et al. 1999), it has to be further confirmed by additional Mg measurements before interpreting it as a peculiar nucleosynthetic effect similar to the observations of some Galactic stars (e.g. McWilliam et al. 1995)."
    },
    "0207/astro-ph0207215_arXiv.txt": {
        "abstract": "Shocks are a ubiquitous consequence of cosmic structure formation, and they play an essential role in heating galaxy cluster media. Virtually all of the gas in clusters has been processed by one or more shocks of at least moderate strength. These are collisionless shocks, so likely sites for diffusive shock acceleration of high energy particles. We have carried out numerical simulations of cosmic structure formation that directly include acceleration and transport of nonthermal protons, as well as primary and secondary electrons. Nonthermal emissions have also been computed from the resulting particle spatial and energy distributions. Here we outline some of our current findings, showing that nonthermal protons may contribute a significant pressure in cluster media, and that expected radio, X-ray and $\\gamma$-ray emissions from these populations should be important cluster diagnostics. ",
        "introduction": "At least two lines of reasoning lead us to examine the properties of energetic particle acceleration at structure formation shocks. First, virtually all the gas in filaments, groups and clusters\\footnote[1]{To simplify the discussion below we do not distinguish groups from clusters.} has at some time passed through one, or more probably, several structure shocks (e.g., Quilis et al 1998; Cen \\& Ostriker 1999; Miniati et al. 2000). Structure shocks involve very diffuse plasmas, so will be ``collisionless''.  Thus, they are likely to accelerate high energy particles (call them ``cosmic rays'' or CRs) through the so-called ``diffusive shock acceleration'', provided there is a weak magnetic field present (e.g., Blandford \\& Eichler 1987). As discussed by Blasi at this meeting and outlined briefly below, clusters should be good reservoirs of CR protons, so that over time they will accumulate, possibly even to a level where they can contribute significantly to the cluster pressure (e.g., Berezinsky et al 1997). CRs have a softer equation of state than nonrelativistic thermal plasma, but CR protons are effectively immune to nonadiabatic cooling. Consequently, they can change the thermodynamic properties of the ICM. In that event it becomes important to include them in consideration of cluster dynamics, especially in cooling flows (see, e.g., the contribution by Ryu et al in these proceedings).  The second rationale for understanding particle acceleration at structure shocks comes from diffuse nonthermal emissions seen in at least some clusters. The most compelling such evidence, known for some time, is the existence of diffuse cluster radio halos and so-called ``radio relic'' sources (e.g., Feretti \\& Giovannini 1996; Feretti 1999). The detailed properties of these two classes of radio source are different in some respects, such as polarization and location inside the clusters, but both involve substantial volumes in their host clusters. The radio halos tend to be centered on the cluster cores and are unpolarized, while the relics are most likely to be found on the perimeters of clusters and can be highly polarized. Both types of radio sources result from synchrotron radiation by substantial populations of $\\gsim$GeV electrons.  As discussed at this meeting by Fusco-Femiano, X-ray observations now show convincing evidence for diffuse, nonthermal hard X-ray emission in at least Coma and A2256 (e.g., Rephaeli et al 1999; Fusco-Femiano et al 2000). Again, this implies nonthermal electron populations. Here, however, as discussed by several other speakers at this workshop, the origin of that emission and the energy of the electrons are less certain. If the emission is nonthermal bremsstrahlung the electrons are only a little more energetic than the thermal electrons responsible for the soft X-ray continuum. On the other hand, one of the prime candidates for the hard X-ray excesses is inverse-Compton scattered CMB photons, again involving roughly GeV electrons. So far, the evidence requires only nonthermal electrons, since $\\pi^0$ decays coming from inelastic p-p collisions have not yet been detected.  However, for particle acceleration in normal plasmas we should expect energetic hadrons as well. If the accelerators behave similarly to galactic CR accelerators the energy carried by hadronic CRs is likely to be one to two orders of magnitude greater than for electrons, in fact.  To facilitate what follows it may be useful here to review very quickly the key issue of CR longevity in clusters, since that largely controls the needs for extended accelerators of CRs. First, it has long been noted that electrons responsible for the observed radio halos have such short lifetimes to radiative losses that they cannot possibly fill a cluster from a single point source (e.g., Jaffe 1977). Supposing for example, that the cluster magnetic field is $1 \\mu$G, then electrons radiating at 1 GHz have Lorentz factors, $\\gamma \\sim 2 \\times 10^4$. For Lorentz factors above a few hundred the dominant energy losses will come from inverse-Compton scattering in this case (e.g., Sarazin 1999), leading to lifetimes $t_r \\sim 4\\times 10^{12} \\gamma^{-1}$ yrs, or about $2\\times 10^8$ yrs in this case. It is simple to show that if we fix the observed radio frequency at 1 GHz, this is about the maximum lifetime of the relevant radiating electrons against combined inverse-Compton and synchrotron radiation.  Although in free flight relativistic electrons could cross a cluster in a few million years, diffuse radio emissions from the clusters and Faraday rotation through some clusters reveal the clear presence of weak magnetic fields (see the contributions by Feretti,  Clarke, and Kronberg at this meeting, for example). The observations point to a tangled, perhaps turbulent field. MHD fluctuations will severely restrict their propagation. As an example consider Bohm diffusion, corresponding to scattering on saturated field fluctuations, so that a particle mean free path approximates the particle gyro radius; that is, $D_B = \\frac{1}{3} c r_g$, for relativistic particles. This would allow GeV particles to diffuse less than a kiloparsec in a Hubble time. Bohm diffusion is a limiting case, and advective motions will surely carry CR electrons farther than that by a large factor even in the much smaller radiative lifetime of the electrons. But, the very small range of these electrons remains valid under any reasonable set of circumstances. The main consequence of this result is that CR electrons responsible for observed diffuse cluster radio and probably X-ray and predicted $\\gamma$-ray emissions must be continuously replenished somehow. The observed hard X-ray luminosity of Coma exceeds $10^{43}$erg s$^{-1}$, placing a lower bound on the replenishment rate in that cluster.   Note as well that hadronic CRs, and protons, in particular, are confined similarly by MHD wave scattering. In that case, on the other hand, radiative losses are negligible up to extremely high energies, as are losses due to inelastic collisions with the CMB and the cluster thermal plasma (e.g., Berezinsky et al 1997). So, the key consequence is that CR protons below about $10^{15}$eV = 1 PeV will be essentially locked to the ICM forever, once they are introduced.  The importance of shock heating to the ICM makes diffusive shock acceleration (DSA) an immediate candidate for production of cluster CRs, since DSA can transfer some tens of percents of the energy dissipated at the shock to the CR population (e.g., Blandford \\& Eichler 1987). For a typical cluster shock that would correspond to a CR  energy input rate $\\sim 0.1 \\rho u^3 R^2 \\sim 10^{45}$ erg/sec (e.g., Jones et al 2001). Over a Hubble time this could amount to as much as $10^{63}$ergs of CR energy. We emphasize from the start that our aim is not to argue for structure shocks as the only source of diffuse CRs in clusters, but that it is one very likely to be there and to be important. Others have estimated that radio galaxies can contribute a similar CR energy flux (e.g., En$\\ss$lin et al 1997). Intense starburst activity during galaxy formation has also been proposed (e.g., V\\\"olk et al 1996). Turbulent acceleration in the ICM may also play a significant role, especially as a means to reaccelerate CRs introduced by some other means (e.g., Jaffe 1977; Eilek \\& Wetherall 1999; Brunetti et al 2001).  In the remainder of this paper we will outline the role of structure shocks in determining the conditions in the ICM, review briefly the relevant properties of shock acceleration physics and discuss the results of some numerical simulations of structure formation that include treatment of CR acceleration via DSA, followed by advective transport and relevant energy loss mechanisms. Finally, since the ultimate test of such calculations is their match to observed cluster properties, we have computed ``synthetic observations'' of nonthermal emissions from simulated clusters, so describe some of those results here, as well. More extensive discussions of most of these issues are contained in our cited works, as well. ",
        "conclusions": "Shocks are a ubiquitous consequence of cosmic structure formation, and they play an essential role in heating of cluster media . Virtually all of the gas in cluster media has been processed by at one or more shocks of at least moderate strength. Since these shocks involve highly tenuous ionized media, they are collisionless in nature, so will not fully thermalize the plasmas passing through them. One likely consequence is the acceleration of relativistic particles, or cosmic rays. We have begun to explore through numerical simulations the roles that particle acceleration in structure shocks may play. Our current conclusions are:  $\\bullet$The shocks that are primarily responsible for heating cluster ICMs can be efficient particle accelerators, possibly generating nonthermal proton pressures on the order of 10\\% or more of the total virial pressure in cluster cores.  $\\bullet$Cluster ICMs are very good reservoirs for energetic protons, so the cosmic ray populations there reflect the full history of the cluster medium more than any single event, such as a recent merger.  $\\bullet$Relativistic electrons at most energies have loss lifetimes so short that they must either be accelerated locally or be secondary products from energetic hadronic cosmic rays in order to have populations great enough to account for detectable nonthermal emissions.  $\\bullet$There are two main regions for production of nonthermal radiation in clusters; the X-ray bright core and the outskirts where strong shocks are most likely.  $\\bullet$Inverse-Compton emission and $\\pi^0$ decays dominate the production of $\\gamma$-rays in typical clusters.  Inverse-Compton from primary electrons dominates in the outskirts, provided electrons are injected in proportion to ions comparably to the galactic cosmic rays; that is, so that a fraction of a percent of the energy flux through shocks is transferred to electrons.  $\\bullet$In cluster cores $\\gamma$-ray emission above a few hundred MeV should be dominated by $\\pi^0$ decays, whereas inverse-Compton emission from secondary electrons dominates in these regions at lower $\\gamma$-ray energies.  $\\bullet$Primary and secondary electrons may also contribute substantially to nonthermal radio synchrotron emissions in clusters. Primary electron emissions are confined to volumes close to contemporary shocks, so should be seen mostly in cluster outskirts, contributing to radio relic sources. Secondary electronic emissions should again be concentrated in cluster cores, contributing to radio halos."
    },
    "0207/gr-qc0207122_arXiv.txt": {
        "abstract": "We will expose a preliminary study on the feasibility of an experiment leading to a direct measurement of the gravitomagnetic field generated by the rotational motion of the Earth. This measurement would be achieved by means of an appropriate coupling of a TELEscope and a Foucault PENdulum in a laboratory on ground, preferably at the SOUTH pole. An experiment of this kind was firstly proposed by Braginski, Polnarev and Thorne, 18 years ago, but it was never re-analyzed. ",
        "introduction": "The search for measurable effects of a  gravitational field due to the angular momentum of the source, within the framework of General Relativity (GR), continues. In the weak and slow motion approximation of GR, the gravitomagnetic part of the gravitational potential gives rise to the Lense-Thirring effect \\cite{len}. The actual detection of this effect is entrusted both to Earth  satellites experiments and to Earth based laboratory experiments. So far, the only positive indirect result concerns an experiment of the first kind, the precession of the nodes of the orbit of the LAGEOS satellite \\cite{ciufolini}.  On the other hand, in the next years the space mission Gravity Probe B (GPB) is planned to fly, carrying gyroscopes which should verify the Lense-Thirring precession effect directly \\cite{GPB1},\\cite{GPB2}. Moreover, different possibilities connected both with the clock effect and the gravitational Sagnac effect have been considered \\cite{mashhoon},\\cite{tartaglia}.  Recently, after the completion of this work, a Earth based laboratory experiment to test directly the quadratic terms in the angular momentum of a gravitational potential, has been proposed by Tartaglia (see \\cite{Tarta1}). This proposal deserves further study. However, in what follows, I will remind a different Earth based laboratory experiment to test directly the Lense-Thirring effect, which was firstly proposed by Braginski, Polnarev and Thorne, 18 years ago \\cite{Bra}, but never reconsidered or re-analyzed. ",
        "conclusions": ""
    },
    "0207/astro-ph0207486_arXiv.txt": {
        "abstract": "{We present results of the spectroscopical follow-up observations of QSO candidates from a combined variability and proper motion (VPM) survey in a $\\sim 10$ square degrees region centered on the globular cluster M\\,3. The search is based on a large number of digitised Schmidt plates with a time-baseline of three decades. This paper reviews the candidate selection, the follow-up spectroscopy, and general properties of the resulting QSO sample. In total, 175 QSOs and Sey1s were identified among the objects from the VPM survey, with 114 QSOs and 10 Sey1s up to the pre-estimated 90\\% completeness limit of the survey at $B_{\\rm lim} \\approx 19.7$. The redshift range of the QSOs is $0.4 <z<3$. Among the 80 QSO candidates of highest priority we confirm 75 QSOs/Sey1s and 2 NELGs. We present magnitudes, colours, redshifts, and variability indices for all 181 identified QSOs/Sey1s/NELGs and spectra for the 77 QSOs/Sey1s/NELGs from our spectroscopic follow-up observations. The VPM survey uses selection criteria that are not directly relying on the spectral energy distribution of QSOs. It is therefore remarkable that the properties of the VPM QSOs do not significantly differ from those of samples from colour selection or slitless spectroscopy. In particular, we do not detect a substantial number of unusually red QSOs. The total surface density of the brighter QSOs ($17 \\le B \\le 18.5$) in our search field is found to be by a factor of $\\sim 1.8$ larger than that derived from previous optical surveys. ",
        "introduction": "The variability and proper motion (VPM) survey is a QSO search project that is based on optical long-term variability and non-detectable proper motions. Variability of flux densities and stationarity of positions are two fundamental properties of quasars, and therefore well suited as selection criteria of a QSO search (e.\\,g., Kron \\& Chiu \\cite{Kro81}; Hawkins \\cite{Haw83}; Majewski et al. \\cite{Maj91}; V\\'eron \\& Hawkins \\cite{Ver95}; Bershady et al. \\cite{Ber98}). However, due to the special demands on the number and the time-baseline of the available observations such attempts must be limited to comparatively small and confined areas. We performed the VPM technique  in two Schmidt fields of $\\sim 10$ \\,square degrees each on the basis of a large number of altogether more than 200 digitised Tautenburg Schmidt plates in the B band with a time-baseline of three decades (Meusinger et al. \\cite{Meu97}, \\cite{Meu02a}).  It is not the primary aim of this project to increase the number of known QSOs by an insignificant fraction; the problems of detecting substantial numbers of QSOs have long been overcome. Over the last decade, among others, the Durham/AAT survey (Boyle et al. \\cite{Boy90}), the Large Bright Quasar Survey (Hewett et al. \\cite{Hew95}), the Edinburgh Quasar Survey (Goldschmidt \\& Miller \\cite{Gol98}), and the Hamburg/ESO survey (Wisotzki et al. \\cite{Wis00}) have been completed. Presently, the 2dF Quasar Survey (Croom et al. \\cite{Cro01}) and the Sloan Digital Sky Survey (Schneider et al. \\cite{Sch02}) are extremely efficient at identifying very large numbers of quasars. The INT Wide Angle Survey (Sharp et al. \\cite{Sha01}) is expected to detect a statistically significant sample of high-redshift quasars. Very deep quasar samples were obtained in the Lockman hole via the X-ray satellite ROSAT (Hasinger et al. \\cite{Has98}) and in the optical domain with the Hubble Space Telescope (e.\\,g., Conti et al. \\cite{Con98}), respectively. Further, the VLA FIRST Bright Quasar Survey (e.\\,g., White et al. \\cite{Whi00}) will define a radio-selected QSO sample that is competitive in size with current optically selected samples.  Most of the criteria for the selection of QSO candidates rely upon differences in the broad-band spectral energy distribution of QSOs and stars. Despite the large number of QSOs now catalogued, the selection effects of the conventional surveys are not yet fully understood. It is therefore important to perform QSO surveys that are based on different selection methods. The VPM survey  does not directly invoke the spectral energy distribution as the primary selection criterion and provides therefore an interesting opportunity to evaluate the selection effects of more conventional optical QSO searches. For instance, a serious question concerns the possible existence of a substantial population of red QSOs. Extinction-reddened QSOs are suggested both from the AGN unification model (e.\\,g., Antonucci \\cite{Ant93}; Maiolino \\cite{Mai01}) and from the hypothesis of ultra-luminous IR galaxies (ULIRGs) as QSOs in the making (Sanders et al. \\cite{San88}). The vast majority of catalogued QSOs have uniform spectral energy distributions with a blue continuum and broad absorption lines. Over the last few years, QSOs with extreme red colours have been detected on the basis of their X-ray emission (e.\\,g., Risaliti et al. \\cite{Ris01}) or by radio surveys (Webster et al. \\cite{Web95}; Francis et al. \\cite{Fra00}; White et al. \\cite{Whi00}; Menou et al. \\cite{Men01}; Gregg et al. \\cite{Gre02}). The fraction of unusually red objects among the whole QSO population is however unknown. When compared to other optical surveys, the VPM technique has the advantage that it can discover such red QSOs as long as (1.) they are not too faint in the B band and (2.) they are not much less variable than the conventional QSOs.  The VPM survey was started in the high-galactic latitude field around M\\,3 (Meusinger et al. \\cite{Meu95}). Half of this field is covered by the CFHT blue grens survey (e.\\,g., Crampton et al. \\cite{Cra90}). The CFHT QSOs could serve as a training set and were used to define the selection thresholds for the VPM survey in such a way that a $\\sim90$\\% completeness is expected up to $B \\approx 19.7$. The strategy, the observational material,  and the data reduction for the M\\,3 field were presented in detail in Paper~1 (Scholz et al. \\cite{Sch97}). The procedures and results for the second search field, around M\\,92, are described in a series of papers (Brunzendorf \\& Meusinger \\cite{Bru01}, \\cite{Bru02}; Meusinger \\& Brunzendorf \\cite{Meu01}, \\cite{Meu02b}). A brief review  of the whole VPM project is given by Meusinger et al. (\\cite{Meu02a}). The present paper presents the QSO sample in the M\\,3 field. In Section~2, we briefly discuss the candidate selection. The spectroscopic follow-up observations are described in Section~3. Section~4 gives an overview of the properties of the QSO sample. Conclusions are given in Section~5. As in the previous papers, we adopt $H_0 = 50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $q_0=0$. ",
        "conclusions": "Conclusions}   We performed a VPM QSO survey with a limiting magnitude of $B_{\\rm lim} \\approx 20$ in a 10 square degrees field at high galactic latitude. The VPM technique proved to be an efficient method for finding QSOs. As the result of the spectroscopic follow-up observations of 198 candidates and the identification of further candidates in the NED, a sample of 175 QSOs/Sey1s with $0.4<z<3$ is available. With a stellar contamination of only about 4\\%, the high-priority QSO subsample from the VPM search is very clean. For the combined sample of high-and-medium-priority objects the fraction of established QSOs/Sey1s among all candidates is still as high as 63\\%. The completeness of the VPM QSO sample with $B \\le 19.7$ is estimated to be 94\\%. The number-magnitude relation for that sample is in good agreement with the one expected from the relation derived by Wisotzki (\\cite{Wis98}) from various QSO samples. At brighter magnitudes ($B<18.5$), we find a somewhat higher QSO surface density.   The optical broad-band colours and the spectra of the VPM-selected QSOs are not significantly different from those of QSOs selected by other optical surveys, in agreement with what we found in the M\\,92 field (Meusinger \\& Brunzendorf \\cite{Meu01}; Brunzendorf \\& Meusinger \\cite{Bru02}). Such a result can not be a priori expected since the selection criteria of the VPM survey are completely different from those in most other optical surveys. Although there is a large fraction of objects with red colours among the VPM QSO candidates, all candidates with extremely red colours were proved to be stellar contaminants. We estimate that the fraction of QSOs with unusualy red optical colours is at most a few per cent up to the limit of the survey, provided that their variability properties are not significantly different from those of the other QSOs. Some BAL QSOs are known to be considerably redder than the targets of most QSO surveys (e.\\,g., Weymann et al. \\cite{Wey91}; Menou et al. \\cite{Men01}; Hall et al. \\cite{Hal02}). The fraction of such unusual QSOs in the (incomplete) SDSS Early Data Release is less than 1\\% (Hall et al. \\cite{Hal02}), in good agreement with our result. In this context it is notable that all VPM QSOs with indication for substantial absorption are strongly variabel ($I_{\\rm var} > 1.5$). A VPM search is thus expected to be essentially unbiased against strongly absorbed QSOs, apart from the bias introduced by the band-pass of the search. The general agreement of the properties of the  VPM QSO sample with those from more conventional optical surveys suggests that the latter  do obviously not ignore a substantial number of red QSOs up to $B \\approx 20$. On the other hand, we can conclude that the VPM survey can be combined with colour search criteria in order to reach a very high efficiency without a significant loss of  completeness. Of course, we can not exclude the existence of substantial numbers of obscured red QSOs that are fainter than the current survey limit. Such objects can be found by a deeper VPM survey.       \\begin{table*}[ht] {\\renewcommand{\\baselinestretch}{0.95}\\footnotesize \\begin{tabular}{rlllrrrrrrrrr} \\toprule No. & $\\alpha$\\,(J2000) & $\\delta$\\,(J2000) & \\multicolumn{1}{c}{$z$} & run & $M_B$ & type & \\multicolumn{1}{c}{$B$} & $U-B$ & $B-V$ & $I_{\\rm pm}$ & $I_{\\rm var}$ & $I_{\\rm ltvar}$\\\\ \\midrule 1 &  13 35 14.55 & 29 44 55.3 &  1.853 &  4  & -26.39 & QSO & 19.49 & -0.54 & -0.50  & 2.33 & 1.62  &  3.57\\\\%  2291 2 &  13 35 42.48 & 28 53 30.5 &  0.647 &  2  & -24.73 & QSO & 18.40 & -0.64 &  0.09  & 1.38 & 3.19  &  6.36\\\\%  9849 3 &  13 35 50.27 & 28 48 08.7 &  1.071 &  1  & -26.94 & QSO & 17.49 & -0.68 &  0.43  & 3.32 & 3.34  &  8.13\\\\% 10601 4 &  13 36 09.57 & 29 09 16.3 &  0.560 &  5  & -23.23 & QSO & 19.55 & -0.89 &  0.36  & 1.93 & 1.69  &  1.82\\\\%  7678 5 &  13 36 10.89 & 27 10 52.8 &  1.718 &  2  & -26.95 & QSO & 18.69 & -0.62 &  0.31  & 1.67 & 3.05  &  6.14\\\\% 32600 6 &  13 36 15.06 & 28 59 11.9 &  1.419 &  4  & -25.78 & QSO & 19.37 & -0.71 &  0.54  & 1.27 & 1.95  &  3.40\\\\%  9076 7 &  13 37 51.14 & 27 14 21.1 &  1.70: &  5  & -26.09 & QSO & 19.52 & -0.26 &  1.06  & 0.22 & 2.51  &  2.47\\\\% 32136 8 &  13 37 56.25 & 29 38 39.6 &  1.216 &  3  & -25.45 & QSO & 19.28 & -0.83 &  0.43  & 1.32 & 2.77  &  2.12\\\\%  3344 9 &  13 37 59.79 & 28 17 05.0 &  0.888 &  3  & -24.83 & QSO & 19.14 & -0.87 &  0.71  & 1.69 & 2.37  &  4.45\\\\% 21261 10 &  13 38 14.68 & 28 46 12.7 &  0.376 &  4  & -22.65 &Sey1 & 19.23 & -0.62 &  0.55  & 0.50 & 2.49  &  8.56\\\\% 10927 11 &  13 38 26.47 & 28 36 37.8 &  1.534 &  3  & -26.06 & QSO & 19.28 & -0.55 &  0.35  & 0.54 & 1.65  &  2.09\\\\% 12558 12 &  13 38 44.36 & 29 01 49.2 &  1.052 &  2  & -26.02 & QSO & 18.37 & -0.64 &  0.28  & 1.94 & 3.48  & 10.00\\\\%  8747 13 &  13 38 48.93 & 28 54 59.0 &  1.281 &  4  & -25.38 & QSO & 19.50 & -0.89 &  0.23  & 1.00 & 1.94  &  3.39\\\\%  9733 14 &  13 38 52.81 & 29 33 35.3 &  1.322 &  4  & -25.61 & QSO & 19.35 & -0.78 &  0.22  & 2.33 & 1.91  &  3.55\\\\%  4148 15 &  13 39 01.52 & 29 19 11.0 &  0.137:&  5  & -21.61 &NELG & 18.06 &  0.86 &  1.13  & 7.11 & 1.22  &\t  *\\\\%  6270 16 &  13 39 13.30 & 27 18 18.5 &  0.680 &  3  & -24.25 & QSO & 19.01 & -0.55 &  0.12  & 0.79 & 1.65  &  1.81\\\\% 31613 17 &  13 39 18.46 & 29 29 52.4 &  1.743 &  4  & -26.20 & QSO & 19.49 & -0.70 &  0.13  & 2.45 & 2.40  &  3.10\\\\%  4684 18 &  13 39 19.27 & 28 39 08.9 &  2.514 &  3  & -27.86 & QSO & 18.98 & -0.31 &  0.27  & 0.69 & 1.91  &  3.74\\\\% 12061 19 &  13 39 58.17 & 29 37 54.0 &  0.441 &  6  & -22.73 &Sey1 & 19.48 & -0.44 &  0.27  & 2.68 & 1.69  &  2.43\\\\%  3485 20 &  13 40 11.33 & 29 24 00.5 &  1.527 &  4  & -25.92 & QSO & 19.41 & -0.43 & -0.39  & 1.80 & 6.06  & 34.40\\\\%  5562 21 &  13 40 33.34 & 29 23 16.4 &  0.841 &  2  & -25.11 & QSO & 18.72 & -0.81 &  0.34  & 1.01 & 3.00  &  4.19\\\\%  5674 22 &  13 41 04.99 & 29 17 08.0 &  1.100 &  3  & -25.17 & QSO & 19.32 & -0.96 &  0.43  & 2.25 & 2.09  &  2.93\\\\%  6597 23 &  13 41 05.08 & 28 38 29.7 &  0.580 &  5  & -23.50 & QSO & 19.36 & -0.61 &  0.26  & 1.68 & 1.71  &  1.87\\\\% 12223 24 &  13 41 07.37 & 28 39 35.7 &  1.569 &  4  & -26.50 & QSO & 18.89 & -0.59 &  0.40  & 0.30 & 2.80  &  5.79\\\\% 12006 25 &  13 41 29.81 & 29 42 31.0 &  0.330 &  3  & -22.61 &Sey1 & 19.00 & -0.65 &  0.55  & 1.65 & 1.87  &  1.96\\\\%  2795 26 &  13 42 00.57 & 29 47 31.9 &  0.190 &  5  & -20.86 &NELG & 19.53 &  0.09 &  1.27  & 1.07 & 2.45  &  1.28\\\\%  2016 27 &  13 42 01.08 & 29 00 33.4 &  1.016 &  4  & -24.93 & QSO & 19.38 & -0.82 &  0.28  & 0.46 & 2.56  &  7.18\\\\%  8951 28 &  13 42 08.01 & 29 44 34.7 &  1.528 &  5  & -24.90 & QSO & 20.42 &     * &  1.13  & 3.32 & 0.95  &     *\\\\%  2465 29 &  13 42 08.28 & 27 09 31.0 &  1.190 &  3  & -26.55 & QSO & 18.13 & -0.82 &  0.38  & 0.67 & 3.21  & 14.90\\\\% 32910 30 &  13 42 29.81 & 28 42 48.9 &  1.184 &  5  & -24.93 & QSO & 19.74 & -0.77 &  0.70  & 1.31 & 3.49  &  6.29\\\\% 11479 31 &  13 42 30.41 & 28 56 26.4 &  1.264 &  4  & -25.34 & QSO & 19.50 & -0.88 &  0.21  & 0.83 & 1.88  &  4.06\\\\%  9561 32 &  13 42 30.90 & 28 37 25.1 &  0.589 &  3  & -24.75 & QSO & 18.14 & -0.71 &  0.28  & 1.71 & 3.42  &  9.15\\\\% 12465 33 &  13 42 32.55 & 28 55 19.3 &  2.172 &  5  & -26.77 & QSO & 19.58 & -0.69 &  0.27  & 2.13 & 2.50  &  3.14\\\\%  9709 34 &  13 42 39.19 & 29 36 13.4 &  0.235 &  3  & -21.48 &Sey1 & 19.36 & -0.84 &  0.73  & 2.80 & 5.66  & 10.70\\\\%  3775 35 &  13 42 53.05 & 29 33 42.8 &  1.939 &  4  & -26.66 & QSO & 19.37 & -0.82 &  0.56  & 2.00 & 2.39  &  1.85\\\\%  4167 36 &  13 42 58.55 & 27 54 31.2 &  1.799 &  2  & -26.91 & QSO & 18.88 & -0.81 &  0.36  & 1.31 & 2.78  &  4.44\\\\% 26352 37 &  13 43 02.76 & 28 49 43.8 &  1.581 &  2  & -26.57 & QSO & 18.83 & -0.34 &  0.14  & 2.05 & 1.73  &  1.80\\\\% 10474 38 &  13 43 05.17 & 29 09 21.2 &  1.679 &  5  & -26.00 & QSO & 19.58 & -0.74 &  0.15  & 2.14 & 3.17  &  9.97\\\\%  7755 39 &  13 43 24.63 & 29 39 07.1 &  1.332 &  3  & -25.83 & QSO & 19.15 & -0.81 & -0.14  & 0.78 & 2.17  &  4.09\\\\%  3326 40 &  13 43 47.46 & 28 35 07.3 &  0.477 &  4  & -22.97 & QSO & 19.43 & -0.80 &  0.00  & 1.86 & 3.25  &  6.87\\\\% 12880 41 &  13 43 48.08 & 28 23 53.2 &  0.733 &  3  & -24.28 & QSO & 19.18 & -0.52 &  0.47  & 2.29 & 2.30  &  3.73\\\\% 17649 42 &  13 43 50.03 & 28 02 05.7 &  0.211 &  1  & -22.69 &Sey1 & 17.92 & -0.67 &  0.48  & 1.53 & 3.81  &  5.24\\\\% 25116 43 &  13 43 58.99 & 28 35 48.8 &  0.193 &  4  & -21.04 &NELG & 19.38 & -0.55 &  1.11  & 1.33 & 1.83  &  2.34\\\\% 12755 44 &  13 43 59.43 & 29 29 33.3 &  2.109 &  5  & -26.85 & QSO & 19.44 & -0.71 &  0.58  & 1.47 & 1.51  &  1.08\\\\%  4752 45 &  13 44 03.07 & 29 20 05.3 &  1.506 &  4  & -25.79 & QSO & 19.50 & -0.43 &  0.08  & 0.53 & 4.19  &  4.65\\\\%  6156 46 &  13 44 12.71 & 28 52 16.0 &  2.375 &  2  & -28.20 & QSO & 18.40 & -0.33 &  0.05  & 0.94 & 1.36  &  2.34\\\\% 10103 47 &  13 44 22.39 & 28 56 42.9 &  0.481 &  3  & -23.07 & QSO & 19.35 & -0.58 & -0.14  & 1.50 & 4.21  &  9.53\\\\%  9523 48 &  13 44 33.95 & 28 27 23.7 &  1.007 &  3  & -24.93 & QSO & 19.36 & -0.81 &  0.50  & 0.92 & 2.48  &  2.69\\\\% 15572 49 &  13 45 04.05 & 28 48 18.6 &  1.275 &  3  & -25.59 & QSO & 19.27 & -0.81 &  0.30  & 2.59 & 1.73  &  1.98\\\\% 10669 50 &  13 45 11.65 & 28 13 60.0 &  1.053 &  5  & -24.71 & QSO & 19.68 & -0.92 &  0.53  & 1.85 & 1.89  &  3.00\\\\% 22609 51 &  13 45 13.81 & 29 06 03.0 &  0.587 &  2  & -24.53 & QSO & 18.36 & -0.47 & -0.01  & 0.82 & 4.19  & 11.00\\\\%  8192 52 &  13 45 38.38 & 28 49 35.5 &  0.454 &  2  & -23.56 & QSO & 18.74 & -0.22 &  0.53  & 0.59 & 1.17  &  1.27\\\\% 10481 53 &  13 45 47.42 & 28 19 09.8 &  0.457 &  4  & -22.91 & QSO & 19.39 & -0.40 &  0.32  & 2.71 & 1.72  &  1.81\\\\% 20225 54 &  13 45 57.82 & 28 32 06.7 &  1.700 &  2  & -26.64 & QSO & 18.97 & -0.75 &  0.30  & 1.08 & 2.06  &  3.09\\\\% 13610 55 &  13 46 22.15 & 28 52 14.4 &  0.758 &  5  & -23.77 & QSO & 19.78 & -0.51 &  0.51  & 1.12 & 3.21  &  3.82\\\\% 10086 \\bottomrule \\end{tabular} \\caption{\\label{QSO_list} QSOs, Seyfert\\,1, and NELGs from the follow-up spectroscopy of the present study. $I_{\\rm pm},\\, I_{\\rm var},\\, I_{\\rm ltvar}$ are the indices for proper motion, overall variability, and long-term variability, respectively. A colon behind the redshift symbolises uncertain data, an asterisk indicates missing data. (Note that $I_{\\rm ltvar}$ has been computed only for objects with $B<20$ and $I_{\\rm pm}<4$.)} } \\end{table*}  \\begin{table*}[ht] {\\renewcommand{\\baselinestretch}{0.95}\\footnotesize \\begin{tabular}{rrrrrrrrrrrrr} \\toprule No. & $\\alpha$\\,(J2000) & $\\delta$\\,(J2000) & \\multicolumn{1}{c}{$z$} & run & $M_B$ & type & \\multicolumn{1}{c}{$B$} & $U-B$ & $B-V$ & $I_{\\rm pm}$ & $I_{\\rm var}$ & $I_{\\rm ltvar}$\\\\ \\midrule 56 &  13 46 38.31 & 29 45 54.7 &  1.336 &  4  & -25.58 & QSO & 19.41 & -0.77 &  0.19  & 1.06 & 1.93  &  4.83\\\\%  2232 57 &  13 46 48.04 & 29 46 21.0 &  1.594 &  5  & -25.90 & QSO & 19.52 & -0.62 & -0.08  & 0.68 & 1.88  &  2.43\\\\%  2164 58 &  13 47 04.33 & 29 35 23.3 &  0.625 &  2  & -24.40 & QSO & 18.65 & -0.15 &  0.25  & 1.41 & 1.55  &  2.87\\\\%  3866 59 &  13 47 14.11 & 28 00 08.0 &  1.566 &  3  & -26.09 & QSO & 19.29 & -0.43 & -0.14  & 2.97 & 1.59  &  2.47\\\\% 25407 60 &  13 47 15.42 & 29 26 23.9 &  1.918 &  2  & -27.25 & QSO & 18.75 & -0.85 &  0.11  & 1.14 & 3.50  &  4.54\\\\%  5171 61 &  13 47 23.48 & 29 27 22.2 &  1.223 &  2  & -25.78 & QSO & 18.97 & -0.74 & -0.06  & 1.93 & 2.07  &  3.83\\\\%  5038 62 &  13 47 35.79 & 29 00 12.3 &  1.398 &  3  & -25.85 & QSO & 19.23 & -0.71 &  0.12  & 0.65 & 1.56  &  2.40\\\\%  8960 63 &  13 47 37.92 & 27 14 50.5 &  0.787 &  5  & -24.07 & QSO & 19.58 & -0.44 &  0.14  & 2.01 & 2.78  &  5.67\\\\% 32067 64 &  13 47 42.88 & 28 56 52.1 &  0.985 &  3  & -24.97 & QSO & 19.28 & -0.72 &  0.64  & 1.85 & 1.95  &  3.76\\\\%  9454 65 &  13 47 56.49 & 28 35 20.7 &  2.458 &  3  & -27.75 & QSO & 18.98 & -0.27 & -0.12  & 0.55 & 2.92  &  3.20\\\\% 12783 66 &  13 47 58.90 & 27 24 50.7 &  1.048 &  2  & -26.16 & QSO & 18.22 & -0.38 &  0.09  & 3.35 & 2.12  &  2.10\\\\% 30618 67 &  13 48 02.84 & 28 27 15.9 &  0.864 &  2  & -25.57 & QSO & 18.32 & -0.52 &  0.03  & 0.50 & 2.51  &  3.22\\\\% 15460 68 &  13 48 04.29 & 28 40 25.0 &  2.471 &  1  & -28.91 & QSO & 17.85 & -0.08 & -0.11  & 1.33 & 3.32  &  6.88\\\\% 11804 69 &  13 48 08.94 & 28 02 13.4 &  1.643 &  1  & -27.55 & QSO & 17.96 & -0.32 &  0.03  & 0.88 & 1.63  &  2.00\\\\% 25035 70 &  13 48 11.70 & 28 18 01.5 &  2.941 &  1  & -29.87 & QSO & 17.72 &  0.17 &  0.24  & 2.22 & 1.81  &  1.14\\\\% 20727 71 &  13 48 14.87 & 29 02 56.3 &  0.938 &  3  & -24.76 & QSO & 19.35 & -0.81 & -0.20  & 2.26 & 4.28  & 11.50\\\\%  8580 72 &  13 48 18.75 & 29 10 41.9 &  1.932 &  4  & -26.79 & QSO & 19.24 & -0.77 & -0.09  & 0.90 & 3.11  &  4.56\\\\%  7496 73 &  13 48 20.29 & 29 34 03.5 &  2.041 &  6  & -26.56 & QSO & 19.65 & -0.76 &  0.30  & 1.41 & 2.26  &  2.07\\\\%  4049 74 &  13 48 22.58 & 28 39 43.2 &  1.140 &  4  & -25.36 & QSO & 19.21 & -0.29 & -0.08  & 1.02 & 2.23  &  3.59\\\\% 11924 75 &  13 48 24.33 & 28 32 50.3 &  2.141 &  5  & -26.76 & QSO & 19.57 & -0.06 &  0.62  & 0.94 & 1.98  &  2.08\\\\% 13333 76 &  13 48 26.47 & 28 43 17.2 &  2.031 &  5  & -27.47 & QSO & 18.72 & -0.81 &  0.04  & 3.53 & 1.67  &  1.26\\\\% 11359 77 &  13 48 26.60 & 29 06 22.7 &  2.819 &  3  & -28.17 & QSO & 19.18 &  0.02 &  0.08  & 1.22 & 1.85  &  1.77\\\\%  8096 \\bottomrule \\end{tabular} \\addtocounter{table}{-1} \\caption{ continued} } \\end{table*} \\afterpage  \\begin{table*}[htbp] {\\renewcommand{\\baselinestretch}{0.95}\\footnotesize \\begin{tabular}{rlllrrrrrrrrr} \\toprule No. & $\\alpha$\\,(J2000) & $\\delta$\\,(J2000) & \\multicolumn{1}{c}{$z$} & $M_B$ & type & \\multicolumn{1}{c}{$B$} & $U-B$ & $B-V$ & $I_{\\rm pm}$ & $I_{\\rm var}$ & $I_{\\rm ltvar}$\\\\ \\midrule 1 & 13 35 11.82 & 27 56 51.7 & 2.425 &  -26.83  & QSO & 19.85 & -0.74 &      * &   2.07 &   2.02 &  3.33\\\\%  25856 2 & 13 35 23.65 & 28 08 38.9 & 0.904 &  -24.15  & QSO & 19.87 & -0.87 &   0.60 &   1.60 &   1.80 &  2.38\\\\%  23814 3 & 13 35 24.10 & 27 17 21.0 & 0.879 &  -24.21  & QSO & 19.74 & -0.78 &   0.33 &   3.61 &   1.14 &  2.65\\\\%  31661 4 & 13 35 35.92 & 27 21 23.8 & 0.783 &  -24.55  & QSO & 19.08 & -0.58 &   0.10 &   1.97 &   2.54 &  1.88\\\\%  31062 5 & 13 35 39.68 & 28 05 04.7 & 1.095 &  -24.22  & QSO & 20.26 & -1.02 &   0.52 &   1.55 &   1.36 &\t *\\\\%  24479 6 & 13 35 46.10 & 28 05 51.1 & 2.960 &  -28.30  & QSO & 19.32 & -0.12 &   0.81 &   1.97 &   1.53 &  1.29\\\\%  24360 7 & 13 35 48.08 & 27 28 35.5 & 1.331 &  -25.10  & QSO & 19.88 & -0.96 &   0.44 &   2.71 &   1.82 &  1.73\\\\%  30028 8 & 13 35 54.38 & 27 53 23.2 & 1.886 &  -26.64  & QSO & 19.30 & -0.81 &   0.60 &   1.79 &   1.32 &  2.42\\\\%  26420 9 & 13 36 02.73 & 27 27 51.8 & 1.117 &  -24.78  & QSO & 19.74 & -1.10 &   0.37 &   2.46 &   2.23 &  4.04\\\\%  30140 10 & 13 36 05.91 & 28 13 24.1 & 2.388 &  -27.58  & QSO & 19.04 & -0.32 &   0.21 &   2.92 &   1.54 &  1.52\\\\%  22627 11 & 13 36 11.38 & 27 10 13.0 & 2.425 &  -28.74  & QSO & 17.94 & -0.47 &   0.21 &   3.52 &   1.68 &  2.71\\\\%  32692 12 & 13 36 13.21 & 28 24 58.6 & 1.908 &  -27.17  & QSO & 18.81 & -1.19 &   0.36 &   1.18 &   1.24 &  1.82\\\\%  16656 13 & 13 36 16.62 & 28 04 52.8 & 0.271 &  -21.81  & Sey1& 19.35 & -0.71 &   0.51 &   2.15 &   1.47 &  2.51\\\\%  24532 14 & 13 36 23.02 & 27 07 20.2 & 1.931 &  -26.14  & QSO & 19.88 & -0.78 &   0.22 &   1.02 &   1.85 &  3.78\\\\%  33196 15 & 13 36 30.17 & 27 01 01.3 & 0.283 &  -21.08  & NELG& 20.17 & -0.81 &   0.71 &   1.79 &   1.66 &\t *\\\\%  34158 16 & 13 36 43.50 & 27 29 53.6 & 0.780 &  -23.83  & QSO & 19.56 & -1.02 &  -0.29 &   2.21 &   3.88 & 41.80\\\\%  29860 17 & 13 36 43.59 & 27 10 59.3 & 1.440 &  -25.02  & QSO & 20.17 & -1.22 &   0.39 &   1.74 &   2.67 &\t *\\\\%  32596 18 & 13 36 50.06 & 27 20 54.8 & 1.950 &  -26.11  & QSO & 19.94 & -0.93 &   0.37 &   1.06 &   1.01 &  1.21\\\\%  31175 19 & 13 36 56.79 & 27 25 43.0 & 1.360 &  -26.21  & QSO & 18.83 & -0.87 &  -0.01 &   3.86 &   4.89 &  9.90\\\\%  30470 20 & 13 37 00.88 & 27 13 22.8 & 2.074 &  -26.27  & QSO & 19.98 & -0.76 &   0.53 &   0.94 &   1.95 &  3.25\\\\%  32256 21 & 13 37 14.12 & 27 12 49.8 & 1.909 &  -26.47  & QSO & 19.51 & -0.99 &   0.55 &   1.02 &   1.68 &  1.50\\\\%  32334 22 & 13 37 17.39 & 27 01 07.4 & 0.637 &  -24.45  & QSO & 18.64 & -0.43 &   0.35 &   2.41 &   3.02 &  3.02\\\\%  34165 23 & 13 37 19.98 & 28 09 26.2 & 0.460 &  -22.69  & Sey1& 19.62 & -0.84 &   0.28 &   1.38 &   1.70 &  2.80\\\\%  23707 24 & 13 37 26.09 & 27 36 39.1 & 1.121 &  -25.21  & QSO & 19.32 & -0.72 &   0.29 &   1.22 &   1.59 &  3.84\\\\%  28903 25 & 13 37 35.34 & 27 03 11.2 & 1.762 &  -25.58  & QSO & 20.14 & -0.77 &   0.76 &   1.31 &   2.46 &\t *\\\\%  33856 26 & 13 37 37.95 & 28 13 47.1 & 1.865 &  -25.72  & QSO & 20.13 & -1.11 &   0.58 &   1.66 &   1.51 &\t *\\\\%  22579 27 & 13 37 39.07 & 28 04 46.8 & 1.124 &  -24.24  & QSO & 20.30 & -1.14 &   0.55 &   1.96 &   1.60 &\t *\\\\%  24578 28 & 13 37 44.30 & 27 01 29.2 & 1.928 &  -27.21  & QSO & 18.81 & -0.84 &   0.26 &   3.79 &   3.56 &  8.39\\\\%  34113 29 & 13 37 49.44 & 28 04 01.0 & 1.321 &  -25.19  & QSO & 19.77 & -0.86 &   0.21 &   2.48 &   2.08 &  2.01\\\\%  24722 30 & 13 37 54.57 & 28 09 44.0 & 1.592 &  -26.32  & QSO & 19.10 & -0.60 &   0.22 &   1.39 &   1.96 &  2.31\\\\%  23665 31 & 13 37 54.90 & 27 37 55.6 & 2.362 &  -26.34  & QSO & 20.24 & -1.03 &   0.74 &   1.47 &   1.90 &\t *\\\\%  28747 32 & 13 37 55.12 & 27 22 44.4 & 0.433 &  -21.69  & NELG& 20.48 & -0.97 &   0.60 &   1.23 &   1.29 &\t *\\\\%  30908 33 & 13 38 04.56 & 27 48 34.4 & 0.814 &  -23.33  & QSO & 20.41 & -0.95 &   0.66 &   2.47 &   1.40 &\t *\\\\%  27252 34 & 13 38 21.80 & 29 14 45.7 & 0.646 &  -26.02  & QSO & 17.11 & -0.62 &   0.17 &   1.50 &   2.19 &  6.49\\\\%\t6923 35 & 13 38 29.51 & 28 02 56.4 & 1.116 &  -24.14  & QSO & 20.38 & -1.01 &   0.84 &   1.80 &   1.25 &\t *\\\\%  24927 36 & 13 38 34.08 & 27 30 09.8 & 0.342 &  -21.89  & Sey1& 19.80 & -0.47 &   0.14 &   0.87 &   1.72 &  4.46\\\\%  29857 37 & 13 38 49.26 & 27 49 15.7 & 1.310 &  -24.71  & QSO & 20.23 & -0.94 &   0.66 &   0.51 &   0.96 &\t *\\\\%  27162 38 & 13 38 55.70 & 27 13 05.3 & 0.486 &  -22.42  & Sey1& 20.02 & -0.52 &   0.18 &   1.42 &   1.94 &\t *\\\\%  32327 39 & 13 38 55.91 & 27 48 38.4 & 1.325 &  -25.72  & QSO & 19.25 & -0.67 &   0.35 &   1.21 &   2.88 &  5.78\\\\%  27256 40 & 13 38 57.86 & 27 11 50.3 & 1.922 &  -25.70  & QSO & 20.30 & -1.03 &   0.41 &   0.38 &   0.68 &\t *\\\\%  32514 41 & 13 39 00.98 & 28 14 24.8 & 2.513 &  -27.43  & QSO & 19.40 & -0.17 &   0.06 &   1.89 &   1.68 &  3.14\\\\%  22431 42 & 13 39 02.19 & 27 36 43.2 & 2.530 &  -27.11  & QSO & 19.76 & -0.47 &   0.56 &   0.89 &   1.84 &  2.86\\\\%  28927 43 & 13 39 09.63 & 28 05 26.5 & 1.113 &  -25.29  & QSO & 19.23 & -0.93 &   0.26 &   0.42 &   1.97 &  2.22\\\\%  24501 44 & 13 39 16.57 & 27 28 16.1 & 1.047 &  -25.16  & QSO & 19.22 & -0.71 &   0.55 &   1.95 &   2.61 &  4.25\\\\%  30152 45 & 13 39 24.84 & 27 23 35.8 & 1.056 &  -24.50  & QSO & 19.90 & -0.83 &   0.49 &   1.25 &   1.01 &  1.20\\\\%  30814 46 & 13 39 38.80 & 27 15 29.3 & 1.750 &  -25.22  & QSO & 20.48 & -0.82 &   0.43 &   0.63 &   1.19 &\t *\\\\%  32006 47 & 13 39 45.89 & 27 11 00.4 & 1.120 &  -24.65  & QSO & 19.88 & -0.82 &   0.46 &   2.61 &   1.71 &  4.83\\\\%  32649 48 & 13 39 55.56 & 27 47 58.7 & 0.657 &  -24.06  & QSO & 19.11 & -0.45 &   0.49 &   1.17 &   1.35 &  2.00\\\\%  27348 49 & 13 40 00.10 & 28 14 04.4 & 1.760 &  -25.74  & QSO & 19.98 & -0.93 &   0.36 &   1.43 &   1.42 &  1.60\\\\%  22573 50 & 13 40 04.87 & 28 16 53.2 & 2.517 &  -29.45  & QSO & 17.39 & -0.29 &   0.20 &   3.58 &   1.29 &  2.88\\\\%  21467 51 & 13 40 09.35 & 27 18 39.6 & 0.327 &  -21.92  & Sey1& 19.66 & -0.74 &   0.51 &   1.52 &   2.03 &  1.77\\\\%  31581 52 & 13 40 13.61 & 28 15 20.7 & 1.467 &  -25.35  & QSO & 19.88 & -0.85 &   0.40 &   3.10 &   1.92 &  2.26\\\\%  22138 53 & 13 40 20.41 & 27 20 41.7 & 1.140 &  -24.63  & QSO & 19.94 & -0.76 &   0.67 &   1.26 &   1.70 &  0.59\\\\%  31279 54 & 13 40 22.78 & 27 40 58.7 & 0.172 &  -21.19  & NELG& 18.98 & -0.81 &   0.72 &   1.67 &   7.43 & 12.60\\\\%  28342 55 & 13 40 31.59 & 27 05 41.9 & 0.280 &  -21.59  & Sey1& 19.65 & -0.66 &   0.96 &   1.40 &   0.96 &  0.97\\\\%  33518 \\bottomrule \\end{tabular} \\caption{\\label{QSO_NED_list} As Table\\,\\ref{QSO_list} but for the objects from our basic candidate sample that were identified with QSOs, Seyfert\\,1s, or NELGs from the NED.} } \\end{table*}  \\begin{table*}[ht] {\\renewcommand{\\baselinestretch}{0.95}\\footnotesize \\begin{tabular}{rrrrrrrrrrrrr} \\toprule No. & $\\alpha$\\,(J2000) & $\\delta$\\,(J2000) & \\multicolumn{1}{c}{$z$} & $M_B$ &type & \\multicolumn{1}{c}{$B$} & $U-B$ & $B-V$ & $I_{\\rm pm}$ & $I_{\\rm var}$ & $I_{\\rm ltvar}$\\\\ \\midrule 56 & 13 40 48.23 & 27 40 09.3 & 1.900  &  -25.57  & QSO & 20.39 & -0.90 &   0.51 &   1.48 &   1.47 &\t  * \\\\%   28457 57 & 13 40 54.50 & 27 17 53.7 & 1.252  &  -24.80  & QSO & 20.02 & -0.93 &   0.21 &   1.43 &   2.13 &\t  * \\\\%   31684 58 & 13 41 00.11 & 27 25 17.9 & 1.175  &  -24.19  & QSO & 20.46 & -1.04 &   0.65 &   1.27 &   1.87 &\t  * \\\\%   30582 59 & 13 41 16.27 & 28 16 09.2 & 1.310  &  -25.03  & QSO & 19.91 & -0.30 &   0.58 &   0.53 &   1.20 &  2.28 \\\\%   21818 60 & 13 41 19.35 & 27 29 26.7 & 0.480  &  -22.29  & QSO & 20.12 & -0.72 &   0.41 &   0.80 &   1.87 &\t  * \\\\%   30000 61 & 13 41 23.26 & 27 49 55.3 & 1.045  &  -25.36  & QSO & 19.01 & -0.56 &   0.28 &   4.24 &   2.94 &\t  * \\\\%   27096 62 & 13 41 53.82 & 27 35 56.2 & 1.552  &  -25.03  & QSO & 20.33 & -0.85 &   0.40 &   0.24 &   1.24 &\t  * \\\\%   29070 63 & 13 42 03.24 & 26 58 46.7 & 0.530  &  -23.18  & QSO & 19.47 & -0.60 &   0.44 &   2.30 &   1.46 &  1.75 \\\\%   34601 64 & 13 42 11.61 & 28 28 49.8 & 0.330  &  -22.29  & Sey1& 19.31 & -0.19 &   0.53 &   1.13 &   0.84 &  1.31 \\\\%   14851 65 & 13 42 14.31 & 27 37 55.0 & 0.770  &  -23.68  & QSO & 19.90 & -0.84 &   0.38 &   1.44 &   1.17 &  0.89 \\\\%   28789 66 & 13 42 21.55 & 27 26 21.4 & 2.240  &  -26.95  & QSO & 19.47 & -0.69 &   0.38 &   0.40 &   1.52 &  1.50 \\\\%   30446 67 & 13 42 23.62 & 27 04 34.6 & 2.827  &  -27.63  & QSO & 19.75 &  0.07 &   0.06 &   1.57 &   1.37 &  2.61 \\\\%   33711 68 & 13 42 24.34 & 27 20 40.0 & 0.704  &  -25.29  & QSO & 18.05 & -0.45 &   0.34 &   0.85 &   1.90 &  2.32 \\\\%   31302 69 & 13 42 37.30 & 27 15 41.4 & 1.229  &  -24.34  & QSO & 20.43 & -0.89 &   0.66 &   1.70 &   2.24 &\t  * \\\\%   31999 70 & 13 42 44.42 & 27 37 37.2 & 1.721  &  -25.65  & QSO & 20.00 & -0.89 &   0.29 &   2.03 &   1.65 &\t  * \\\\%   28832 71 & 13 42 52.90 & 27 36 15.8 & 1.444  &  -25.13  & QSO & 20.07 & -0.79 &   0.32 &   0.54 &   2.46 &\t  * \\\\%   29020 72 & 13 42 54.02 & 27 33 10.0 & 0.810  &  -24.72  & QSO & 19.01 & -0.62 &   0.19 &   0.68 &   2.02 &  3.16 \\\\%   29443 73 & 13 42 54.30 & 28 28 05.7 & 1.037  &  -25.98  & QSO & 18.37 & -0.65 &   0.35 &   1.57 &   2.31 &  3.55 \\\\%   15224 74 & 13 42 54.72 & 26 58 04.9 & 2.737  &  -27.43  & QSO & 19.79 & -0.33 &   0.70 &   3.21 &   1.47 &  0.92 \\\\%   34708 75 & 13 42 55.06 & 27 53 31.5 & 1.527  &  -26.18  & QSO & 19.14 & -0.56 &  -0.14 &   0.98 &   3.03 &  8.27 \\\\%   26516 76 & 13 42 59.18 & 27 47 23.1 & 1.164  &  -25.13  & QSO & 19.50 & -0.71 &   0.44 &   1.72 &   4.99 &  3.89 \\\\%   27472 77 & 13 43 00.10 & 28 44 07.4 & 0.905  &  -26.88  & QSO & 17.14 & -0.75 &   0.37 &   1.87 &   2.55 &  3.94 \\\\%   11281 78 & 13 43 14.11 & 27 38 12.3 & 2.491  &  -26.67  & QSO & 20.12 & -0.47 &   0.44 &   2.63 &   0.82 &\t  * \\\\%   28755 79 & 13 43 23.05 & 26 57 16.5 & 1.274  &  -25.54  & QSO & 19.32 & -0.97 &   0.58 &   3.28 &   2.18 &  3.04 \\\\%   34820 80 & 13 43 30.20 & 27 44 55.3 & 2.424  &  -26.56  & QSO & 20.12 & -0.51 &   0.19 &   0.31 &   1.51 &\t  * \\\\%   27816 81 & 13 44 05.32 & 27 26 33.4 & 1.312  &  -25.13  & QSO & 19.81 & -0.98 &   0.42 &   1.39 &   2.24 &  1.79 \\\\%   30409 82 & 13 44 05.78 & 28 00 04.7 & 0.733  &  -23.70  & QSO & 19.75 & -0.55 &   0.50 &   1.36 &   1.07 &  1.34 \\\\%   25453 83 & 13 44 26.34 & 27 58 45.0 & 0.900  &  -24.24  & QSO & 19.77 & -0.79 &   0.55 &   1.02 &   1.83 &  1.50 \\\\%   25693 84 & 13 44 51.91 & 28 11 13.8 & 2.401  &  -27.61  & QSO & 19.03 & -0.18 &   0.20 &   0.54 &   1.47 &  1.66 \\\\%   23398 85 & 13 45 15.59 & 27 26 17.3 & 1.884  &  -26.71  & QSO & 19.23 & -0.70 &   0.09 &   2.25 &   1.84 &  3.14 \\\\%   30443 86 & 13 45 40.92 & 27 55 24.6 & 0.453  &  -23.14  & QSO & 19.15 & -0.48 &  -0.03 &   2.41 &   2.58 &  8.77 \\\\%   26177 87 & 13 45 42.81 & 27 22 19.5 & 1.183  &  -25.69  & QSO & 18.98 & -0.72 &   0.36 &   2.42 &   2.39 &  2.57 \\\\%   31019 88 & 13 45 46.05 & 28 09 17.0 & 2.225  &  -26.62  & QSO & 19.79 & -0.50 &   0.20 &   0.89 &   1.43 &  1.33 \\\\%   23817 89 & 13 45 54.54 & 27 41 01.4 & 1.038  &  -25.25  & QSO & 19.11 & -0.71 &   0.46 &   1.81 &   2.13 &  1.64 \\\\%   28330 90 & 13 46 06.68 & 27 11 22.2 & 1.806  &  -25.67  & QSO & 20.13 & -1.05 &   0.52 &   2.05 &   1.86 &\t  * \\\\%   32595 91 & 13 46 14.57 & 28 13 52.5 & 0.659  &  -24.75  & QSO & 18.43 & -0.39 &   0.21 &   1.62 &   2.37 &  4.81 \\\\%   22629 92 & 13 46 33.56 & 27 05 58.3 & 2.340  &  -27.23  & QSO & 19.32 & -0.21 &   0.16 &   0.54 &   1.60 &  1.07 \\\\%   33455 93 & 13 46 38.35 & 27 57 41.2 & 2.439  &  -27.61  & QSO & 19.09 &  0.03 &   0.15 &   2.07 &   1.64 &  1.73 \\\\%   25829 94 & 13 46 42.71 & 27 16 12.9 & 1.612  &  -25.78  & QSO & 19.67 & -0.49 &   0.28 &   1.18 &   2.19 &  2.65 \\\\%   31903 95 & 13 46 44.32 & 28 01 30.0 & 1.127  &  -25.68  & QSO & 18.86 & -0.69 &   0.45 &   1.30 &   3.84 &  6.32 \\\\%   25190 96 & 13 47 02.80 & 26 59 35.1 & 1.504  &  -25.63  & QSO & 19.66 & -0.50 &   0.26 &   1.47 &   2.62 &  3.06 \\\\%   34457 97 & 13 47 03.74 & 27 09 24.8 & 1.932  &  -26.70  & QSO & 19.32 & -0.69 &   0.39 &   1.47 &   3.56 &  1.26 \\\\%   32887 98 & 13 47 05.51 & 28 18 05.0 & 0.255  &  -21.20  & NELG& 19.83 & -0.36 &   1.17 &   0.57 &   0.92 &  1.52 \\\\%   20773 99 & 13 47 10.23 & 28 03 54.6 & 0.992  &  -25.26  & QSO & 19.00 & -0.45 &   0.29 &   1.05 &   1.85 &  1.36 \\\\%   24764 100 & 13 47 16.51 & 27 14 20.1 & 2.530  &  -27.21  & QSO & 19.65 & -0.41 &   0.34 &   3.21 &   1.65 &  1.52 \\\\%   32151 101 & 13 47 26.33 & 27 04 33.1 & 2.212  &  -27.66  & QSO & 18.74 & -0.24 &   0.31 &   3.07 &   1.75 &  3.07 \\\\%   33668 102 & 13 47 30.86 & 27 13 58.5 & 2.118  &  -26.68  & QSO & 19.62 & -0.68 &   0.42 &   1.08 &   1.64 &  2.80 \\\\%   32198 103 & 13 48 05.16 & 27 47 13.1 & 1.430  &  -26.14  & QSO & 19.02 & -0.53 &   0.10 &   1.87 &   1.99 &  2.12 \\\\%   27430 104 & 13 48 25.07 & 27 06 16.4 & 2.600  &  -29.33  & QSO & 17.66 & -0.20 &  -0.06 &   1.26 &   2.58 &  8.82 \\\\%   33384 \\bottomrule \\end{tabular} \\addtocounter{table}{-1} \\caption{ continued} } \\end{table*}"
    },
    "0207/astro-ph0207165_arXiv.txt": {
        "abstract": "}[1]{{ \\footnotesize \\noindent {\\bf Abstract} #1\\\\}}  \\renewcommand{\\author}[1]{\\subsection*{#1}} \\newcommand{\\address}[1]{\\subsection*{\\it#1}}  \\begin{document}    \\chapter*{Galaxy Clusters with Chandra\\footnote{Contribution to XIII Rencontres de Blois 2001, ed. L. M. Celnikier}} \\author{W. Forman$^1$, C. Jones$^1$, M. Markevitch$^{1,3}$, A. Vikhlinin$^{1,3}$, \\& E. Churazov$^{2,3}$}    \\address{1) Smithsonian Astrophysical Observatory, Cambridge, MA, USA\\\\ 2) MPI fur Astrophysik, Garching, Germany\\\\ 3) Space Research Institute, Moscow, Russia}   \\abstract{We discuss Chandra results related to 1) cluster mergers and cold fronts and 2) interactions between relativistic plasma and hot cluster atmospheres. We describe the properties of cold fronts using NGC1404 in the Fornax cluster and A3667 as examples. We discuss multiple surface brightness discontinuities in the cooling flow cluster ZW3146. We review the supersonic merger underway in CL0657. Finally, we summarize the interaction between plasma bubbles produced by AGN and hot gas using M87 and NGC507 as examples.} ",
        "introduction": " ",
        "conclusions": "We did not expect the rich variety of new structures seen in the Chandra high angular resolution observations of clusters and early type galaxies.  Instead of confirming our prejudices, Chandra has brought us a wealth of new information on the interaction of radio sources with the hot gas in both galaxy and cluster atmospheres. We see ``edges'' in many systems with hot and cold gas in close proximity and have been able to extract important new parameters of the ICM.  We have only barely begun to digest the import of the Chandra cluster and galaxy observations.  We can only expect the unexpected as Chandra observations continue and as our understanding of how best to use this new observatory matures.  We acknowledge support from NASA contract NAS8 39073, NASA grants NAG5-3065 and NAG5-6749 and the Smithsonian Institution.   \\centerline{\\Large \\bf Note Added on ZW3146}  One of the puzzles relating to the X-ray observation of ZW3146 was the apparent misalignment by a few arcseconds between the X-ray brightness peak and the optical center of the central, dominant cD galaxy. The solution to this apparent misalignment was resolved by the FIRST radio image which shows two radio sources with flux densities of $\\sim2$ mJy separated by $14''$ (see Fig.~\\ref{wrfzw3146_new}a). As Fig.~\\ref{wrfzw3146_new}b shows, the northern component of the pair of radio sources is precisely aligned with the optical cD galaxy at the cluster center.  Fig.~\\ref{wrfzw3146_new}c shows that the radio source, and therefore the center of the cD galaxy, lies at the approximate center of irregular X-ray features.  The bright X-ray structure northwest of the cD is elongated perpendicular to the direction to the galaxy center. A second X-ray bright feature lies southeast of the radio peak.  We suggest that the radio emitting plasma formed a bubble that produced a trough in the X-ray surface brightness at the galaxy center. The bubble is surrounded by an irregular shell of X-ray emission. Thus, the brightest X-ray structures do not lie at the galaxy center, but instead form a partial shell surrounding the radio emitting plasma~\\cite{wrfzw3146}.  Such structures are similar to the X-ray surface brightness enhancements seen around radio plasma bubbles in nearby galaxies e.g., M84~\\cite{wrffinoguenov2001}.  \\begin{figure} [bh] \\vspace*{0.2in} \\centerline{\\includegraphics[width=0.8\\textwidth]{formanwnewz.ps} } \\caption{(\\textbf{a - upper left}) The FIRST radio image of ZW3146 shows two radio sources with integrated fluxes of $\\sim2$ mJy. (\\textbf{b - lower left}) Contours of the FIRST radio image superposed on an optical image (S. Allen and A. Edge, private communication) shows that the northern radio source is nearly perfectly aligned on the center of the optical cD galaxy at the cluster center. The second radio source may be a background object, unrelated to the cluster. (\\textbf{c - upper right}) Contours from the FIRST radio image superposed on the Chandra full resolution ($0.492''$ pixel) image show that the radio emission is centered within an irregular shell of X-ray emission. (\\textbf{d - lower right}) Large scale X-ray image of ZW3146.} \\label{wrfzw3146_new} \\end{figure}"
    },
    "0207/astro-ph0207353_arXiv.txt": {
        "abstract": "We develop a new method to estimate the redshift of galaxy clusters through resolved images of the Sunyaev-Zel'dovich effect (SZE). Our method is based on morphological observables which can be measured by actual and future SZE experiments. We test the method with a set of high resolution hydrodynamical simulations of galaxy clusters at different redshifts. Our method combines the observables in a principal component analysis. After {\\it calibrating} the method with an independent redshift estimation for some of the clusters, we show-- using a Bayesian approach-- how the method can give an estimate of the redshift of the galaxy clusters. Although the error bars given by the morphological redshift estimation are large, it should be useful for future SZE surveys where thousands of clusters are expected to be detected; a first preselection of the high redshift candidates could be done using our proposed morphological redshift estimator. Although not considered in this work, our method should also be useful to give an estimate of the redshift of clusters in X-ray and optical surveys. ",
        "introduction": "\\label{section_introduction}  The advent of new experiments dedicated to the observation of the Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich, 1972) (SZE hereafter), demands the development of new techniques to best analyze these new and exciting data. With the SZE it is possible to probe the hot plasma in  galaxy clusters, which shifts the spectrum of the cosmic background radiation. This shift is redshift independent, and it is proportional to the temperature of the plasma and its electron density ($n_e$). This characteristic ($z$-independent, and $\\propto n_e$) makes the SZE an ideal way to explore the high redshift population of galaxy clusters.  However, the fact that the SZE distortion is independent of the redshift of the cluster makes the determination of the redshift of the cluster a challenging task. Redshift information is crucial if one attempts to use cluster surveys to study the evolution of our universe. The evolution of the cluster number counts ($dN/dz$) is a very sensitive indicator of the cosmological model \\citep{eke98,mathiesen98,henry00}. The local abundance of clusters shows a degeneracy in $\\Omega -- \\sigma_8$ \\citep{eke96,bahcall97} but this degeneracy can be broken with an accurate estimation of $dN/dz$ up to moderate or high redshifts ($z \\approx 0.5 - 1$) \\citep[see e.g][]{Bahcall98, borgani01}. The cluster redshift distributions in suitably large SZE cluster surveys can potentially provide precise constraints on the amount and nature of the dark energy in the universe \\citep{haiman01,holder01b,weller01,majumdar02}. Redshifts are also necessary to study the evolution of the cluster structure and dynamics.  One normally determines the redshift using photometric and spectroscopic observations of the galaxies in the cluster. Spectroscopic observations of galaxies in relatively nearby clusters are straightforward, but for distant clusters it is challenging even with the largest telescopes.   For large solid angle surveys, photometric redshifts will be of critical importance, allowing redshift determination for far less time invested at the telescope.  However, photometric redshifts are also time consuming and for clusters above redshift $\\approx$1, photometric redshifts require large telescopes (see for instance Diego et al. 2002 where the authors show the selection function for a galaxy cluster survey with a 10-m telescope and photometric redshift estimations). Future SZE experiments will detect hundreds and perhaps thousands of galaxy clusters. The {\\it Planck} Surveyor SZE survey is expected to detect more than 10$^4$ clusters with redshifts extending to $\\sim$2 \\citep[depending on the cosmological model,][]{diego02}.  A planned, arcminute resolution SZE survey from the South Pole will detect similar numbers of clusters with a much larger fraction at high redshift.  Measuring redshifts for large solid angle, high redshift cluster surveys is a daunting task. An optimal solution may be to combine small and medium--sized telescopes to determine the redshifts of the low and intermediate $z$ clusters, reserving the redshift measurements of the most distant clusters for the largest available telescopes.  Clearly this strategy requires crude {\\it a priori} knowledge of the cluster redshifts. The motivation of the work described here is to examine whether it is possible to make this preselection of the low, intermediate and high redshift clusters using SZE data alone.  Our method is based only on observed SZE properties of galaxy clusters. These include  the observed shape and size of the cluster, which do have some dependence on the redshift.  For instance, the apparent size of a particular cluster will decrease when increasing its redshift.  So, an apparent size will, in principle, constrain the cluster redshift.  However, the apparent size of the cluster also depends on its total mass. Two clusters with different redshifts and masses can have the same apparent size, provided the more distant cluster has a larger mass that exactly compensates for the decrease in the apparent size due to the increased redshift. There is, therefore, a degeneracy between the cluster redshift and mass.  The question is whether we can break this degeneracy by using additional information. A resolved SZE image of a cluster provides information not only about the cluster size, but also about the shape of the cluster gas distribution. The total observed flux of the cluster, for instance, depends on the total cluster mass, the redshift and the temperature. The central SZE decrement depends on the core radius and the electron central density, but it is, in principle, independent of the redshift. Our method incorporates these and other observables to break the mass--redshift degeneracy.  This method requires resolved SZE images. Therefore, it should be useful for arcmin and sub-arcmin resolution experiments but not for experiments like {\\it Planck},  where the best resolution will be 5~arcmin.  In this work we will not consider the effects of the relativistic corrections and the kinematic effect, because they are small compared with the non-relativistic thermal SZE. Their effect will be discussed in a later paper.  In $\\S$\\ref{sec:clusters} we outline the connections between cluster morphology and redshift from a structure formation viewpoint. $\\S$\\ref{sec:method} discusses some of the weaknesses of this theoretical perspective and then provides a detailed description of a method that overcomes these weaknesses.  A demonstration of the degree to which the method works is contained in $\\S$\\ref{sec:application}, and a discussion of conclusions follows in $\\S$\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have developed a means of estimating galaxy cluster redshifts using only observed SZE properties of the clusters. Using a toy model we show how morphological quantities associated to clusters may contain redshift information. We also show how modelling of the morphological quantities can lead to systematic errors in the redshift estimation. We then propose an alternative method which is model independent. Specifically, we have combined several redshift sensitive SZE observables using a standard principal component analysis (PCA).  The PCA led to significant compression, showing that most of the redshift information contained in the 11 SZE observables can be expressed in three orthogonal linear combinations.  The use of the PCA has several advantages. These include (i) no required assumptions about cluster scaling relations, (ii) straightforward to use of direct observables (like the isophotal quantities), and (iii) orthogonality of the principal components. The method must be {\\it calibrated}, and we suggest using a cluster training set that has redshift estimates from photometric or spectroscopic means. This training set is required to build the likelihoods of the principal components as a function of redshift.  \\begin{figure} \\ifthenelse{\\equal{\\figtype}{EPS}}{ \\begin{center} \\epsfxsize=8.cm \\begin{minipage}{\\epsfxsize}\\epsffile{z_true_vs_z_recov_3PC_1PC.eps}\\end{minipage} \\end{center}} {\\myputfigure{f8.pdf}{3.25}{0.50}{-85}{-30}} \\caption{\\label{fig_z_true_vs_z_recov} Mean recovered redshift and error bar (dispersion) as a function of redshift. The solid line represents the ideal situation where the recovered redshift equals the true one. For comparisson we also show the correponding recovered redshifts when only the first principal component is used in the Bayesian approach (dotted error bars). The error bars for this case have been displacced 0.05 units in redshift to the right. } \\end{figure}  In our analysis we include 11 different {\\it observables}: isophotal flux, isophotal size, central amplitude, second derivative at the center, the mean of the first derivative in the region where the second derivative vanishes, the ellipticity and perimeter of the isophote, the number of subgroups above the isophote, and three Mexican--hat wavelet coefficients evaluated at the cluster center. Principle components were determined, and the first three components had $\\approx 90$ \\% of the variance of the data. Application of our redshift estimator using these three components indicates that the method can distinguish between clusters at low, intermediate and high redshift.  Although the error bar for a specific cluster redshift is fractionally large, our method should be useful for future SZE surveys, providing a preselection of low, intermediate and high redshift clusters. This preselection can be used to optimize the optical followup.  Because of the smoothly varying nature of the cluster redshift distribution expected in future surveys, it may also be possible to obtain cosmological constraints directly with these morphological redshifts. As shown in \\citet{fan01}, the ratio of the number of clusters above and below a given redshift can be a useful cosmological discriminator. This kind of analysis could be well suited to our morphological redshift estimates.   \\begin{figure} \\ifthenelse{\\equal{\\figtype}{EPS}}{ \\begin{center} \\epsfxsize=8.cm \\begin{minipage}{\\epsfxsize}\\epsffile{z_true_vs_z_recov_3PC_3vs8Observables.eps}\\end{minipage} \\end{center}} {\\myputfigure{f8.pdf}{3.25}{0.50}{-85}{-30}} \\caption{\\label{fig_z_true_vs_z_recov_3vs8Obs} Recovered vs true redshift in the case where only three observables: central amplitude, isophotal flux, and isophotal size are considered in the PCA analysis (solid line). This result is comparable with what the one obtained when the 11 observables are considerd but only the first PC was used in the redshift estimation. The dashed lines show the corresponding redshift estimation when the remaining 8 observables are considered in the analysis (first and second derivatives, 3 MHW coefficients, ellipticity, number of subgroups and perimeter) and the central amplitude, isophotal flux, and isophotal size are excluded. The true redshift has been displaced 0.05 to the right to avoid overlapping. Note the good constraints on $z$ obtained by these eight observables at low $z$. } \\end{figure}  A requirement for morphological redshifts is resolved, SZE cluster images.  Our estimates were carried out assuming an instrument resolution of 25~arcsec.  This resolution requirement makes our method inappropriate for application to clusters detected in the {\\it Planck} Surveyor mission, but there are several planned interferometric and single dish SZE surveys which could take advantage of our method.  Although in this work we have only considered the case of the SZE, our method can be extended to X-ray and optical cluster surveys. The main difference would be that the flux in the X-ray and optical bands are inversely proportional to the luminosity distance squared and the region of the spectrum observed by a particular instrument also varies with redshift. The difference between the luminosity distance and the angular diameter distance is a factor $(1 + z)^2$. In general, the X-ray and optical flux is much more sensitive to the cluster redshift than is the SZE flux. Although the redshift of galaxy clusters in X-rays can be obtained, for some clusters, directly from their X-ray spectrum (with typical errors of $\\Delta z \\approx 0.2$), for many clusters with a low SNR the redshift can not be obtained from this method.  Large, planned X--ray surveys will have a preponderance of low signal to noise detections, making the use of morphological redshifts (alone or combined with photometric redshifts) very promising."
    },
    "0207/astro-ph0207435_arXiv.txt": {
        "abstract": "\\noindent We develop a new nonlinear mean field dynamo theory that couples field growth to the time evolution of the magnetic helicity and the turbulent electromotive force, $\\emfb$. We show that the difference between kinetic and current helicities emerges naturally as the growth driver when the time derivative of $\\emfb$ is coupled into the theory. The solutions predict significant field growth in a kinematic phase and a saturation rate/strength that is magnetic Reynolds number dependent/independent in agreement with numerical simulations. The amplitude of early time oscillations provides a diagnostic for the closure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207603_arXiv.txt": {
        "abstract": "During the course of an investigation of the interaction of the radio galaxy M84 and its ambient cluster gas, we found excess X-ray emission aligned with the northern radio jet.  The emission extends from the X-ray core of the host galaxy as a weak bridge and then brightens to a local peak coincident with the first detectable radio knot at $\\approx~2.5^{\\prime\\prime}$ from the core.  The second radio knot at 3.3$^{\\prime\\prime}$ is brighter in both radio and X-rays.  The X-ray jet terminates 3.9$^{\\prime\\prime}$ from the core.  Although all the evidence suggests that Doppler favoritism augments the emission of the northern jet, it is unlikely that the excess X-ray emission is produced by inverse Compton emission.  We find many similarities between the M84 X-ray jet and recent jet detections from Chandra data of low luminosity radio galaxies.  For most of these current detections synchrotron emission is the favored explanation for the observed X-rays. ",
        "introduction": "The radio galaxy M84 is a low luminosity (FRI type) radio galaxy in the Virgo cluster.  We obtained Chandra observations in order to study the interaction of the radio structures with the hot intra cluster medium (ICM) and that work was reported in Finoguenov \\& Jones (2001). In this paper we report on X-ray emission detected from the inner 300 pc of the northern radio jet.  X-ray emission from radio jets presents us with the problem of identifying the emission process but once this process is determined, we can then obtain new constraints on physical parameters (Harris and Krawczynski, 2002).  With the introduction of the relativistic beaming model of Celotti (Celotti, Ghisellini, \\& Chiaberge, 2001) and Tavecchio (Tavecchio, et al. 2000), most X-ray emission from jets has been interpreted as indicating either synchrotron emission or inverse Compton scattering off the cosmic microwave background (CMB).  For M84, we show that synchrotron emission is the probable process, as has been found for a number of other FRI radio galaxies (Worrall, Birkinshaw, \\& Hardcastle, 2001; Hardcastle, Birkinshaw, \\& Worrall 2001).  The implications of the detected X-ray emission are discussed in sec.~\\ref{sec:disc}.   We take the distance to the Virgo cluster to be 17 Mpc so that one arcsec corresponds to 82 pc.  We follow the usual convention for flux density, S$\\propto~\\nu^{-\\alpha}$. ",
        "conclusions": "\\label{sec:disc}  We believe the evidence favors synchrotron emission for the observed X-rays although to sustain this model, higher s/n X-ray data and optical detections are required.  Undoubtedly there are bulk relativistic velocities in the jet producing the observed intensity differences between the N and S jets, but with velocity vectors not too far from the plane of the sky, we see only mild boosting and the parameters for IC/CMB emission are completely at odds with all other evidence.  To check the 'mild beaming' synchrotron model, we note that the ratio of the net counts in the N jet (r=1.05$^{\\prime\\prime}$ aperture) to that found in the same sized circle at the same distance to the south  is $>$4.1 (where we used the 2$\\sigma$ upper limit for the south value).  Thus the X-ray ratio (North/South) is consistent with the radio value (16) measured at the same distance from the core (3$^{\\prime\\prime}$).  Even in fields of order 100 $\\mu$G, the half-life for electrons producing X-rays is so short that they could travel no more than $\\approx$~3pc from their acceleration region.  Thus the X-rays clearly demarcate that sort of acceleration region.   The M84 jet is considerably weaker than other FRI detections.  The X-ray luminosity is $\\approx$~6 times less than that of Cen A, the jet with the lowest luminosity listed in a table of 7 FRI jets in Harris, Krawczynski, \\& Taylor (2002).  The combined flux densities of N2.5 and N3.3 are a factor of 84 less than that from knots HST-1 and D in the M87 jet (Marshall et al. 2002); these two knots are within 3$^{\\prime\\prime}$ of the nucleus of M87 although their physical distance from the core may be larger than that for the M84 knots owing to a larger projection effect.  Although some of the disparity in X-ray luminosity between M84 and the other FRI jets may be caused by differences in beaming factors, most of the current sample are believed to have rather large angles between the line of sight and the jet axis and thus beaming is generally moderate to low in all of them."
    },
    "0207/astro-ph0207329_arXiv.txt": {
        "abstract": "{New Integral Field Spectroscopy of the central region of NGC 7331 reveals strong H$\\alpha$ emission in the well-known CO and HI ring of NGC 7331. The [NII]/H$\\alpha$ ratio indicates that a large scale stellar formation process is taken place at the ring in agreement with previous hypothesis about the exhaustion of gas in the inner to the ring region. The dynamics of stars and gas are not coupled. There is a ring of peculiar velocities in the ionized gas velocity map. These peculiar velocities can be well interpreted by the presence of an axisymmetric inflow of 40 km/s at the inner boundary of the large-scale gaseous ring. We infer an inwards total flux of 1.6 M$_{\\odot}$yr$^{-1}$. This value is typical of the accretion rates in hypothetical {\\bf large} nuclear black holes. Despite the large differences in the scales of the nucleus and the gas ring of NGC 7331, we suggest that this inwards flux is feeding the nucleus. ",
        "introduction": "Peculiar motions are frequently detected in active galaxies where the existence of several kinematically distinct gaseous systems, some of them suffering radial outward movements, has been attributed to the influence of the active nucleus (Garc\\'{\\i}a-Lorenzo et al. 1999, 2001, Arribas et al. 1997, Mediavilla \\& Arribas 1993). The outflow of gas close to the nucleus is relatively easy to detect because of the large velocities involved (of the order of 100 km/s) and of the direct illumination from the strong active source. Inflow of gas seems to be much more difficult to observe as it is presumably caused by less outstanding mechanisms than the nuclear activity. However, it is very important to detect since inflow should provide the matter that will eventually accrete into the nuclear black hole. To study the influence of the environment on the nucleus we will present new results about the kinematics in the central region of NGC 7331. Some peculiarities previously detected in the ionized gas velocity field of this galaxy (see below) make it specially interesting for this kind of studies.  The large scale gas distribution in NGC 7331 is ring-like as first detected by Bosma (1978) in HI, Telesco et al. (1982) from NIR colors and Young and Scoville (1982) in CO. The ring is confirmed in other wavelengths and is taken today as a prototype gas ring.  Bower et al. (1993) used optical long-slit spectra to show that their models without a central black hole fit the observational data. However, there is an unresolved LINER nucleus (see, for instance, Cowan et al., 1994), so this galaxy could harbour a massive black hole of $5\\times10^8 M_{\\odot}$. This is supported by the motion of ionized gas ([NII] + H$\\alpha$ emission line) reported by Afanasiev et al. (1989) in the 0.2-0.4 kpc (2.8-5.6 arcsec) zone and by the discovery of a nuclear X-ray source by Stockdale et al. (1998) using ROSAT.  The existence of a bar has been suggested from non-circular motions by Marcelin et al. (1994) and von Linden et al. (1996). However the I and K band photometry in the central regions carried out by Prada et al. (1996) did not indicate any significant barred morphology.  In the central kpc, Bottema (1999) found that the emission line gas (for R $\\leq$ 40 arcsec) seems to rotate slower than the stars. An explanation of the observations would consist in an inclined and warped gaseous plane.  Using 2D spectroscopy, Mediavilla et al. (1997) concluded that the kinematics of the stars and most of the ionized gas are decoupled. These authors found that the kinematic axes of the ionized gas velocity map are distorted and rotated with respect to the stars, something that can be interpreted in terms of radial movements. However, the reduced field of view of the available velocity maps (about 7'' $\\times$  7'') makes difficult to verify this interpretation.  In this letter we are going to present new Integral Field spectroscopy covering the central 30'' $\\times$ 30'' of NGC 7331, with the aim of understanding the kinematics of the region surrounding the nucleus. ",
        "conclusions": "Some conclusions are worthy to be summarized:  We have detected an H$\\alpha$ ring coincident with the inner part of the well known CO and HI ring. The [NII]/H$\\alpha$ ratio at the ring is typical of HII-like ionization and indicates that a large scale stellar formation process is taken place at the ring. This is in agreement with the hypothesis of that the gas in the inner to the ring region has been exhausted by massive stellar formation, that continues at the ring. The physical conditions in the inner to the ring region are typical of LINERS and the transition between starburst and LINER takes place very sharply at the inner boundary of the H$\\alpha$ ring.  Ionized gas peculiar velocities are found near to the inner boundary of the CO-HI ring, compatible with an axisymmetric morphology superimposed to the otherwise regular map of the ionized gas. The comparison with the regular stellar velocity map makes the assumption of an axisymmetric contraction very plausible. The infalling matter is observed at the inner boundary of the large-scale gaseous ring and likely continues towards the inner region. Although several mechanisms are possible to explain the origin of the inflow, the existence of strong winds at the ring wall fits naturally with the presence of massive stellar formation in the ring.  Even if the accreting flux could have very different causes we would be observing infalling velocities of the order of 50 km/s and infalling gaseous fluxes of the order of 1 M$_{\\odot}$yr$^{-1}$.  This is higher than the required accretion rate for a central black hole. Probably we are observing the infalling motion feeding the black hole having its source as far as 2 kpc away from the center.  The inwards flux axisymmetry seems to rule out, in this galaxy, mechanisms based on bars. Bars have been proposed as sources for feeding AGN's and this have been observed to operate in the central regions of some active galaxies (See Perez et al. 2000, and references therein). NGC 7331, without a clearly observed bar, and with an axisymmetric inflow suggest that other ways of loosing angular momentum could also be important."
    },
    "0207/astro-ph0207573_arXiv.txt": {
        "abstract": "{ We suggest that magnetic fields in the accretion disks of AGN reach into the coronae above and have a profound effect on the mass flow rate in the corona. This strongly affects the location where the accretion flow changes from a geometrically thin disk to a pure vertically extended coronal or advection-dominated accretion flow (ADAF). We show that this can explain the different disk truncation radii in elliptical galaxies and low luminosity AGN with about the same mass flow rate, a discrepancy pointed out by Quataert et al. (1999). Without disk magnetic activity the disk truncation is expected to be uniquely related to the mass flow rate (Meyer et al. 2000b). Whether dynamo action occurs depends on whether the electrical conductivity measured by a magnetic Reynolds number surpasses a critical value (Gammie \\& Menou 1998). In elliptical galaxies the disk is self-gravitating at the radii where the truncation should occur. It is plausible that instead of a cool disk a ``layer of clouds'' may form (Shlosman et al. 1990, Gammie 2001) for which no dynamo action is expected. For low luminosity AGN the magnetic Reynolds number is well above critical. Simple model calculations show that magnetic fields in the underlying disks reduce the strength of the coronal flow and shift the truncation radius significantly inward. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207059_arXiv.txt": {
        "abstract": "Time series of {\\it B,V,I} CCD photometry and radial velocity  measurements from high resolution  spectroscopy (R=30,000) covering the full pulsation cycle are presented for the  field RR Lyrae star CM Leonis. The photometric data span a 6 year interval  from 1994 to 1999,  and allow us to firmly establish the pulsation mode and periodicity of the  variable.   The derived period {\\it P}=0.361699 days  ($\\pm 0.000001$) is very close to the value published   in the Fourth Edition of the General Catalogue of Variable Stars ({\\it P}=0.361732 days). However, contrary to what was previously found,  the amplitude and shape of the light curve qualify CM Leo as a very regular first overtone pulsator  with a prominent  hump on the rising branch of its multicolour light curves. According to an  abundace analysis performed on three spectra taken near minimum light (0.42 $< \\phi <$ 0.61), CM Leo is  a metal-poor star  with metal abundance [Fe/H]=$-1.93 \\pm 0.20$. The photometric and radial  velocity curves of CM Leo have been compared with the predictions of suitable pulsational  models to infer tight  constraints on the  stellar mass, effective temperature, and distance modulus of the star.  We derive a true distance modulus of CM Leo of $\\mu_0$=13.11 $\\pm$0.02 mag and a  corresponding absolute magnitude of $M_V$=0.47$\\pm$0.04.  This absolute magnitude, once corrected for  evolutionary and metallicity  effects, leads to a true distance modulus of the Large Magellanic Cloud of  $\\mu_0$=18.43$\\pm$ 0.06 mag, in better agreement with the {\\it long} astronomical distance scale. ",
        "introduction": "This is the third in a series of papers dealing with the detailed study of  the pulsational  properties of a sample of field RR Lyrae stars that were known to exhibit anomalous  scatter and variations in  the shape and amplitude of their light curves. In Clementini et al. (1995b; hereinafter Paper I) we  reported results  from a first photometric study of a sample  of 8 field RR Lyrae stars which, according to the literature, were classified  as fundamental mode pulsators (RRab) with short periods and located at large  heights {\\it z } from the Galactic plane (Castellani, Maceroni \\& Tosi 1983).  The circumstance of being far from the disc but with a period that for RRab usually corresponds to roughly solar metallicity made these stars fairly  anomalous. The photometry presented in Paper I demonstrated however that  five out of the eight variables were more likely to pulsate in the first  overtone mode, and that, therefore, the short period should not be taken  as indicative of high metallicity. Furthermore, it  showed that some of the stars in the sample (namely CM Leonis, CU, BS  and BE Comae) exhibited irregularities in their light curves. In order  to investigate  the actual nature of these anomalies new observing campaigns were conducted  on these four  stars. In Clementini et al. (2000, Paper II) we  reported the discovery that CU Comae is a {\\it double-mode} RR Lyrae, the sixth detected in the field of our Galaxy and the most metal-poor ever found. The present paper is  devoted to  the discussion of the results obtained for CM Leonis (CM Leo) from the  combination of its Paper I photometry with new data taken from 1995 to 1999,  and from   high resolution spectroscopy obtained in 1999. The observations and data sets are presented in Section 2. In Section 3 we discuss the results of  the analysis of the complete photometric data set. In Section 4 we report the results  of our spectroscopic analyses (derivation of a radial  velocity curve over the full pulsation cycle and  elemental abundance analysis of the spectra of CM Leo  taken near minimum light).  In section 5, we present results of the modeling of  the observed light and radial velocity curves of CM Leo,  based on nonlinear pulsation models, computed with the same physical and numerical assumptions as in Bono, Castellani \\& Marconi (2000), and derive an estimate of  the stellar mass, effective temperature, and distance modulus of the star. Finally, in Section 6 we summarize the main derived quantities of CM Leo and  discuss the absolute magnitude we derive for the star in the framework of the  {\\it short} and {\\it long} distance scale dichotomy. ",
        "conclusions": "CM Leo is a very regular first overtone RR Lyrae (RRc) with  period P=0.$^d$361699, epoch E=2450841.$^d$362285,  a prominent hump  on the raising branch of its light curves, and metal abundance [Fe/H]=$-$1.93 $\\pm 0.20$. Final results  as well as the average quantities derived for the star  from our photometric, pulsational and spectroscopic analyses  are summarized in Table~8.  We have shown that the irregularities found by Clementini et al. (1995b) in the  light curve of CM Leo disappear when star C3 is used as comparison star, since they were in fact produced by variations of about 0.1 mag  in the light of the reference star used in Paper I (star C1). The true distance modulus   $\\mu_0=13.11 \\pm$ 0.02 mag we derive for CM Leo leads to  a distance from the Galactic  disc\\footnote{The distance from the galactic plane almost coincides  with the distance from the Sun since CM Leo is located in the  direction of the Galactic pole.} of $d=4.2$ kpc. This distance is fully   consistent with the rather low metallicity we found for the star, thus removing  the previously reported metallicity-distance anomaly of CM Leo. We recall that a similar  result was found also for CU Com (see Paper II).  RR Lyrae stars are known to be powerful standard candles through their absolute magnitude that is only slightly dependent on metal abundance. A {\\it pulsational} estimate  of the absolute magnitude of an RR Lyrae star provides an independent  calibration of these standard candles and can thus be used  to infer the  distance of any other stellar system where RR Lyrae's are found, the  Large Magellanic Cloud in particular. The absolute magnitude we derive for CM Leo  from the  modeling of the multicolour light and radial velocity  curves is M$_V$=0.47$\\pm$0.04 mag, where the error includes the random uncertainty contributions of the photometry (0.03 mag), and of the fitting with the theoretical pulsational models (0.02 mag in distance modulus and 0.01 mag in $\\log{L}$, respectively).   \\begin{table} \\begin{center} \\caption{Properties of CM Leo} \\vspace{5 mm} \\begin{tabular}{c c} \\hline \\hline Type  & RRc \\\\ ${\\rm [Fe/H]}$& $-1.93 \\pm 0.20$\\\\ Epoch & 2450841.$^d$362285 \\\\ P & 0.$^d$361699$\\pm$ 0.000001 \\\\ $< V >$ & 13.66\\\\ $< B >$ & 13.97\\\\ $< I >$ & 13.32\\\\ $< B > - < V >$ & 0.31\\\\ $< V > - < I >$ & 0.34\\\\ $A_V$ & 0.498\\\\ $A_B$ & 0.631\\\\ $A_I$ & 0.301\\\\ $A_{\\rm RV}$ & ~~26.55 km s$^{-1}$\\\\ $\\gamma$ & $-$24.47 km s$^{-1}$\\\\ M$_V$ & 0.47\\\\ $\\mu_0(V)$ & 13.11\\\\ \\hline \\end{tabular} \\end{center} \\label{t:tab6} \\end{table}  Allowing for an evolution  of about 0.08 mag off the Zero Age  Horizontal Branch for the RR Lyrae stars, as suggested by  Caputo \\& Degl'Innocenti (1995), and Caloi, D'Antona  \\& Mazzitelli (1997), and  a slope of 0.2 mag/dex for the luminosity-metallicity relation of  RR Lyrae's to correct CM Leo absolute magnitude to the   average metal abundance of the RR Lyrae in the Large Magellanic Cloud  ([Fe/H]=$-$1.50: Alcock et al. 1996, Bragaglia et al. 2001a),  the absolute magnitude we derive for CM Leo implies a value for  the true distance modulus of the LMC of  $\\mu_0$(LMC) = 18.43$\\pm$0.06 mag, for an average dereddened luminosity of the RR Lyrae in the LMC bar $<V_0(RR)>$=19.07$\\pm$0.05 (Clementini et al. 2001),  in better agreement with the {\\it long} astronomical distance scale.   \\bigskip\\noindent  ACKNOWLEDGEMENTS  This paper is based on observations  obtained with the 1.52 m telescope of the Bologna Observatory in Loiano, the Michigan State University 60 cm telescope, and the 2.7 m telescope of the  McDonald Observatory. This work was partially funded by MURST-Cofin00 under the project ``Stellar Observables of Cosmological Relevance\". SDT was partially supported with an Italian CNAA fellowship at the University of Texas at Austin. She thanks all the Texas Staff for the warm hospitality. III gratefully acknowledges support from Continuing Fellowships at the University of Texas at Austin. HAS thanks C. Wilkinson for assistance in obtaining the 1998 MSU observations and thanks the US National Science Foundation for support under grant AST9986943. CS is happy to acknowledge that this research was partially funded by NSF grants, most recently AST9987162.  We warmly thank Raffaele Gratton for interesting suggestions and support."
    },
    "0207/astro-ph0207084_arXiv.txt": {
        "abstract": "{We show that the absence of outbursts during low states of VY Scl stars is easily explained if white dwarfs in these systems are weakly magnetized ($\\mu\\gta 5\\times 10^{30}$ G cm$^{3}$). However, some of the VY Scl stars are observed to have very slow declines to minimum and similarly slow rises to maximum. The absence of outbursts during such {\\sl intermediate} (as opposed to {\\sl low}) states, which last much longer than typical disc's viscous times, can be explained only if accretion discs are absent when their temperatures would correspond to an unstable state. This requires magnetic fields stronger than those explaining outburst absence during low states, since white dwarfs in this sub-class of VY Scl stars should have magnetic moments $\\mu\\gta 1.5\\times 10^{33}$ G cm$^{3}$ i.e. similar to those of Intermediate Polars. Since at maximum brightness several VY Scl stars are SW Sex stars, this conclusion is in agreement with recent claims about the magnetic nature of these systems. ",
        "introduction": "Cataclysmic variables (CVs) are close binary systems in which a white dwarf primary star accretes matter lost from a Roche-lobe filling, low-mass secondary companion. In most CVs, matter transferred from the secondary forms an accretion disc around the white dwarf. Orbital periods of CVs extend from a sharp minimum at $\\sim 80$ min to $\\sim 10$ hrs (a few systems with giant or sub-giant secondaries have longer periods), with a prominent deficit of systems between $\\sim 2$ and $\\sim 3$ hours (the ``period gap''). CVs are divided into different types according to their observed properties. Some of these properties reflect the way the systems are seen (e.g. their inclination) but most are a reflection of the physical parameters of the system. Here we will be mainly interested in two broad types of CVs which are classified according to the value of the rate at which matter is transferred from the secondary. These are the Dwarf Novae (DNs) and Nova-Like (NLs) variables -- a nomenclature reflecting misunderstandings of historical interest only (see Warner \\cite{w95} for details). DNs show 2-5 mag (up to 8 mag in some cases) outbursts lasting from a couple of days up to more than a month. Outbursts are separated by quiescent intervals lasting from $\\sim 10$ d to tens of years. It is believed that DN outbursts are due to a thermal-viscous accretion-disc instability occurring at temperatures where hydrogen recombines. According to the disc instability model (DIM) which is currently used to describe DN outburst cycle (see Lasota \\cite{l01} for a review), the condition for the instability to occur can be expressed in terms of a critical rate at which matter is fed to the outer disc's regions. Above this critical rate (whose value strongly increases with the disc's radius) accretion discs are hot and stable. CVs with such discs belong to the class of Nova-Like variable, which are defined as `non-eruptive' CVs. There exists also a second critical accretion rate below which a CV disc is cold and stable. The two critical rates have the same radial dependence.  The membership of a class is not permanent but depends on the actual mass-transfer rate. It seems, for example, that a NL should become a DN if a fluctuation brings the mass-transfer rate below the critical value for instability (the opposite fluctuation would transform a DN into a NL). The behaviour of VY Scl stars contradicts, however, this apparently reasonable conclusion. These are very bright NLs which occasionally undergo a fall in brightness by more than one magnitude. Although such drops in luminosity bring them into the DN instability strip, they show no outbursts. In fact, during their low states, which may last weeks or several years, VY Scl stars may spectroscopically appear like quiescent DNs but they {\\sl do not} become `regular' DNs (see Warner \\cite{w95} and references therein). At best they show once or twice a DN-type outburst but not an outburst cycle one could expect from systems at mass-transfer rates lower than the critical one for the disc instability to occur.  It would be tempting to say that VY Scl stars during low states are stable with respect to the DN-type instability, i.e. the effective temperature of their disc is below the instability strip. Since they are also stable in their high  states, they would just oscillate between two stable steady-state configurations, explaining the general absence of outbursts. However, this simple explanation cannot be correct because during the transition from one stable state to the other, the disc configuration (its accretion rate) must cross the dwarf-nova instability strip and therefore show outbursts.  \\subsection{Mass transfer fluctuations}  Before trying to explain this strange behaviour, one should determine the cause of brightness fluctuations in these systems. They seem to be due to mass-transfer rate fluctuations. Indeed, there is compelling evidence for mass-transfer fluctuations of various amplitudes occurring on various time-scales. Of course, in systems with an accretion disc, it might be difficult to make the difference between variations in the {\\sl accretion} and the {\\sl mass-transfer} rates, unless one can observe the mass-transfer-stream impact region in the outer disc (the so called `bright' or `hot' spot) whose changes of brightness would reflect mass-transfer variations. However, in polars (AM Her stars), a class of strongly magnetized CVs, no accretion disc is present so that observed changes in brightness by 1.5 do 4.5 mag must result directly from mass-transfer variations (see Hessman et al. \\cite{hgm00} and Buat-M\\'enard et al. \\cite{bhl01b} for an evaluation of the effect in DNs of mass transfer variations similar to those observed in AM Her systems). Also in VY Scl stars it is obvious that the mass-transfer rate fluctuates between a very high value (comparable or higher than the {\\sl accretion rate} at DN's maximum; Warner \\cite{w95}) and a value so low that the mass-transfer rate can be considered to have stopped.  Let us first consider what is happening when the mass-transfer rate is rapidly switched off. In such a case, as nicely explained by Leach et al. (\\cite{lhk99}, hereafter L99), the moment the mass-transfer rate drops below the (upper) critical value the outer disc becomes unstable and a cooling front propagates inwards, bringing the whole disc into a cold state. This state, however, is {\\sl non-stationary} as the accretion rate in the disc is not constant but increases with radius (see e.g. Figs. 10 and 17 in Lasota \\cite{l01}). Since in such a disc most of the mass is in its outer regions, matter will diffuse inwards and inevitably cross the critical line as in Figs. 10 and 17 of Lasota (\\cite{l01}) and the system will go into outburst.  L99 suggested that in VY Scl irradiation by white dwarfs heated by accretion to very high temperatures ($\\gta 40,000$ K) during high states, stabilizes the inner disc thus preventing outbursts. Indeed, as noticed by Lasota et al. (\\cite{lhh95}) a cold, dwarf-nova-type accretion disc (i.e. a disc whose structure results from the passage of a cooling front) can be stable if its inner part is removed. In L99 it is assumed that the inner disc is kept hot by white-dwarf irradiation which, as far as the cooling-front propagation is concerned, has the same effect a disc truncation: the cooling front has to stop at the transition radius between the cold and the hot part of the disc. The transition radius has to be large enough to allow the cold disc to settle to a quasi-stationary stable configuration (see L99 for details). As we will see in Section 2 the L99 model cannot work for white-dwarf masses larger than $\\sim 0.4$ \\msol, which makes its application to {\\sl all} VY Scl stars rather questionable. In Section 3 we will show, however, that inner disc truncation by magnetic fields of not extravagantly large values (magnetic moments $\\mu\\gta 5\\times 10^{30}$ G cm$^{3}$, i.e. magnetic fields $\\gta 40$ kG for a $\\sim 0.7$ \\msol white dwarf) also suppress outbursts in VY Scl low states and does not require unrealistic parameters. The proposal that VY Scl systems may contain magnetized white dwarfs is not new: using arguments based on the observed spectral or timing properties of these systems, Jameson et al. (\\cite{jsk82}), and Hollander \\& van Paradijs (\\cite{hv92}) proposed that TT Ari might be magnetized, while Voikhanskaia \\cite{v88} suggested that accretion in MV Lyr occurs on a magnetic pole during its low states.  In Section 4 we point out that although such a truncation suppresses outbursts during the minimum light, it cannot explain the absence of outbursts during the very slow decays from maximum to minimum light and very slow rises back to maximum. For example, Honeycutt et al. (\\cite{hlr93}) observed a $\\sim 500$ day, 4 mag rise of \\object{DW UMa}, which is very long compared to the disc viscous time expected to be rather 10 - 20 days. It is instructive to have a closer look at Fig. 4 of L99. It shows part of the \\object{MV Lyr} light-curve and the light-curve calculated by L99. Their model nicely fits the first rapid drop in luminosity but fails to reproduce the following slow rise and had these authors continued their calculation it would miserably fail to account for next, very slow (50 days) decay (see our Fig. \\ref{fig:slowvar}), contrary to their expectations.  The reason is simple: when the mass-transfer rate remains in the instability strip longer than the viscous time, the system must produce a series of outbursts. We will show that in order to suppress outbursts during very slow rises and decays one has to suppress the disc as soon as the mass-transfer rate enters the instability range. This, as we argue, requires magnetic moments similar to those of Intermediate Polars,  In Section 5 we recall the independent evidence for the magnetic nature of VY Scl stars. In Section 6, after discussing the possibility that in some VY Scl stars the magnetized white dwarf could be in the propeller regime (we modify the standard inner boundary condition of the model in order to account for a source of angular momentum and matter ejection), we decide this would not change our main conclusion about the magnetic moment required to stabilize intermediate luminosity states. In Section 7 we discuss the relation between VY Scl stars and other types of CVs. Section 8 summarizes our results. ",
        "conclusions": "We found that irradiation by a hot white dwarf cannot account for the whole VY Scl syndrome. Even if these systems had the required (low) white dwarf masses and very high white dwarf temperature to prevent dwarf-nova type outbursts during their low states, the absence of outbursts during long intermediate states observed in some VY Scl stars cannot be accounted for by white dwarf irradiation. Therefore in such binaries the white dwarf has to possess a magnetic moment similar to that of IPs. This would also prevent outbursts in the low states. On the other hand in VY Scl systems in which fluctuations of the mass-transfer rate happen on timescales shorter than the disc's viscous time and/or outbursts occur during the long transitions, the stabilizing effect could be produced by a much weaker magnetic field.  Of course we have not {\\sl proven} that VY Scl stars must be magnetic. In principle one could imagine an ``evaporation mechanism\" so tuned that it would get rid of the disc in the right moment. Our model, however, provides the most conservative solution of the VY Sc puzzle. In addition there exist independent indication of the magnetic nature of least some bright nova-like binaries.  Timing and polarization observations of VY Scl stars will provide the ultimate test of our model."
    },
    "0207/astro-ph0207421_arXiv.txt": {
        "abstract": "We explore the origin of the flat spectrum seen in some T~Tauri stars by considering a three-component structure: a central star, a circumstellar disk, and a dusty halo. The radiative energy transport is faithfully treated by solving the angle- and frequency-dependent radiative transfer equation in two space dimensions assuming axisymmetry, and hence the radiative equilibrium temperature in the disk and halo is determined simultaneously. The disk is effectively heated by the scattering and reprocessing of stellar radiation through the halo. The large mid- to far-infrared excess originates from the photosphere of the warmed disk, resulting in a flat spectrum, as observed. The halo which we consider is observed as a compact reflection nebula, and is discriminated from extended, disk-like envelopes around flat-spectrum T~Tauri stars. We show that the overall spectral shape of flat-spectrum T~Tauri stars can be reproduced by the present {\\it disk--halo} model. ",
        "introduction": "It seems reasonable to suppose that the infrared excesses of T~Tauri stars can be attributed to thermal dust emission from circumstellar disks. By modeling the observed spectral energy distributions (SEDs), one can derive the disk properties, such as masses, radii, and temperature distributions (Adams et al.\\ 1988; Strom et al.\\ 1989; Beckwith et al.\\ 1990). Such spectral modeling, however, has revealed an extreme class of T~Tauri stars, namely flat-spectrum T~Tauri stars. The large mid- to far-infrared excesses of flat-spectrum T~Tauri stars require their disks to have a temperature distribution of the form $T \\propto R^{-1/2}$ (Adams et al.\\ 1988), where $R$ is the distance from the rotation axis, whereas a $T \\propto R^{-3/4}$ dependence is predicted by both a standard accretion disk (Lynden-Bell, Pringle 1974) and a spatially flat reprocessing disk (Adams, Shu 1986). This implies that the disks of flat-spectrum T~Tauri stars are warmer in the outer regions than predicted by the simple disk models; we must therefore consider the mechanisms which heat the outer region of the disk. Disk flaring allows a disk to receive more emission from the central star, and to produce temperature distributions shallower than $T \\propto R^{-3/4}$ (Kusaka et al.\\ 1970; Kenyon, Hartmann 1987), but does not suffice to reproduce the observed flux.   The infrared excesses may originate from another circumstellar dust component. Calvet et al.\\ (1994) invoked infalling envelopes, which were originally applied to the spectral modeling of protostars (Adams, Shu 1986; Kenyon et al. 1993), and showed that the infalling envelopes can produce the mid- to far-infrared excesses of flat-spectrum T~Tauri stars. In fact, observations have shown that there exists an extended, disk-like structure of radius $\\sim 1000\\ {\\rm AU}$ around a typical flat-spectrum T~Tauri star, HL~Tau (Sargent, Beckwith 1991). Furthermore, Hayashi, Ohashi, and Miyama (1993) have revealed an infalling motion in the disk-like structure of HL~Tau, suggesting that it is a remnant of an infalling envelope.   Although the infalling envelope model of Calvet et al.\\ (1994) is successful in reproducing the flat spectrum, it does not take into account the following two important effects. First, as pointed out by Natta (1993), envelopes scatter and reprocess the stellar radiation toward the disk, and thereby alter the temperature distribution in the disk (see also Butner et al.\\ 1994; D'Alessio et al.\\ 1997). Second, the disk, itself, significantly influences the temperature distribution in the envelope. Since these two effects are coupled with each other, they cannot be treated separately. Hence, to elucidate the substantial mechanism for the flat spectrum, the temperature structure of the disk and envelope should be solved simultaneously. In the present analysis, we consistently handle the radiative energy transport in the disk and envelope by solving the angle- and frequency-dependent radiative transfer equation in two space dimensions by assuming axisymmetry.   In this paper, we demonstrate that the mid- to far-infrared excesses of flat-spectrum T~Tauri stars can originate from the {\\it disk}. Since the envelope which heats the disk can be as compact as the disk itself, we henceforth call it a {\\it halo} to distinguish it from an extended infalling envelope. The reflection nebula of HL~Tau revealed by the Hubble Space Telescope (Stapelfeldt et al.\\ 1995) may be regarded as an observational counterpart of the halo. ",
        "conclusions": "\\begin{enumerate} \\item We have shown that disks heated by the scattering and reprocessing of the stellar radiation through the halo can have flat infrared spectra.  \\item Local viscous heating is not sufficient to produce large mid- to far-infrared emission from the disk if we consider a reasonable rate of mass accretion in disks around classical T~Tauri stars. However, the accretion energy released in the innermost region of the star/disk system is transported by radiation through the halo to heat the outer region of the disk, resulting again in a flat spectrum.  \\item We examined the sensitivity of the SED to the assumed halo structure, and found that density distributions with a power-law index $\\le 3/2$ can provide the backwarming needed for flat infrared spectra. However, we have found that it is difficult to constrain the halo structure only from the SED in the infrared.  \\item We have discussed that the halo will be observed as a reflection nebula often associated with a T~Tauri star, indicating that it is the inner part of a remnant of an infalling envelope. A more detailed test of the model will be made by comparisons with imaging observations of near-infrared scattered light. \\end{enumerate}   \\vspace{1pc}  We are grateful to M.\\ Hayashi and M.\\ Umemura for valuable discussions. We also thank an anonymous referee for suggestions which improved the paper. The computations were performed on CP-PACS at the Center for Computational Physics in University of Tsukuba, and on the Fujitsu VPP300/16R at the Astronomical Data Analysis Center of the National Astronomical Observatory, Japan. TN was partially supported by the Grant-in-Aid for Scientific Research on Priority Areas (10147105) and for Scientific Research (10740093) of the Ministry of Education, Culture, Sports, Science, and Technology, Japan."
    },
    "0207/astro-ph0207171_arXiv.txt": {
        "abstract": "We present new  X-ray and radio observations of the Wolf-Rayet star EZ CMa (HD 50896) obtained with {\\em XMM-Newton} and the {\\em VLA}. This WN4 star exhibits optical and UV variability at a period of 3.765 d whose cause is unknown. Binarity may be responsible but the existence of a companion has not been proven. The radio spectral energy distribution of EZ CMa determined from {\\em VLA} observations at five frequencies is in excellent agreement with predictions for free-free wind emission and the ionized mass-loss rate allowing for distance uncertainties is $\\mathrm{\\dot{M}}$ = 3.8 ($\\pm$2.6) $\\times$ 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$. The CCD X-ray spectra show prominent Si XIII and S XV emission lines and can be acceptably modeled as an absorbed multi-temperature optically thin plasma, confirming earlier {\\em ASCA} results. Nonsolar abundances are inferred with Fe notably deficient. The X-ray emission is dominated by cooler plasma at a temperature  kT$_{cool}$ $\\approx$ 0.6 keV, but a harder component is also detected and the derived temperature is kT$_{hot}$ $\\approx$ 3.0 - 4.2 keV if the emission is thermal. This is too high to be explained by radiative wind shock models and the X-ray luminosity of the hard component is three orders of magnitude lower than expected for accretion onto a neutron star companion. We show that the hard emission could be produced by the Wolf-Rayet wind shocking onto a normal (nondegenerate) stellar companion at close separation. Finally, using comparable data sets we demonstrate that the X-ray and radio properties of EZ CMa are strikingly similar to those of the WN5-6 star WR110. This similarity points to common X-ray and radio emission processes in WN stars and discredits the idea that EZ CMa is anomalous within its class. ",
        "introduction": "Theories of the origin of X-ray emission in massive stars are now being reexamined  in light of new discoveries by the {\\em XMM-Newton} and {\\em Chandra} observatories. Traditionally, the X-ray emission of {\\em single} OB and Wolf-Rayet (W-R) stars without companions has been attributed to shocks distributed throughout their winds that form as a result of line-driven flow instabilities (Lucy \\& White 1980; Lucy 1982; Baum et al. 1992; Gayley \\& Owocki 1995; Feldmeier et al. 1997; Owocki, Castor, \\& Rybicki 1988). Such emission is predicted to be relatively soft (kT $<$ 1 keV) and X-ray emission lines  formed in an optically thick outflowing wind are expected to be blueshifted and asymmetric due to higher attenuation of the redward portion of the line by receding material on the far side of the star (MacFarlane et al. 1991; Owocki \\& Cohen 2001).  Recently obtained X-ray grating spectra of several OB  stars reveal properties that  conflict with the predictions of traditional radiative wind shock theory. In particular, the emission lines of the O9 supergiant $\\zeta$ Ori (Waldron \\& Cassinelli 2000) and the B0V star $\\tau$ Sco (Cohen et al. 2002) are unshifted and show no obvious asymmetries. The lack of asymmetries is difficult to explain if the lines are indeed formed in shocked winds (Owocki \\& Cohen 2001).  Also puzzling is the presence of narrow high-temperature Fe XXIII and Fe XXIV lines in $\\tau$ Sco indicative of hot plasma well in excess of $\\sim$1 keV. In contrast, much-needed support for the radiative wind shock paradigm has been received from {\\em Chandra} and {\\em XMM-Newton} grating observations of the O4 supergiant $\\zeta$ Puppis, whose emission lines are blueshifted and asymmetric (Cassinelli et al. 2001; Kahn et al. 2001).  Although grating observations are providing stringent tests of wind shock theories in OB stars, the observational picture for  W-R stars is much less complete due to the lack of suitable targets bright enough for X-ray grating observations. At current sensitivity levels, grating spectra of sufficient quality for emission line analysis can only be obtained for a few of the brightest W-R $+$ OB binary systems such as $\\gamma^{2}$ Velorum (Skinner et al. 2001; Dumm et al. 2002). Such binaries provide crucial information on extrastellar emission from colliding wind shocks between the stars (Cherepashchuk 1976; Prilutskii \\& Usov 1976; Usov 1992), but this emission is difficult to decouple from any intrinsic stellar emission that may be present since the binary components are usually not spatially resolved in X-rays.  We are thus undertaking an observing program using {\\em XMM-Newton} aimed at acquiring moderate resolution CCD spectra of fainter W-R stars that are currently beyond the reach of gratings. These include high-interest objects such as the putatively single WN5-6 star WR 110 (Skinner et al. 2002, hereafter SZGS02) and possible binary systems such as EZ CMa discussed here. The large effective area of {\\em XMM-Newton} is particularly well-suited for obtaining good-quality CCD spectra of such fainter objects in reasonable exposure times. Although CCD spectra do not provide the detailed information on line widths and profiles that can only be obtained with gratings, they are capable of discerning stronger emission lines and provide meaningful constraints on X-ray absorption and the overall distribution of plasma with temperature. Such information is more than adequate to discriminate between the relatively cool emission (kT $<$ 1 keV) that is expected from filamentary shocks distributed throughout the wind and harder emission at several keV that is predicted for colliding wind  binaries. As such, CCD spectra provide an ideal means for identifying candidate colliding wind binaries that may be amenable to study at higher spectral resolution with next-generation X-ray telescopes.  We present here new {\\em XMM-Newton} and  {\\em VLA} observations of the nitrogen-type W-R star EZ CMa (= HD 50896 = WR6). This WN4 star has been extensively studied at all wavelengths and shows optical and ultraviolet variability at a well-documented 3.765 day optical period, the origin of which is still not understood. Various explanations have been proposed including an as yet undetected  companion (Firmani et al. 1980; Lamontagne, Moffat \\& Lamarre 1986; Georgiev et al. 1999) and a rotationally modulated wind  (St.-Louis et al. 1995).  The {\\em XMM-Newton} observations provide broader energy coverage and higher signal-to-noise (S/N) spectra than in previous observations, giving more accurate measurements  of the X-ray absorption and plasma temperature distribution for comparison with emission models. The {\\em VLA} data yield the first reliable determination of the radio spectral index based on single-epoch multifrequency data. \\\\ ",
        "conclusions": "The new observational results discussed above provide the most detailed picture to date of the X-ray and centimeter radio properties of EZ CMa. Below, we make comparisons with previous studies and comment on specific emission models.  \\subsection{WN Stars: Comparative Spectroscopy}  The {\\em XMM-Newton} and VLA data  for EZ CMa provide a good basis for comparison with similar data recently obtained for WR 110 (SZGS02). For the first time we are able to directly compare the X-ray and radio properties of two WN stars having similar spectral types using analogous data sets.  Table 4 summarizes the X-ray and radio properties of EZ CMa and WR 110, and their X-ray spectra are compared in Figure 8. Overall, the two stars are strikingly similar and the same physical processes are very likely responsible for the X-rays and radio emission in both stars. One notable difference is the larger N$_{H}$ for WR 110, attributable to its stronger interstellar absorption toward the Galactic center (SZGS02).  The uncertain distance for EZ CMa introduces some ambiguity into the calculation of $\\mathrm{\\dot{M}}$ and L$_{x}$. But, as Table 4 shows, if an intermediate distance of 1.2 kpc is adopted for EZ CMa  then $\\mathrm{\\dot{M}}$ and L$_{x}$ agree with WR 110 to better than a factor of two.    \\subsection{Radiative Wind Shocks}  The {\\em XMM-Newton} spectra provide some constraints on wind shock models. The dominant cool emission component peaking near kT$_{cool}$ $\\approx$ 0.6 keV can potentially be explained by radiative wind shock models, but the hot component requires a different explanation.  Assuming that the cool component is due to shocks distributed throughout the wind, then the observed X-ray temperature provides constraints on the shock speed v$_{s}$ using the adiabatic shock formula kT$_{s}$ = (3/16)\\={m}v$_{s}^2$. For a helium-rich WN wind \\={m} = (4/3)m$_{p}$ where m$_{p}$ is the proton mass. The value kT$_{cool}$ $\\approx$ 0.6 keV gives a typical shock speed v$_{s}$ $\\approx$ 480 km s$^{-1}$ but the range in temperatures of $\\approx$0.3 - 0.8 keV inferred from the FWHM of the DEM model implies a range of shock speeds v$_{s}$ $\\approx$ 340 - 550 km s$^{-1}$. These values overlap the range of shock velocity jumps $\\Delta$v = 500 - 1000 km s$^{-1}$ derived in numerical simulations of radiatively driven winds (Owocki et al. 1988).   The average filling factor $f$ of X-ray emitting plasma in the wind can be estimated using the procedure given in SZGS02 and the mass-loss parameters from Section 4.2.1. We define $f$ = EM$_{x}$/EM$_{tot}$ where EM$_{x}$ is the volume emission measure of the X-ray emitting plasma and EM$_{tot}$ is the total volume emission measure in the wind. For a He-dominated wind we obtain $f$ = 5.34 $\\times$ 10$^{-3}$(R$_{*}$/R$_{\\odot}$)$norm$/d$_{kpc}$, where $norm$ is the normalization factor from XSPEC VAPEC models. Since only the cool X-ray component can be attributed to radiative wind shocks, we set $norm$ = $norm_{cool}$ from Table 3 and obtain $f$ = 1.8 $\\times$ 10$^{-7}$(R$_{*}$/R$_{\\odot}$)/d$_{kpc}$. For radii of a few solar radii (Hillier 1987) and the range of estimated distances 0.58 - 1.8 kpc, the filling factor need not be larger than $f$ $\\sim$ 10$^{-6}$.   The emergent cool emission detected by {\\em XMM-Newton} must originate at large distances from the star unless the wind is clumped. Using the same procedure as in SZGS02 along with an intermediate distance d = 1.2 kpc and the mass loss parameters in Sec. 4.2.1, the radius of optical depth unity at 1 keV is R$_{\\tau = 1}$(E = 1 keV) = 1.58 $\\times$ 10$^{14}$ cm $\\approx$ 10.6 AU. Assuming R$_{*}$ $\\approx$ 2  R$_{\\odot}$ as a representative value (Hillier 1987; Hamann \\& Koesterke 1998), then R$_{\\tau = 1}$(E = 1 keV) $\\approx$ 1130 R$_{*}$. The harder emission could be coming from much smaller radii. Specifically, R$_{\\tau = 1}$(E = 4 keV) $\\approx$ 0.3 AU assuming that the wind absorption cross-section for X-rays $\\sigma_{w}$ scales with energy according to $\\sigma_{w}$ $\\propto$ E$^{-2.5}$ (Fig. 1 of Ignace, Oskinova, \\& Foullon 2000).  The shock speeds, filling factor and value of R$_{\\tau = 1}$(E = 1 keV) computed above are nearly identical to those derived previously for WR 110 (SZGS02). The unit optical depth calculations for EZ CMa and WR 110 suggest that soft X-rays ($\\sim$1 keV) emerge at many hundreds of stellar radii in WR stars assuming homogeneous winds, but smaller emergent radii would be possible for clumped winds. Similar conclusions have been reached for some O-type stars (e.g. Hillier et al. 1993). Detailed hydrodynamic simulations are now needed to determine if instability-generated wind shocks can persist to hundreds of radii in WR stars, analogous to those recently undertaken for OB stars by Runacres \\& Owocki (2002).    \\subsection{On the Possibility of a Compact Companion}  Binarity has been suggested as one possible means of explaining the 3.765 day optical variability of EZ CMa. An estimate of the mass of the putative companion M$_{comp}$ = 1.3 ($\\pm$0.4) M$_{\\odot}$ derived by Firmani et al. (1980) raised speculation that EZ CMa could have a compact companion (c), making it a rare WR $+$ c system (Lamontagne et al. 1986; White \\& Long1986). But, it was shown that the X-ray luminosity of EZ CMa is about three orders of magnitude lower than the accretion luminosity L$_{x}^{(acc)}$ $\\sim$10 $^{36}$ ergs s$^{-1}$ expected for accretion of the W-R wind onto a neutron star (Stevens \\& Willis 1998). Similar arguments against a black hole companion based on {\\em ASCA} luminosities were given by Skinner et al. (1997).   The {\\em XMM-Newton} results  confirm the above luminosity deficit, giving an unabsorbed luminosity L$_{x}$ = (0.2 - 10 keV) = 3.46 $\\times$ 10$^{32}$d$_{kpc}^{2}$ ergs s$^{-1}$. At the upper end of current distance estimates, d = 1.8 kpc and L$_{x}$ = 10$^{33.0}$ ergs s$^{-1}$. If only the contribution of the hard component is considered, then L$_{x,hard}$(0.2 - 10 keV) = 10$^{32.5}$ ergs s$^{-1}$. Thus, if a neutron star companion is present then some mechanism such as rapid rotation near breakup (Davidson \\& Ostriker 1973) is needed to inhibit accretion.  The strong similarity between EZ CMa and WR 110 (Table 4) presents a new challenge for the compact companion hypothesis.  WR 110 has so far shown no clear signs of binarity or periodic optical variability and has not previously been proposed as a candidate WR $+$ c system. If EZ CMa has a compact companion, then it is not clear why its X-ray and radio properties would so closely mimic those of another WN star for which evidence of a compact companion is lacking.    \\subsection{On the Possibility of a Normal Stellar Companion}  An interpretation of the X-ray emission in terms of accretion onto a compact companion is questionable on the above grounds. However, it is  more difficult to rule out a normal (nondegenerate) stellar companion. Such a companion is expected to be much less massive than the W-R star given that the radial velocity variations reported by Firmani et al. (1980) are of low amplitude.  Arguments for binarity have recently been strengthened on the basis of long-term  coherent optical variability seen in data sets spanning more than 15 years  with a well-determined period P = 3.765 $\\pm$ 0.0001 days (Georgiev et al. 1999). Georgiev et al. have argued that if the optical variability is due to a  stellar companion then it is most likely orbiting very close to EZ CMa near the base of the wind in order to explain variations in the N V lines. We show below that a close companion could account for the hard X-ray emission detected by {\\em XMM-Newton}.  To constrain the separation, we assume that P = 3.765 d is an orbital period and that the companion mass is much less than that of the W-R star M$_{wr}$ $\\approx$ 16 M$_{\\odot}$  (Hamann \\& Koesterke 1998). Kepler's third law then gives the separation in AU as a$_{AU}$ $\\approx$ 0.12, or equivalently a $\\approx$ 25 R$_{\\odot}$. A nearly identical separation is obtained if one uses the values of M$_{wr}$ $\\approx$  10 M$_{\\odot}$ and the companion mass M$_{comp}$ $\\approx$ 1.3 M$_{\\odot}$ adopted by Firmani et al. (1980).  We now assume that the hard X-ray component is produced by the W-R wind shocking onto the lower mass companion and that the W-R wind is dominant. In this case the contact surface is the surface of the companion star, as discussed in more detail by Luo, McCray, \\& MacLow (1990). Because of the close separation, radiative cooling may be important (eq. [8] of Stevens et al. 1992; eq. [52] of Usov 1992). We thus consider both adiabatic and radiative shocks. The close separation also raises the question of whether hard X-rays could escape from the overlying W-R wind and be detected. As already noted (Sec. 5.2), the radius of optical depth unity at 4 keV is R$_{\\tau = 1}$(4 keV) $\\approx$ 0.3 AU, so some absorption of the hard X-rays could occur. However, the absorption will depend critically on wind properties such as the clumping factor, and the actual value of R$_{\\tau = 1}$ in a clumped wind would be less than that given above, which is based on the assumption of a spherical, homogeneous wind. Thus, without more specific information on wind geometry and homogeneity the escape of hard X-rays ($\\sim$4 keV) from radii smaller than 0.3 AU is not precluded.  The inferred radius of the companion R$_{comp}$ is obtained by equating the unabsorbed luminosity of the hard component (Table 3) with the predicted shock luminosity, where the predicted shock luminosity is different in the adiabatic and radiative cases (eqs. [79], [80] of Usov 1992, respectively). Using the mass loss parameters in Sec. 4.2.1 and the unabsorbed hard-component flux F$_{x}^{(hot)}$(0.2 - 10 keV) = 7.8 $\\times$ 10$^{-13}$ ergs cm$^{-2}$ s$^{-1}$, the adiabatic case gives a$_{AU}$ = 0.20d$_{kpc}^{0.25}$(R$_{comp}$/R$_{\\odot}$)$^{0.75}$. For a$_{AU}$ = 0.12 and the probable distance range d$_{kpc}$ = 0.575 - 1.8 the companion radius is R$_{comp}$ $\\approx$ 0.4 - 0.6 R$_{\\odot}$. A similar calculation in the radiative case gives R$_{comp}$ $\\approx$ 0.12 - 0.16 R$_{\\odot}$. In the above, we have used the luminosity in a specific energy range (0.2 - 10 keV) as an approximation for the total X-ray luminosity integrated over all frequencies. In general, the luminosity within a specific bandpass will be less than the value over all frequencies so the derived values of R$_{comp}$  are in fact lower limits.   The observed temperature kT$_{hot}$ $\\approx$ 3.5 [3.0 - 4.2] keV is also compatible with a shocked companion interpretation. The observed temperature derived from spectral fits is an average that includes contributions from the hottest shocked plasma at kT$_{s,max}$ along the line-of-centers and cooler plasma downstream. Thus, the observed value will in general be {\\em less} than the maximum temperature. The difference between the observed temperature and kT$_{s,max}$ will depend on several poorly known factors including wind chemical composition and the importance of radiative cooling. As a rough estimate, one expects an observed temperature kT $\\approx$ 0.8kT$_{s,max}$ in the adiabatic case (eq. [83] ff. of Usov 1992).  For EZ CMa, the maximum predicted temperature for an adiabatic shock in a He-dominated wind is kT$_{s,max}$ $\\approx$ 7.5 keV (Sec. 5.2), assuming that the wind has reached terminal speed v$_{\\infty}$ = 1700 km s$^{-1}$ at the shock. However, at the close separation of interest here (a $\\approx$ 25 R$_{\\odot}$), the wind may not have reached terminal speed. Using the velocity profile and stellar radius R$_{wr}$ = 2.5 R$_{\\odot}$ given by Hillier (1987), we obtain a/R$_{wr}$ $\\approx$ 10 and the wind is at or near terminal speed when it impacts the companion. But, using the radius R$_{wr}$ = 3.5 R$_{\\odot}$ and $\\beta$ = 3 wind velocity law of Schmutz (1997) then a/R$_{wr}$ $\\approx$ 7.1 and the wind speed is v$_{wr}$ $\\approx$ 0.64 v$_{\\infty}$. In this case, kT$_{s,max}$ $\\approx$ 3.1 keV. Thus, based on the rough approximation that the observed temperature should be $\\approx$0.8 kT$_{s,max}$, we would expect observed values in the range $\\approx$2.5 - 6 keV. The values measured from spectral fits are well within this range.  We thus conclude that the X-ray luminosity and observed temperature of the hard component are compatible with the shocked companion hypothesis. If the companion mass is much less than that of the W-R star then the inferred companion radius is at least R$_{comp}$ $\\approx$ 0.2 - 0.6 R$_{\\odot}$. This would correspond to a main-sequence (MS) M-type star, but clearly such a low mass star could not yet have reached the MS if it formed contemporaneously with the W-R star. Thus, a pre-main-sequence (PMS) companion would seem more likely. It may be relevant here that some W-R stars such as $\\gamma^{2}$ Vel are now believed to be associated with PMS objects (Pozzo et al. 2000; Skinner et al. 2001). If the putative companion is indeed a late-type star then interaction of the W-R wind with the companion's magnetic field would be an important factor to include in more detailed models. Our analysis shows that the observed temperature of the hot component is somewhat less than that expected for an adiabatic shock if the wind has reached terminal speed. This could be an indication that radiative cooling is important and the adiabatic approximation is breaking down at close separation, or perhaps that the W-R wind has not reached terminal speed at the shock. Radiative braking of the WR wind by the the companion star can also lead to lower X-ray temperatures than predicted by simple models (Gayley, Owocki, \\& Cranmer 1997) but this effect would be of little importance for a lower mass companion that is less luminous than the WR star."
    },
    "0207/astro-ph0207492_arXiv.txt": {
        "abstract": "It is shown explicitly that if only the zeroth Landau level is occupied by electrons in  strongly magnetized neutron star matter in $\\beta$-equilibrium, then both the proton and neutron matter sectors become spin symmetric. Whereas, the study of Pauli para-magnetism of neutron matter sector shows that such a scenario is physically impossible. It is also shown that in dense stellar hadronic matter in $\\beta$-equilibrium in presence of strong quantizing magnetic field, with $\\sigma-\\omega-\\rho$ exchange type mean field interaction and with the inclusion of magnetic dipole moments does not allow electrons to occupy only the zeroth Landau level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207201_arXiv.txt": {
        "abstract": "{ The total and polarized radio continuum emission of 20 barred galaxies was observed with the Very Large Array (VLA) at $\\lambda$3, 6, 18 and 22~cm and with the Australia Telescope Compact Array (ATCA) at $\\lambda$6~cm and 13~cm. Maps at 30\\arcsec\\ angular resolution are presented here. Polarized emission (and therefore a large-scale regular magnetic field) was detected in 17 galaxies. Most galaxies of our sample are similar to non-barred galaxies with respect to the radio/far-infrared flux correlation and equipartition strength of the total magnetic field. Galaxies with highly elongated bars are not always radio-bright. We discuss the correlation of radio properties with the aspect ratio of the bar and other measures of the bar strength. We introduce a new measure of the bar strength, $\\Lambda$, related to the quadrupole moment of the bar's gravitational potential. The radio surface brightness {\\rm I} of the barred galaxies in our sample is correlated with $\\Lambda$, $I\\propto\\Lambda^{0.4\\pm0.1}$, and thus is highest in galaxies with a long bar where the velocity field is distorted by the bar over a large fraction of the disc. In these galaxies, the pattern of the regular field is significantly different from that in non-barred galaxies. In particular, field enhancements occur upstream of the dust lanes where the field lines are oriented at large angles to the bar's major axis. Polarized radio emission seems to be a good indicator of large-scale non-axisymmetric motions. ",
        "introduction": "Polarization observations in radio continuum have revealed basic properties of interstellar magnetic fields in a few dozen spiral galaxies (Beck et al.\\ \\cite{beck+96}, Beck\\ \\cite{beck00}). Large-scale regular fields form spiral patterns with pitch angles similar to those of the optical spiral arms. The strongest {\\it regular\\/} fields usually occur between the optical arms, sometimes concentrated in `magnetic arms' (Beck \\& Hoernes \\cite{beck+hoernes96}). The {\\it total} (= polarized + unpolarized) nonthermal (synchrotron) radio emission is a tracer of the {\\it total\\/} field which comprises both regular and random field components. It generally peaks on the optical arms because the random field is strongest there. This distinction implies that the regular and random magnetic fields are maintained and affected by different physical processes.  Spiral patterns of the regular magnetic field are believed to be generated by dynamo action in a differentially rotating disc (Beck et al.\\ \\cite{beck+96}). The dynamo reacts or interacts with non-axisymmetric disturbances like density waves (Mestel \\& Subramanian\\ \\cite{mestel+91}, Rohde et al.\\ \\cite{rohde+99}), but little is known about the effects of bar-like distortions. Chiba \\& Lesch (\\cite{chiba+lesch94}) suggested that a bar may excite higher dynamo modes, while Moss et al. (\\cite{moss+98}) found from their models a mixture of modes with rapidly changing appearance.  Radio observations of barred galaxies are rare. The angular resolution of the maps in Condon's (\\cite{condon87}) atlas was insufficient to distinguish emission from the bar, the spiral arms and the halo. Another survey of barred galaxies in radio continuum by Garc\\'{\\i}a-Barreto et al.\\ (\\cite{garcia+93}) had even lower angular resolution; neither survey included polarization. The first high-resolution radio map of a barred galaxy, NGC~1097 (Ondrechen \\& van der Hulst\\ \\cite{ondrechen+83}), showed narrow ridges in total intensity coinciding with the dust lanes, which are tracers of compression regions along the leading (with respect to the sense of rotation) edge of the bar. A similar result was obtained for M83 (Ondrechen\\ \\cite{ondrechen85}) which hosts a bar of smaller size than NGC~1097. No polarization could be detected in NGC~1097 by Ondrechen \\& van der Hulst (\\cite{ondrechen+83}). Radio observations of NGC~1365 at $\\lambda\\lambda20,\\ 6$ and $2\\cm$, restricted to a central region, have revealed similar features (J\\\"ors\\\"ater \\& van Moorsel\\ \\cite{joersaeter+moorsel95}). The first detection of polarized radio emission from the bar region was reported by Ondrechen (\\cite{ondrechen85}) for M83, with a mean fractional polarization at $\\lambda6$~cm of 25\\%.  Neininger et al.\\ (\\cite{neininger+91}) mapped the polarized emission from M83 at $\\lambda2.8$\\,cm. They showed that the regular magnetic field in the bar region is aligned with the bar's major axis.  Observed with higher resolution, the regular field is strongest at the leading edges of the bar of M83 (Beck\\ \\cite{beck00}).  Another barred galaxy which has been studied in detail in radio polarization is NGC~3627 (Soida et al.\\ \\cite{soida+01}). The regular field in the bar region is again aligned parallel to the bar's major axis, being strongest at the leading edges of the bar. However, east of the bar the field behaves anomalously, forming a `magnetic arm' crossing the gaseous arm.  The first high-resolution polarization observations of a galaxy with a massive bar, NGC~1097, were presented by Beck et al. (\\cite{beck+99}). The magnetic field lines in and around the bar appear to follow the velocity field of the gas expected from a generic gas dynamic model (Athanassoula \\cite{atha92}). The regular magnetic field outside the bar region has a spiral pattern similar to that seen optically. A narrow ridge of greatly reduced polarized intensity indicates the deflection of the field lines in a shear shock (the dust lane), but the magnetic field lines turn more smoothly than the gas streamlines (Moss et al.\\ \\cite{moss+01}, hereafter Paper II). Velocity fields are available from HI observations only for the outer parts of NGC~1097 (Ondrechen et al.\\ \\cite{ondrechen+89}) and from CO observations only for the circumnuclear ring (Gerin et al.\\ \\cite{gerin+88}).  NGC~1097 is one of the objects in our sample of barred galaxies observed with the VLA and the ATCA. In this paper we present the full set of radio maps of our survey, smoothed to a common resolution of 30\\arcsec, and give an overview of their salient properties. Higher-resolution maps of NGC~1097, 1365 and 7479 will be presented and discussed in subsequent papers. New dynamo models for barred galaxies are discussed in Paper II. Further details on the magnetic fields in NGC1672, 2442 and 7552 will be given by Harnett et al.\\ (\\cite{harnett+02}, hereafter Paper III).  \\scriptsize \\begin{table*}[htb] \\caption{The VLA sample of barred galaxies} \\label{tab:vla-sample} \\scriptsize \\begin{tabular}{llllllllllllllll} \\hline \\noalign{\\smallskip} NGC &Hubble&Lum.&RC3 &R.A. &Dec. &$d_{25}$ &$q_{25}$ &$v_{\\rm GSR}$ &$D$ &$i$ &PA &$b/a$ &$2a/$ &$S_{60\\mu \\rm m}$ &$S^{\\rm tot}_{20\\rm cm}$ \\\\ &type&class&class&(2000)&(2000)&[\\arcmin ] & &[km/s]&[M&[\\degr ] &[\\degr ] & &$d_{25}$ &[Jy] &[mJy]\\\\ &(1)&(1)&(2) &[h m s]&[\\degr\\ \\arcmin\\ \\arcsec] &(2) &(2) &(2) &pc]& & & & &(3) &(4)\\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 1097  &SBbc(rs) &I-II &SBS3 &02 46 19.0 &$-$30 16 21 &\\pheins 9.3 &1.48 &        1193 & 16 & 45 & 135  & [0.4] & 0.37 & 45.9 & 415\\\\ 1300  &SBb(s) &I.2     &SBT4 &03 19 40.9 &$-$19 24 41 &\\pheins 6.2 &1.51 &        1496 & 20 & 35 &\\pheins  86 & [0.3] & 0.41 &\\pheins 2.4 &\\pheins  35\\\\ 1365  &SBb(s) &I      &SBS3 &03 33 36.7 &$-$36 08 17 &    11.2 &1.82 &        1541 & 19 & 40 &\\pheins  40  & 0.51 & 0.47 & 78.2 & 530\\\\ 2336  &SBbc(r) &I      &SXR4 &07 27 04.4 &$+$80 10 41 &\\pheins 7.1 &1.82 &        2345 & 31 & 59 & 178  & 0.41 & 0.17 &\\pheins 1.0 &\\pheins  18\\\\ 3359  &SBc(s) &I.8     &SBT5 &10 46 37.8 &$+$63 13 22 &\\pheins 7.2 &1.66 &        1104 & 15 & 55 & 170  & 0.32 & 0.25 &\\pheins 4.1 &\\pheins  50\\\\ 3953  &SBbc(r) &I-II  &SBR4 &11 53 49.6 &$+$52 19 39 &\\pheins 6.9 &2.00 &        1122 & 15 & 61 &\\pheins  13 & 0.89 & 0.17 &\\pheins 2.9 &\\pheins 41\\\\ 3992  &SBb(rs) &I      &SBT4 &11 57 36.3 &$+$53 22 31 &\\pheins 7.6 &1.62 &        1121 & 15 & 59 &\\pheins  67 & 0.58 & 0.27 &\\pheins $\\simeq$3 &\\pheins 21\\\\ 4535  &SBc(s) &I.3     &SXS5 &12 34 20.4 &$+$08 11 53 &\\pheins 7.1 &1.41 &        1892 & 16 & 26 &\\pheins  28  &[0.6] &[0.1] &\\pheins 6.5 &\\pheins 65\\\\ 5068  &SBc(s) &II-III &SXT6 &13 18 55.4 &$-$21 02 21 &\\pheins 7.2 &1.15 &\\pheins  550 &\\pheins 7 & 29 & 110  & 0.44 & 0.16 &\\pheins 2.3 &\\pheins 39\\\\ 7479  &SBbc(s) &I-II  &SBS5 &23 04 57.2 &$+$12 19 18 &\\pheins 4.1 &1.32 &        2544 & 34 & 45 &\\pheins  25  & 0.41 & 0.46 & 12.1 & 109 \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\end{tabular} \\vbox{ \\noindent References: (1) Sandage \\& Tammann (\\cite{sandage+tammann81}); (2) de Vaucouleurs et al.\\ (\\cite{vaucouleurs+91}); (3) Fullmer \\& Lonsdale (\\cite{fullmer+lonsdale89}); (4) Condon (\\cite{condon87}). } \\end{table*}    \\begin{table*}[htb] \\caption{The ATCA sample of barred galaxies} \\label{tab:atca-sample} \\scriptsize \\begin{tabular}{llllllllllllllll} \\hline \\noalign{\\smallskip} NGC &Hubble&Lum.&RC3 &R.A. &Dec. &$d_{25}$ &$q_{25}$ &$v_{\\rm GSR}$ &$D$ &$i$ &PA &$b/a$ &$2a/$ &$S_{60\\mu \\rm m}$ &$S^{\\rm tot}_{6\\rm cm}$ \\\\ &type&class&class&(2000) &(2000) &[\\arcmin ] & &[km/s]&[M&[\\degr ]&[\\degr ]& &$d_{25}$ &[Jy] &[mJy] \\\\ &(1)&(1)&(2) &[h m s]&[\\degr\\ \\arcmin\\ \\arcsec] &(2) &(2) &(2) &pc]& & & & &(3) &(4)\\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\pheins 986   &SBb(rs) &I--II  &SBT2  &02 33 34.3 &$-$39 02 43 &\\pheins 3.9 &1.32 &       1907 & 25 &?&\\pheins?& [0.5]& 0.46 & 23.1 &\\pheins 40\\\\ 1313  &SBc(s) &III-IV &SBS7  &03 18 15.5 &$-$66 29 51 &\\pheins 9.1 &1.32 &\\pheins 292 &\\pheins4 & 38 & 170 & 0.63 & 0.31 & 10.4 &\\pheins 59\\\\ 1433  &SBb(s) &I-II   &PSBR2 &03 42 01.4 &$-$47 13 17 &\\pheins 6.5 &1.10 &\\pheins 920 & 12 & 27 &\\pheins  17 & 0.33 & 0.36 &\\pheins 3.3 &\\pheins --\\\\ 1493  &SBc(rs) &III    &SBR6  &03 57 27.9 &$-$46 12 38 &\\pheins 3.5 &1.07 &\\pheins 900 & 12 & 30 &\\pheins ?  & 0.32 & 0.18 &\\pheins 2.2 &\\pheins -- \\\\ 1559  &SBc(s) &II.8    &SBS6  &04 17 37.4 &$-$62 47 04 &\\pheins 3.5 &1.74 &       1115 & 15 & 55 &\\pheins  65 &[0.3] &[0.2] &23.8 & 120\\\\ 1672  &SBb(rs) &II     &SBS3  &04 45 42.2 &$-$59 14 57 &\\pheins 6.6 &1.20 &       1155 & 15 & 39 & 170  & 0.41 & 0.68 & 34.8 & 100\\\\ 2442  &SBbc(rs) &II    &SXS4P &07 36 23.9 &$-$69 31 50 &\\pheins 5.5 &1.12 &       1236 & 16 & 24 &\\pheins  40 & [0.5]& 0.42 &$\\simeq$22&\\pheins 70\\\\ 3059  &SBc(s) &III     &SBT4  &09 50 08.1 &$-$73 55 17 &\\pheins 3.6 &1.12 &       1056 & 14 &?&\\pheins?&[0.3] &[0.2] &\\pheins 9.6 &\\pheins --\\\\ 5643  &SBc(s) &II-III &SXT5  &14 32 41.5 &$-$44 10 24 &\\pheins 4.6 &1.15 &       1066 & 14 &?&\\pheins?&[0.4] &[0.35] & 18.7 &\\pheins 64\\\\ 7552  &SBbc(s) &I-II  &PSBS2 &23 16 11.0 &$-$42 35 01 &\\pheins 3.4 &1.26 &       1568 & 21 & 31 &\\pheins\\pheins 1  & 0.29 & 0.59 & 72.9 & 140\\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\end{tabular} \\vbox{ \\noindent References: (1)--(3) see Table 1; (4) Whiteoak (\\cite{whiteoak70}). } \\end{table*} \\normalsize ",
        "conclusions": "We observed a sample of 20 barred galaxies with the VLA and ATCA radio telescopes. Polarized radio emission was detected in 17 galaxies.  The flux densities in the radio continuum and the far-infrared spectral ranges are closely correlated in our sample. The average radio/far-infrared flux density ratio and equipartition strength of the total magnetic field are similar to those in non-barred galaxies. These properties are apparently connected to the star formation rate and possibly controlled by the density of cool gas. Radio surface brightness and present star formation activity are highest for galaxies with a high content of molecular gas and long bars where the velocity field is distorted over a large volume. The radio surface brightness is correlated with a newly introduced measure of bar strength proportional to the quadrupole moment of the gravitational potential. However, a few galaxies with strong bars are not radio bright, possibly because their molecular gas has been depleted in a star formation burst.  In barred galaxies with low or moderate radio surface brightness, the regular field (traced by the polarized radio emission) is strongest between the optical spiral arms (e.g.\\ NGC~3359 and 4535) or has a diffuse distribution (e.g.\\ NGC~3059 and 3953). In radio-bright galaxies, the pattern of the regular field can, however, be significantly different: the regular magnetic field may have a broad local maximum in the bar region upstream of the dust lanes, and the field lines are oriented at large angles with respect to the bar (NGC~1097, 1365, 1672 and 7552). We propose that shear in the velocity field around a large bar may enhance dynamo action and explain the observed strong regular fields. Strong bar forcing induces shear in the velocity field and enhancements in the regular magnetic field, and polarized emission traces such shear motions.  The southern galaxies NGC~986, 1559, 1672 and 7552 show strong polarization and are promising candidates for further studies with high resolution. Circumnuclear rings are already known to exist in NGC~1097, NGC~1365, NGC~2442 and NGC~7552 and should be searched for in NGC~986, 1559 and 1672."
    },
    "0207/astro-ph0207037_arXiv.txt": {
        "abstract": "{We predict mass-loss rates for the late evolutionary phases of low-mass stars, with special emphasis on the consequences for the morphology of the Horizontal Branch (HB). We show that the computed rates, as predicted by the most plausible mechanism of radiation pressure on spectral lines, are too low to produce EHB/sdB stars. This invalidates the scenario recently outlined by Yong et al. (2000) to create these objects by mass loss {\\it on} the HB. We argue, however, that mass loss plays a role in the distribution of rotational velocities of hot HB stars, and may -- together with the enhancement of heavy element abundances due to radiative levitation -- provide an explanation for the so-called ``low gravity'' problem. The mass loss recipe derived for hot HB (and extreme HB, sdB, sdOB) stars may also be applied to post-HB (AGB-manqu{\\'e}, UV-bright) stars over a range in effective temperatures between 12\\,500 -- 40\\,000 K. ",
        "introduction": "\\label{s_intro}  Over the last decades, both observational and theoretical efforts have been devoted to the investigation of the observed distribution of stars along the Horizontal Branch (HB) of galactic Globular Clusters (GCs). Although canonical stellar evolution theory has provided a general consensus on the evolutionary phase corresponding to the HB sequence, and convincingly demonstrated that its morphology is most strongly affected by cluster metallicity (the \\lq{first}\\rq\\ parameter; Sandage \\& Wallerstein 1960), many problems remain. The most striking controversy involves the wide variety of HB morphologies among clusters with similar metallicities (the \\lq{second parameter}\\rq~problem; Sandage \\& Wildey 1967, van den Bergh 1967). Candidate second parameters are cluster age (e.g. Lee et al. 1994 and references therein), mass loss along the Red Giant Branch (RGB) (Catelan et al. 2001 and references therein), rotation and deep helium mixing (Sweigart 1997), dynamical interactions involving binaries and even planets (Soker 1998), as well as environmental effects in high-density environments (Fusi Pecci et al. 1993).  The identification of the second parameter is especially relevant to the formation of Extremely blue HB (EHB) stars, which are thought to be responsible for the ultraviolet upturn phenomenon in elliptical galaxies (Greggio \\& Renzini 1990; Dorman et al. 1995). The presence of EHB stars as blue \\lq{tails}\\rq~in clusters (Ferraro et al. 1998; Piotto et al. 1999), as well as sdB/sdO stars in the field (Greenstein 1971; Green et al. 1986), has inspired modern-day research to explain their formation both through mechanisms that produce high mass loss along the RGB (Soker et al. 2001 and references therein), as well as through binarity (Mengel et al. 1976, Heber et al. 2002).  Further puzzles in HB morphology concern the issues of HB \\lq{gaps}\\rq~(Newell 1973) -- specific regions along the branch that are significantly underpopulated\\footnote{It has been claimed that the positions of the gaps along the HB in different galactic GCs are the same within current empirical uncertainties (Ferraro et al. 1998), however it is not clear whether these gaps mark regions with specific effective temperatures, or whether they correspond to constant mass loci (Piotto et al. 1999). Note that Catelan et al. (1998) have challenged the existence of gaps at the same positions in all clusters.}, and a relatively new, unexplained, but ubiquitous feature is the so-called Str{\\\"o}mgren u-jump at an effective temperature of \\teff\\ $\\simeq$\\,11\\,000 K (Grundahl et al. 1999), possibly coinciding with a jump in $\\log g$ (Moehler et al. 1995), and an unexplained absence of fast rotators above this temperature (Behr et al. 2000; Recio-Blanco et al. 2002).  As discussed by Grundahl et al. (1999), the Str{\\\"o}mgren u-jump may be due to atmospheric diffusion by radiative levitation of heavy elements, as both Glaspey et al. (1989) and Behr et al. (1999) found striking abundance anomalies in blue HB stars, with iron enhancements of up to three times the solar value. Moehler et al. (2000) have shown that the enhancement of heavy elements in spectroscopic analyses may partially solve the problem of the anomalously low gravities along the blue HB, but the discrepancy is still present at the level of $\\Delta\\log{g}\\,\\approx\\,0.1$ dex for stars in the range $15\\,000\\,<\\teff\\,<\\,20\\,000\\,$K. Even more so, the first two mentioned HB features -- the blue tails and the gaps -- are still an enigma\\footnote{Note that Brown et al. (2001) and Sweigart et al. (2002) have provided a theoretical framework which could explain the hot gap in the HB of NGC\\,2808.}, and it is not at all obvious whether they originate from a mechanism working in a prior evolutionary phase (on the RGB) or if they are due to a process working \\lq{in situ}\\rq\\, once the star has settled on the HB. One of the options that may help in explaining the above-mentioned problems is mass loss {\\it on} the HB.  It is worth mentioning that the colour width of the hottest gap is so small that changes of the order of a few times $10^{-3}$\\,\\Msun\\ in the total mass are capable to move the star away from its initial location, far enough as to produce an underpopulated region in the H-R diagram (HRD). For this to occur one needs to identify a mass-loss mechanism which efficiency rapidly increases at the specific effective temperatures of the gap.  The hypothesis that a mass-loss mechanism may be at work during the HB evolution was first entertained by Wilson \\& Bowen (1984). They suggested that an increased mass-loss efficiency, when crossing the RR Lyrae instability strip, could provide an explanation for the HB mass distribution in a more natural way than the alternative of a stochastic variation in the amount of mass lost during the prior RGB phase. The topic of mass loss on the HB was further addressed by Koopmann et al. (1994), but they concluded that constant mass loss in the RR Lyrae strip was incapable of providing an explanation for the HB mass dispersion or the RR Lyrae period change distribution. Additionally, mass loss during the central He-burning phase was suggested by Michaud et al. (1985) and Bergeron et al. (1988) in order to explain the large silicon underabundances in some HB stars. More recently, Yong et al. (2000) performed accurate evolutionary computations with mass loss, and suggested that mass-loss rates of the order of $10^{-9}$ -- $10^{-10}$ \\msunyr\\ for HB stars in the metal-rich cluster NGC\\,6791 can force these stars to move to a bluer position and thus lead to the production of EHB stars. If correct, this scenario could provide an explanation for the presence of extended blue tails along the HB of some metal-rich GCs, such as NGC\\,6441 and NGC\\,6388 (Rich et al. 1997), although it would not be able to explain the upward sloping of the HB in these clusters (Raimondo et al. 2002 and references therein). The main problem with the proposed scenario, however, is that no physical mechanism for mass loss was proposed and that the adopted mass-loss rates were completely \\lq{ad hoc}\\rq, as there are neither observational data indicative of mass loss on the HB available, nor any predictions.  Our aim in the present paper is to alleviate current shortcomings by computing radiation-driven wind models and mass-loss rates for low-mass blue stars, and to subsequently investigate their influence on HB evolutionary models. Blue HB stars are located in a region of the HRD, where the stars are hot (with \\teff\\ between 10\\,000 and 35\\,000 K), and relatively bright, and radiation pressure forces can therefore be considered a natural driving mechanism. Although there may be other processes that could possibly drive a wind, such as pulsations \\footnote{Recently a new class of variable stars has been discovered in the field, the so-called EC14026 (Kilkenny et al. 1997), which have been identified as hot HB stars and their progeny. Even though no clear identification of similar variables in GCs have been obtained, there is no {\\it a priori} reason for the lack of this kind of pulsation among cluster HB stars.}, all other wind-driving options are much less well-understood than radiation pressure on spectral lines. Radiation-driven wind models have been developed in the 1970s by Lucy \\& Solomon (1970) and Castor et al. (1975). In more recent days, the models have been very successful in predicting the values observed in O supergiants (Vink et al. 2000). The direct application of these predictions to HB stars, such as the use of the mass-loss recipe provided by Vink et al. (2000) would however involve a rather large and dangerous extrapolation by four orders of magnitude in stellar luminosity.  As far as the \\lq{gaps}\\rq\\ along the HB are concerned, radiation-driven wind models for OB supergiants predict that the efficiency of mass loss jumps strongly by a factor of five at spectral type B1 (Vink et al. 1999, 2000). This is close to the position where the evidence for a gap in HB morphology is strongest. A mass-loss rate of the order of $10^{-10} - 10^{-11}$ \\msunyr\\ could be sufficient to explain the presence of the gap located at $\\teff$\\,$\\simeq$\\,20\\,000 K; given an HB evolutionary timescale of $\\approx10^8$ years with mass loss at this rate leads to a total amount of a few times $10^{-3}$ \\msun\\, sufficient to move an HB star by $\\approx$1000 K, and so creating a \\lq{gap}\\rq.  The above-mentioned issues, i.e. the presence of EHB stars, gaps, and anomalous abundances in HBs and sdB stars, prompted us to compute radiation-driven wind models for HB stars; to predict mass-loss rates for these objects, and subsequently explore their influence on evolutionary models. The mass loss computations may also provide valuable ingredients for HB angular momentum evolution and chemical separation calculations of sdB stars (see Unglaub \\& Bues 2001).  The outline of the paper is as follows. In the next section we describe the approach used for computing mass-loss rates, as well as the assumptions adopted in the numerical computations; in Sect.\\,3 we discuss the results concerning the mass-loss efficiency, where the dependence of \\mdot\\ on the main evolutionary parameters, the luminosity, effective temperature and stellar mass, as well as stellar metallicity, is presented. In Sect.\\,4, we provide an analytical relation for \\mdot\\ as a function of the quoted parameters, which is useful for computing the mass-loss rates in evolutionary computations, and we investigate the effects of our recipe on HB stellar evolution (Sect.\\,5). In Sect.\\,6, we study the implications of mass loss regarding the ``zoo'' of problems in HB morphology that occur for effective temperatures larger than $\\simeq$\\,10\\,000 K, in particular the effects of mass loss on rotational velocities and the $\\log g$ jump. Final remarks and conclusions will close the paper. ",
        "conclusions": "\\label{theend}  In this paper we have, for the first time, computed mass-loss rates for HB stars. We have shown that the computed rates, as predicted by the most plausible mechanism of radiation pressure on spectral lines, are too low to produce EHB/sdB stars. This invalidates the scenario outlined by Yong et al. (2000) to create these objects by excessive mass loss {\\it on} the HB. We argue, however, that mass loss plays a role in the distribution of rotational velocities of hot HB stars, and for the so-called $\\log g$ jump. The mass loss recipe derived in this paper is, strictly speaking, only valid for HB stars, but as there are hardly any mass-loss predictions available for low-mass blue stars, the recipe may also be applied to: post-HB, AGB-manqu{\\'e}, UV-bright stars, extreme helium stars, as long as the desired accuracy is within a factor of two, and as long as the effective temperatures are not higher than 40\\,000 K.  Although we have proposed a scenario where winds are ubiquitous for hot HB stars, and subsequently affect the rotational velocities, as well as the atmospheric parameters ($\\log g$), there is still a lot of work to be done.  First and foremost, spectral evidence for mass loss in HB and sdB stars ought to be sought to check whether the mass-loss rates, as derived in this paper, indeed occur. Diffusion calculations including mass loss for sdB stars (Unglaub \\& Bues 2001) suggest that our derived mass-loss rates are reasonable, but this is certainly not a model-independent check.  Second, evolutionary models including rotation (first steps have been undertaken by Sills \\& Pinsonneault 2000) and mass loss should be computed to see whether the absence of fast rotators for stars hotter than 11\\,000 K can indeed be due to the removal of angular momentum due to stellar winds.  Last but not least, systematic atmospheric analyses of hot HB stars accounting for the {\\sl actual} surface heavy elements distribution and including mass loss should be performed to see whether the $\\log g$ jump is indeed an artifact of the adopted hydrostatic model atmospheres. The current situation, where evolutionary models are not in agreement with the spectral analyses is highly undesirable, as this suggests that current stellar evolution theory is not only incapable of producing extreme HB stars, but that even ``normal'' blue HB stars pose a serious problem. In other words, a solution to the ``low gravity'' problem for hot HB stars could significantly enhance our current understanding of the later phases of stellar evolution."
    },
    "0207/hep-ph0207028_arXiv.txt": {
        "abstract": " ",
        "introduction": "CP violation may have played a crucial role in the formation of matter in the universe.  It must have, if the inflationary cosmology is right, thus ruling out the (very unpalatable) possibility that the baryon asymmetry was an initial condition of the Big Bang.  If inflation took place, as most of us believe based on the growing observational evidence, it made the universe devoid of matter and set the stage for reheating.  The baryon asymmetry then must have been produced at some later point through CP-violating processes. ",
        "conclusions": "There is every reason to believe that matter-antimatter asymmetry is a consequence of CP non-conservation in particle physics.  On the other hand, CP violation from the quark mixing is not sufficient for baryogenesis. This implies the existence of new, yet undiscovered, sources of CP violation in nature."
    },
    "0207/astro-ph0207546_arXiv.txt": {
        "abstract": "{We have observed the X-ray transient \\src\\ in quiescence with XMM-Newton. The observed spectrum is highly unusual being dominated by an emission feature at $\\sim$6.5~keV.  The spectrum can be fit using a partially covered power-law and Gaussian line model, in which the emission is almost completely covered (covering fraction of $ 0.98 _{-0.06}^{+0.02}$) by neutral material and is strongly absorbed with an \\nh\\ of ($ 5 _{-2}^{+3}$)~\\ttroisnh.  This absorption is local and not interstellar.  The Gaussian has a centroid energy of $6.4 \\pm 0.1$~keV, a width $ \\sigma <0.28$~keV and an equivalent width of $ 940 ^{+650}_{-460}$~eV.  It can be interpreted as fluorescent emission line from iron.  Using this model and assuming \\src\\ is at a distance of 5~kpc, its 0.5--10~keV luminosity is $ 3.5 \\times 10^{33}$~\\ergs. The Optical Monitor onboard XMM-Newton indicates a V magnitude of $11.86~\\pm~0.03$. The spectra of X-ray transients in quiescence are normally modeled using advection dominated accretion flows, power-laws, or by the thermal emission from a neutron star surface.  The strongly locally absorbed X-ray emission from \\src\\ is therefore highly unusual and could result from the compact object being embedded within a dense circumstellar wind emitted from the supergiant B[e] companion star. The uncovered and unabsorbed component observed below 5~keV could be due either to X-ray emission from the supergiant B[e] star itself, or to the scattering of high-energy X-ray photons in a wind or ionized corona, such as observed in some low-mass X-ray binary systems. ",
        "introduction": "\\label{sec:intro}  \\src\\ was discovered by the All-Sky Monitor onboard RXTE as a soft X-ray transient during an outburst in 1998 March 31 \\citep{0421:smith98iau}. This outburst was observed by CGRO \\citep{0421:paciesas98iau}, RXTE \\citep{0421:revnivtsev99al,0421:belloni99apj}, ASCA \\citep{0421:ueda98apjl} and \\sax\\ \\citep{0421:frontera98aa,0421:orr98aa}. The source brightened rapidly, reaching an intensity of $\\sim$2~Crab after a few hours, then quickly decayed with an initial {\\it e}-folding time of only 0.6~days before reaching quiescence in less than 2 weeks. This was the fastest rise and decay of any outburst from a soft X-ray transient \\citep[see e.g.,][]{chen97apj}. The outburst X-ray spectra from \\src\\ are complex and can not be fit by any of the models usually applied to soft X-ray transients.  The ASCA outburst spectrum was fitted with a two temperature optically thin thermal model and an additional broad Fe-K emission line at 6.4~keV.  Both \\sax\\ outburst spectra were described using a two temperature bremsstrahlung model and narrow emission line features identified with O, Ne/Fe-L, Si, S, Ca and Fe-K \\citep{0421:orr98aa}.  Emission lines at energies of $\\sim$6.5~keV and $\\sim$8~keV are detected in the RXTE outburst spectra.        Optical and radio observations allowed \\src\\ to be rapidly identified with \\cicam, also known as MWC~84 \\citep{0421:wagner98iau,0421:hjellming98iau,0421:robinson98iau}.  It is a frequently observed source in ultra-violet (UV), optical and infra-red (IR) wavelengths.  Its V magnitude over long term observations shows a $\\sim$0.4~magnitude amplitude variability and has a mean value of 11.6, both before (1989--1994) and after (1998--1999) the 1998 X-ray outburst \\citep{bergner95aa,0421:clark00aa}.  \\cicam\\ is a supergiant B[e] star \\citep{clark99aa,0421:robinson02apj}, or a sgB[e] star, following the notation of \\citet{lamers98aa}, i.e. a supergiant showing the B[e] phenomenon.  The B[e] phenomenon concerns many objects of different masses and evolutionary phases \\citep[see e.g.,][]{lamers98aa}. One of the common properties of stars exhibiting the B[e] phenomenon is the presence of forbidden emission lines in their optical spectra (the notation ``[e]'' refers to the one used for forbidden lines). Another common property is to show a strong IR excess attributed to hot circumstellar dust. In these respects, stars with the B[e] phenomenon clearly differ from the ordinary Be stars which are rapidly rotating stars near the main sequence losing mass in an equatorial wind. In practice, the spectroscopic and photometric properties of stars with the B[e] phenomenon are also easily distinguished from those of ordinary Be stars. \\cicam/\\src\\ is the first high-mass X-ray binary (HMXB) with a sgB[e] mass donor companion. Another source suspected to be a HMXB with a mass donor showing the B[e] phenomenon is the optical/X-ray source HD~34921/1H~0521+37 \\citep{clark99aa}.  Adopting the classification criteria and notation of \\citet{lamers98aa}, \\citet{clark99aa} identify the companion star in this system as an ``unclB[e] star'' (unclassified B[e] star).     Optical high-dispersion spectroscopy of \\cicam\\ led \\citet{0421:robinson02apj} to the conclusion that the sgB[e] star emits a two component wind.  One component is a hot, high-velocity wind. The other component is a cool, low-velocity and very dense (electron number density log~\\nel~$>~9.5$) wind. The wind is roughly spherical and continuously replenished. The mass-loss rate due to the wind is very high: \\mdot~$> 10^{-6}$~\\msunyear. This wind fills the space around the sgB[e] star and, from the size of the IR-emitting dust shell, extends to a radius between 13 and 50~AU.  Thus, the circumstellar material around \\cicam\\ is much denser, far more extended, and much less confined to the equatorial plane than the circumstellar material around a Be star.  The passage of the compact X-ray source through such a complex and dense environment is likely to strongly affect the X-ray properties of the source.     Within this picture, \\citet{0421:robinson02apj} suggest that the 1998 outburst was caused by the same disk instability mechanism responsible for the outbursts in X-ray novae, i.e. by an instability in the accretion disk around the compact object \\citep[see e.g.,][]{lasota01nar}. It would thus differ from the outbursts observed in ordinary Be HMXB that recur at multiples of the orbital period, when the compact object comes close to the Be star at periastron and plunges into its equatorial wind.     The distance to \\src\\ is uncertain.  Based on optical spectroscopic properties, on radial velocity measurements of \\cicam\\ and on considerations about the structure of the Galaxy, \\citet[Sect.~2.3]{0421:robinson02apj} estimate that the distance to the source is much larger than the $\\lesssim$2~kpc previously considered.  \\citet{0421:robinson02apj} use a distance of 5~kpc and note that it is likely to be a lower limit to the true distance which could be up to 10~kpc. In this paper, we assume a distance of 5~kpc. This distance makes \\src\\ among the most luminous transients. The 2--25~keV luminosity at the peak of the outburst was $3.0 \\times 10^{38}$~\\ergs, assuming the revised distance of 5~kpc \\citep{0421:orlandini00aa,0421:robinson02apj}.  The unusual nature of \\cicam\\ makes the interstellar absorption towards the star difficult to estimate.  From an UV spectrogram, \\citet{0421:robinson02apj} derive a differential extinction E(B--V) of $0.85 \\pm 0.05$, but do not attempt to separate circumstellar from interstellar extinction.  From an analysis of diffuse interstellar bands in the optical spectrum of \\cicam, \\citet{0421:clark00aa} derive an interstellar E(B--V) of $0.65 \\pm 0.20$ and an ${\\rm A_v}$ of $2.0 \\pm 0.6$, which implies an interstellar X-ray absorption, \\nh, of $(3.6 \\pm 1.1)$~\\tunnh\\ \\citep[Sect. 3]{0421:parmar00aa}.  Extinction at soft X-ray wavelengths yielded an \\nh\\ of $(3.76 \\pm 0.36)$~\\ttnh\\ near the peak of the outburst, and the \\nh\\ decreased to $\\sim$2.2~\\tunnh\\ as \\src\\ approached quiescence \\citep{0421:belloni99apj}. This rapid change in the X-ray extinction, as well as the change in the IR flux after the outburst \\citep{0421:clark00aa}, indicate that much of the extinction to \\cicam\\ is local, not interstellar \\citep{0421:robinson02apj}.   No bursts, pulsations or quasi-periodic oscillations have been detected from \\src\\ \\citep{0421:belloni99apj}.  The large ratio of peak to quiescent luminosity is taken as evidence for the compact object being a black hole \\citep{0421:robinson02apj}. Furthermore, \\cicam\\ has been detected as a relatively bright radio source \\citep{0421:hjellming98iau} which is more typical of black hole candidates than neutron star systems \\citep[see][Sect. 5]{0421:belloni99apj}.   \\src\\ was observed in quiescence by \\sax\\ on 1998 September 03, 1999 September 23 and 2000 February 20 \\citep{0421:orlandini00aa,0421:parmar00aa}. In 1998, the source was soft (power-law photon index, \\phind, of $4.0 ^{+1.9} _{-0.9}$) with a low \\nh\\ of $(1 ^{+5} _{-1}$)~\\tunnh. In 1999, the source had hardened (\\phind~$= 1.86 ^{+0.27} _{-0.32}$) and brightened and became strongly absorbed with an \\nh\\ of $(4.0~\\pm~0.8)$~\\ttroisnh.  There is evidence for a narrow emission line in both spectra at $\\sim$7~keV. In 2000, the source was not detected.  At 5~kpc, the 1--10~keV luminosities were $1.4 \\times 10^{33}$, $2.3 \\times 10^{34}$, and $<$$2.5 \\times 10^{33}$~\\ergs, in 1998, 1999, and 2000, respectively. These results are summarized in Table \\ref{tab:sumxobs}.  \\begin{table} \\caption{Summary of quiescent X-ray observations of \\src. The columns indicate respectively the observatory (SAX for \\sax, XMM for XMM-Newton), the year of observation, the Hydrogen column density derived, the 1--10~keV luminosity at 5~kpc, and the references ([1] for Orlandini et al. (2000), [2] for Parmar et al. (2000), [3] for this work). } \\begin{tabular}{lllll} \\hline \\hline Obs. & Year & \\nh &  L$_{\\rm 1-10\\; keV}$  & Ref.\\\\ & &  (atom cm$^{-2}$) &  (\\ergs)  & \\\\ \\hline SAX & 1998 & $(1 ^{+5} _{-1}) \\times  10^{21}$ & $1.4 \\times 10^{33}$ & [1], [2]  \\\\ SAX & 1999 & $(4.0~\\pm~0.8) \\times  10^{23} $& $2.3 \\times 10^{34}$ & [2] \\\\ SAX & 2000 & & $<$$2.5 \\times 10^{33}$& [2] \\\\ XMM & 2001 & $(5^{+3}_{-2}) \\times  10^{23}$& $3.3 \\times 10^{33}$ & [3]\\\\ \\hline \\hline \\end{tabular} \\label{tab:sumxobs} \\end{table}  Here, we report on the XMM-Newton observation of \\src\\ in quiescence performed on 2001 August 19. We present and discuss the nature of the X-ray spectrum and derive the V magnitude of the source using the Optical Monitor. ",
        "conclusions": "\\subsection{The V magnitude}   Optical observations of \\src\\ (\\cicam) exist prior to the 1998 outburst.  A V magnitude of 11.4 is attributed to \\cicam\\ during observations made in the 1970's \\citep{allen76aa}.  Long term observations of the source were carried out between 1989 and 1994. \\cicam\\ covered a range of V magnitudes between $11.44 \\pm 0.03$ and $11.76 \\pm 0.03$ (47 data points), thus showing a $\\sim$0.4~magnitude amplitude variability on a day-to-day timescale \\citep{bergner95aa}.  The mean V magnitude of \\cicam\\ during this pre-outburst interval was $11.634 \\pm 0.004$.  Performing a Fourier analysis on these data, \\citet{0421:miroshnichenko95aat} found evidence for a 11.7~days quasi-period which he interpreted as a possible orbital period.   \\citet{0421:clark00aa} report on observations of \\cicam\\ performed after the outburst, between 1998 August and 1999 March.  The V magnitude of \\cicam\\ during this post-outburst interval ranges between $11.52 \\pm 0.04$ and $11.70 \\pm 0.01$ (18 data points) and shows the same $\\sim$0.4~magnitude variability with the same mean value as during the pre-outburst period.   Thus, the V magnitude of $11.86 \\pm 0.03$ during the XMM-Newton observations is outside the range of magnitudes reported previously, indicating that the source had become fainter.  However, this magnitude represents only one data point.  So, it is difficult to determine if this higher value is due to variability on a day-to-day timescale, or if the source has become fainter for an extended interval.  \\subsection{The X-ray emission}  The quiescent XMM-Newton spectrum of \\src\\ can be fit by a power-law plus absorbed power-law and Gaussian model or alternatively by a partially covered power-law and Gaussian model. In both models, the high-energy component corresponds to a strongly absorbed continuum (${\\rm N_H}$ of $ 5 _{-2}^{+3}$~\\ttroisnh).  This fitting approach is supported by the fact that a similarly strongly absorbed quiescent emission was also reported from the source during the 1999 \\sax\\ observation.  \\citet{0421:parmar00aa} fit the \\sax\\ 2--10~keV spectrum with an absorbed (\\nh\\ of $(4.0 \\pm 0.8)$~\\ttroisnh) power-law (\\phind~$= 1.86 ^{+0.27}_{-0.32}$) together with a Gaussian emission feature with an energy of 7.3~$\\pm$~0.2~keV and an equivalent width of 620~$\\pm$~350~eV. Since an unabsorbed component is unambiguously detected at low-energy in the XMM-Newton quiescent spectrum of \\src, we re-examined the \\sax\\ spectrum reported in \\citet{0421:parmar00aa} to see if there is evidence for the presence of a similar component. The 0.2--2~keV LECS count rate of (7.5~$\\pm$~4.8)~$\\times$~10$^{-4}$~\\persec\\ suggests that such a component may be present. In order to investigate this further, we fit the partially covered XMM-Newton model discussed above to the 0.5--10~keV \\sax\\ spectrum allowing the spectral parameters to vary. The line energy and width which were poorly constrained were set to the best-fit XMM-Newton values.  This gives a \\rchisq\\ of 1.43 for 38 d.o.f.. The uncovered component contributes (6$\\pm$3)\\% of the total 1--10~keV absorbed flux, consistent with the ratio observed with XMM-Newton.  Thus, the presence in the \\sax\\ 1999 spectrum of an unabsorbed component, similar to that detected in the XMM-Newton spectrum cannot be excluded.    When absorption is added to either the power-law plus absorbed power-law and Gaussian model or the partially covered power-law and Gaussian model used to fit the XMM-Newton spectrum, the \\nh\\ of this component is $<$2.0~\\tunnh\\ (see Table~\\ref{tab:resultscomb}). This value is consistent with that obtained when \\src\\ approached quiescence after the outburst \\citep{0421:belloni99apj}.  These low values of \\nh\\ confirm that the {\\it interstellar} column towards \\src\\ is not high.  On the other hand, X-ray results show that the column density intrinsic to the system can be very high, and as inferred from the large range of absorption obtained (from roughly 0.2 to 50~\\ttnh), very variable. Thus, this confirms the picture that most of the absorption towards \\src\\ is local and not interstellar.  The Gaussian feature observed at 6.4~keV can be interpreted as a fluorescent emission from iron. Such an interpretation is consistent with the modeling of the spectrum using partial covering since it suggests the presence of significant cold absorbing material in the system. The large equivalent width observed can be explained if cold material surrounds the X-ray emitter with a large column density, which is also consistent with our modeling and with the picture, drawn from optical spectroscopy, of a compact object embedded in a very dense wind \\citep{0421:robinson02apj}.  However, the spectral quality is too low to test for the presence of associated signatures of cold material such as absorption edges. A reflection component due to the presence of cold material could be expected as well, but such a component usually peaks above the energy range covered here (between $\\sim$10 and 100~keV). Gaussian emission features were detected during the 1998 and 1999 \\sax\\ observations of \\src\\ in quiescence at energies of $7.0^{+1.6}_{-0.2}$ and 7.3~$\\pm$~0.2~keV respectively. Their different energies as compared to the 6.4~keV feature detected with XMM-Newton may indicate a different origin, or different physical conditions in the emission region.  Emission features at $\\sim$6.4~keV were also detected in outburst spectra from \\src\\ and mostly interpreted as emission lines produced by an optically thin plasma \\citep[see e.g.,][]{0421:ueda98apjl,0421:revnivtsev99al}. Such a mechanism can not be excluded in the case of the XMM-Newton observation of \\src, although the geometry and emission processes involved during quiescence are likely to be very different from those involved during the outburst.   The partially covered power-law  and Gaussian model as well as the powerlaw plus absorbed power-law and Gaussian model suggest the presence of two components: an unabsorbed component mainly observed $\\lesssim$5~keV (the uncovered or low-energy component), and second, a strongly absorbed component mainly observed $\\gtrsim$5~keV (the covered or high-energy component) and dominating the total spectrum.   We propose that the covered component results from the compact object being embedded within the dense circumstellar wind emitted from the sgB[e] companion star, in agreement with the picture drawn from optical spectroscopy of the source \\citep{0421:robinson02apj}.  The large range of observed \\nh\\ at X-ray wavelengths could reflect the complexity of the B[e] star environment in which the compact object is traveling. Regions with different physical properties may be crossed, depending e.g., on the distance of the compact object from the sgB[e] star or from its equatorial plane. This environment may vary with time as well. It may also be modified by the X-rays emitted in the vicinity of the compact object, which are probably variable themselves.  We suggest two possible origins for the low-energy component.  First, it could be due to X-ray emission from the sgB[e] star itself.  The X-ray emission from OB stars is intrinsically soft \\citep[up to $\\sim$4~keV,][]{long80apjl}.  \\citet{0421:orlandini00aa} estimate that the X-ray luminosity of the companion star in \\src\\ could be $\\sim$5~$\\times 10^{32}$~\\ergs, while \\citet[Sect. 2.4]{0421:robinson02apj} estimate that the sgB[e] star could emit up to $10^{34}$~\\ergs\\ in the 0.2--4.0~keV band.  The 0.2--4.0~keV luminosity observed from \\src\\ during the XMM-Newton observation is $1.5 \\times 10^{33}$~\\ergs\\ at 5~kpc. Thus, we cannot exclude that the low-energy emission, or part of it, originates from the companion star.  Another possibility is that the low-energy component is due to the scattering of higher-energy X-ray photons  in a wind or ionized corona such as observed in some low-mass X-ray binaries.  The flux of the low-energy component in \\src\\ is about 10\\% of the total 0.5--10~keV flux.  In dipping, eclipsing or accretion disk corona sources, the ratio observed between the flux attributed to scattered emission and the total flux is usually $\\sim$5\\% \\citep[see e.g.,][]{0748:parmar86apj}. Thus, at least a part of the low-energy emission could be due to scattering in \\src.   We note however that corona have been observed in low-mass X-ray binaries that are much brighter than \\src. So the possible scattering region in \\src\\ may differ in nature and formation from those observed in low-mass X-ray binaries. The scattering region in \\src\\ could be linked to the wind emitted by the B[e] companion star.  Emission from the companion star and scattering could both play a role in the low-energy emission observed from \\src.   \\src\\ is the first identified member of a new class of HMXB with sgB[e] companion. It is the only known system in which the compact object is immersed in a dense and complex circumstellar wind.  Further multiwavelength observations of this source are needed to explore the geometry and the emission processes involved in this system.  Many other stars showing the B[e] phenomenon, and especially sgB[e] stars, could host a compact object. Due to their low X-ray luminosity and absorbed spectra, such objects are unlikely to have been identified in previous low-energy (0.1--2.5~keV) sky surveys such as conducted by ROSAT, and we await future medium energy X-ray surveys to detect further members of this class."
    },
    "0207/astro-ph0207293_arXiv.txt": {
        "abstract": "We review the current status of cosmological parameters, dark energy and large-scale structure, from a theoretical and observational perspective. We first present the basic cosmological parameters and discuss how they are measured with different observational techniques. We then describe the recent evidence for dark energy from Type Ia supernovae. Dynamical models of the dark energy, quintessence, are then described, as well as how they relate to theories of gravity and particle physics. The basic theory of structure formation via gravitational instability is then reviewed. Finally, we describe new observational probes of the large-structure of the universe, and how they constrain cosmological parameters. ",
        "introduction": "In what is now a classic story, lies the foundation of 21st century cosmology.  In 1917 Einstein formulated the field equations which suggested a dynamic universe, in contradiction to the available data which supported a static universe.  Einstein thus added his now famous ``fudge factor'', the cosmological constant.  Following Edwin Hubble's 1926 paper in which he suggested an expanding universe, Einstein removed the cosmological constant.  The subsequent history of cosmology has leapfrogged between experiments to measure the cosmological parameters (expansion rate, mass density, deceleration rate, geometry and now energy density) and theories to organize them.  However, the last few years have been very exciting for cosmologists. Cosmic microwave background experiments, observations of Type Ia supernovae, large scale surveys of galaxies and clusters and the Hubble key project have yielded unhitherto precise measurements of the key cosmological parameters.  Furthermore, these new values have provided a concordance model, a flat universe with a significant component of ``dark energy'' which accelerates the expansion in the current epoch, which challenges our fundamental theories of particle physics.  In this report, we  summarize the current status of the cosmological parameters (the Hubble constant, mass and vacuum energy density (dark energy)), quintessence and large scale structure; and their import on modern cosmology. ",
        "conclusions": ""
    },
    "0207/astro-ph0207400_arXiv.txt": {
        "abstract": "In many GRB inner engine models the  highly  relativistic GRB jets are engulfed by slower moving matter. This could result in different beaming for  the prompt $\\gamma$-ray emission and for the lower energy afterglow. In this case we should expect that some observer will see {\\it on-axis orphan afterglows}: X-ray, optical and radio afterglows within the initial relativistic ejecta with no preceding GRB; The  prompt $\\gamma$-ray emission is pointing elsewhere. We show that the observations of the WFC on BeppoSAX constrain with high certainty the prompt X-ray beaming factor to be less than twice the prompt $\\gamma$-ray beaming. The results of Ariel 5 are consistent with this interpretation.  The RASS from ROTSE and HEAO-1 constrain the X-ray beaming factor  at 400 and 20 minutes after the burst respectively to be comparable and certainly not much larger than the $\\gamma$-ray beaming factor. There is no direct limit on the optical beaming. However, we show that observations of several months with existing hardware could result in a useful limit on the optical beaming factor of GRB afterglows. ",
        "introduction": "The realization that Gamma-Ray bursts (GRBs) are beamed changed our understanding of the phenomenon in many ways. The first and most dramatic  is the drastic revision of  the bursts' energy. The enormous $ 10^{54}$ergs, turned out, after beaming corrections,  to be a ``modest\"  $ 10^{51}$ergs. Naturally the actual GRB rate increased by the inverse factor. Even more surprising was the discovery (Frail et al., 01; Panaitescu \\& Kumar, 01; Piran et al., 01) of the rather narrow total energy distribution. Currently, the evidence for beaming is the 'jet break' in the afterglow's light curve. The jet angle $\\theta_j$ is determined from the time of the break in the light curve, that is interpreted as a jet break (Sari, Piran \\& Halpern, 99). However, even if this interpretation is correct it corresponds to the jet opening angle when most of the emission is in the optical or IR bands.  There is no direct evidence for the beaming factor during the GRB phase, or even during the early afterglow phase when the emission is mostly in X-rays.  It is not clear whether the $ \\gamma $-ray, X-ray, optical and radio emissions have similar initial (before the jet spreading) beaming factors. There are  good physical reasons to question whether there is a common beaming factor in different wavelengths. Different parts of the spectrum dominate the emission at different times and correspond to different physical conditions within the relativistic flow. It is possible, and some will argue even likely, that emission in different energy bands will have different beaming factors.   The origin of the problem in determining the different beaming lies in two related relativistic phenomena, casual connection and relativistic beaming. Regions more than $\\Gamma^{-1}$ apart within the relativistic flow are casually disconnected. Additionally, the light emitted from a relativistic source moving with a Lorentz factor $\\Gamma$ is beamed into an angle of $\\Gamma^{-1}$ along the line of motion. The first phenomenon implies that regions that are more than $\\Gamma^{-1}$ apart could have different physical conditions. The second one implies that an observer could see only one such region at a time and won't know about the other. Both phenomena imply that the $\\gamma$-rays, that are emitted when $\\Gamma \\ge 100$, could have come from a jet with an opening angle of 1/100 and the observer would have no way of telling the difference from a spherical source. The later X-ray (as well as optical and radio) afterglow is emitted from a slower moving material with a lower Lorentz factor. Hence, the X-ray beam could be larger than the $\\gamma$-rays beam. Similarly the region observed by a given observer is larger as well (see fig \\ref{fig:OnAxisOrph}).   \\begin{figure} \\label{fig:OnAxisOrph} {\\par\\centering \\resizebox*{0.9\\columnwidth}{0.5\\textheight}{\\includegraphics{fig1.eps}} \\par} \\caption {A schematic description of on-axis orphan afterglow. Only some regions within the initial ejecta  emit prompt $\\gamma$-ray radiation. The structure of the $\\gamma$-ray emitting regions could be regular (upper left figure) or irregular  (upper right figure). The ellipses describe the area observed by an observer at a given time (The smaller ellipses describe the area observed initially during the $\\gamma$-ray emission, when $\\Gamma$ is large, while the larger one describe the area observed during the X-ray emission, when $\\Gamma$ is smaller.) The lower figure depicts the cross section of the upper left figure. Observer A detects the early emission from a small region within the $\\gamma$-ray emitting region and later  an afterglow from a much larger region. This observer would see a regular GRB and an  afterglow. Observer B does not detect any $\\gamma$-rays but detects a regular afterglow, which we call on-axis orphan afterglow. } \\end{figure}  These relativistic effects allow configurations with narrow $\\gamma$-ray beams, wider X-ray and even wider optical and radio beams. This will produce orphan afterglows - events in which X-ray, optical and radio afterglow is observed while the prompt GRB is pointing elsewhere. There are two types of orphan afterglows: (i) On-axis orphan afterglows (which we introduce and discussed here, see fig. \\ref{fig:OnAxisOrph}) are observed within the initial relativistic jet by observers that miss the narrower $\\gamma$-ray beam. These afterglows follow  the light curves of the standard afterglows (observed following regular GRBs). (ii) Off-axis orphan afterglows (see fig. 2) are the ``traditional\" orphan afterglows (Rohads, 97; Perna \\& loeb, 98; Dalal, Griest \\& Pruet 2001, Granot et al. 2002, Nakar, Piran \\& Granot 2002, Totani \\& Panaitescu 2002) that are observed outside the initial jet. Off-axis orphan afterglows can be seen only after the jet break when the jet expands sideway. Their light curve rises initially reaching a maximal flux (that depends on the observing angle) and then decays following the post-jet-break light curves of  a standard GRB afterglow. To study the initial opening angles of the relativistic jets we must consider the on-axis orphan afterglows.    \\begin{figure} \\label{Off} {\\par\\centering \\resizebox*{0.9\\columnwidth}{0.4\\textheight}{\\includegraphics{fig2.eps}} \\par} \\caption{Off Axis Orphan afterglow is seen by observers that are not within the initial relativistic jet. This emission is seen only after the jet break and the sideways expansion of the relativistic material. The schematic figure depicts three observers. Observer A detects both the GRB and the afterglow. Observer B does not detect the GRB but detects afterglow that is similar to the one observed by A. Observer C detects off-axis orphan afterglow that rises and fall and differs from the afterglow detected by observers A and B.} \\end{figure}   A direct way to determine the beaming ratios is  to compare the rates of detection of transients in different energy bands. However, several confusing factors should be taken into account in such a comparison. First, detectors in different energy bands have different relative thresholds. These should be calibrated using the current GRB and afterglow observations. Second, there are numerous background transients and we have to identify specific transients as afterglows. We show in  section 3 that this problem may not be severe for the X-ray band. Even assuming that all observed transients (after some basic filtering) are afterglows we find a tight constraint on the ratio of X-ray to $\\gamma$-ray beaming. Optical background transients (e.g. AGNs, stellar flares etc.) are more numerous. Here, we should use the temporal and spectral observations of the afterglows, observed so far, as templates for identification.  A third problem that is unique to afterglows is the possible confusion between the optical and radio\\footnote{The current X-ray observations are before the jet break, when only on-axis afterglows can be seen.} on-axis and off-axis orphan afterglows. The overall light curves of on-axis and off-axis orphan afterglows are significantly different (see Fig \\ref{fig:lightSchematic}). However, the post-jet-break light curves of both kinds of orphan afterglows are similar (Granot et al. 2002). To avoid confusion we must catch the afterglow early before the jet break. This can be done by an appropriate choice of the magnitude of an optical survey. While off-axis afterglows are more numerous, most off-axis orphan afterglows won't be detected in a shallow survey and the on-axis orphan afterglows would govern such a sample. In the radio, the transients would  be generally detected after the jet-break and off-axis orphan afterglows would always govern the sample. An estimate of the rate of transients in a radio survey could not constrain the initial radio beaming, but rather it would provide a measure of the total rate of relativistic ejection events (Perna \\& Loeb, 98; Levinson et al.,2002).  \\begin{figure}  \\label{fig:lightSchematic} {\\par\\centering \\resizebox*{0.9\\columnwidth}{0.4\\textheight}{\\includegraphics{fig3.eps}} \\par} \\caption{Schematic light curves of on-axis and off-axis orphan afterglows. Naturally, early, when the observers is outside the initial beam, off-axis afterglows are much weaker. Later, after the jet break both light curves are similar.} \\end{figure}   Our first goal in this paper is to estimate the X-ray to $\\gamma$-ray beaming factors using  the observed limits on X-ray transients. We show (in section 3) that while there is a weak indication that the $\\gamma$-rays are beamed by a factor of a few relative to the X-rays, even current observations rule out the possibility that the $\\gamma$-rays are  significantly (more than a factor of ten) beamed than the X-rays.   There is no available data on relevant optical transients. In a second part of the paper (section 4) we describe possible optical surveys that, with existing hardware, could limit the optical beaming within several months of observation. We also discuss (in section 5) the implications of observations of radio transients. We summarize our results in section 6.   The observed GRB rate is a basic reference to which we compare the rate of other transients. So far all the detected afterglows are of long bursts (duration longer than 2sec). We exclude, therefore, short GRBs when estimating the GRB rate. We take the (all sky) rate of long GRBs as 600 per year (Fishman \\& Meegan, 1995). ",
        "conclusions": "There are various reasons to expect that GRB jets will have an angle dependent profile. For example, this is  natural in the Collapsar model, in which a jet punches a hole in the surrounding stellar envelope. One can expect that along with the ultra-relativistic motion at the core of the jet there would be slower motion of a thicker envelope. The relations between the different beaming factors have many implications on different parameters of the GRB, most notably on the total energy budget. For example Frail et. al. (2001) and Panaitescu \\& Kumar (2001) calculate the total energy emitted in the prompt emission using: $ E_{\\gamma }=E_{\\gamma iso}f_{b_{opt}} $, where $ E_{\\gamma iso} $ is the isotropic energy in $ \\gamma $-rays (during the several dozen seconds of the burst) and $ f_{b_{opt}}=\\theta^2_j/2 $ is the beaming factor of the optical emission several days after the burst. Clearly, this calculation is relevant only if $ f_{b_{opt}}\\approx f_{b_{\\gamma }} $, namely if the beaming during the afterglow after a day or so reflect the beaming during the prompt $\\gamma$-ray emission. Indeed this fact that $ E_{\\gamma }=E_{\\gamma iso}f_{b_{opt}} $, is narrowly distributed suggests that $f_{b_{opt}}\\approx f_{b_{\\gamma }} $ (or at least $f_{b_{opt}}\\propto f_{b_{\\gamma }} $). Otherwise this narrow distribution would require a miraculous coincidence.   These ideas have motivated our study of the possibility of different beaming factors for different observed energies that correspond to matter moving with different Lorentz factors. Our analysis shows that the results of the three  X-ray surveys, as well as the BeppoSAX data are consistent with a very moderate difference in the $\\gamma$-rays and X-rays beaming. The four surveys shows that both the early and the late X-ray beaming, $f_{b_x}$, are comparable to the $\\gamma$-ray beaming, $f_{b_\\gamma}$. Our most stringent limit $f_{b_x} < (1.34 \\pm .23)  f_{b_\\gamma}$, is obtained  for the prompt X-ray emission from the BeppoSAX WFCs. This is only an upper limit because of the uncertainty in the interpretation of the XRFs as orphan GRB counterparts. The ratio is even lower if XRFs are not related to GRBs.   The highest limit is obtained for the afterglow 400 minutes after the burst from the ROSAT RASS data. With 17 unidentified transients compared to 3 expected we find $f_{b_x} < 8 f_{b_\\gamma}$. However, this is only a weak upper limit as all 6 transients  (out of the original 23 transients) that have been checked  were found to be flaring stars. It is possible that all the remaining 17 transients are not related to GRBs. We conclude that during the X-ray afterglow (20-400 minutes after the burst) $f_{b_x}$ is probably less than twice $f_{b_\\gamma}$ and $f_{b_x} \\approx 10 f_{b_\\gamma}$ is certainly ruled out by current observations. This result shows that the bulk energies at $\\Gamma=200$ and at $\\Gamma =10$ are comparable. This result supports the homogeneous jet approximation. It puts a strong constraint on the $\\theta$ dependent jet model of Rossi et al., (02), as it requires that the ratio of $\\gamma$-ray to X-ray emission is roughly constant throughout this variable jet.  In our analysis we have implicitly assumed (by using Eq. \\ref{xrays}) that the energy per unit solid angle is similar in the on-axis orphan afterglows to the one in regular GRB afterglows. If this energy is lower by a given factor $\\varepsilon_x/\\varepsilon$, then using Eq. \\ref{Fx} we find that $t_{max}$ will decrease by a similar factor. The signal will be detected but for a shorter period, as long as the maximal X-ray flux is above the sensitivity limit of the survey. This, in turn, will decrease the detection rate by $\\varepsilon_x/\\varepsilon$ and will increase the implied limits on the ratio of $\\gamma$-ray beaming to X-ray beaming by $(\\varepsilon_x/\\varepsilon)^{-1}$. The total energy in X-ray emitting matter is proportional to $f_{bx} \\varepsilon_x$. Hence our results impose a direct limit (independent of $\\varepsilon_x$) on this quantity. They show that the total energy in X-ray cannot be significantly larger than the total energy in $\\gamma$-ray emitting matter.   It will be a remarkable achievement to constrain also the optical ($\\Gamma \\sim 2-5$) beaming. We have argued that even existing hardware like Super-LOTIS can rule out (or confirm) $f_{b_{opt}}\\approx 10 f_{b_\\gamma}$ on a time scale of several months. The proposed LTSS can constrain this factor within a week. Radio observations provide a different limit on the total rate of relativistic ejecta events. However, this limit can be obtained only under the assumption of a single beaming factor. If optical afterglow observations would show that $f_{b_{opt}} \\approx f_{b_{\\gamma}}$ this assumption would be reasonable. Under this assumption the comparison of the radio orphan afterglow rate to the GRB rate can set limits on $f_b$. Levinson et al (02) find that current surveys limit $f^{-1}_b > 10$. This is consistent with the rest of the observational results.   One argument that supports a beaming ratio of a few is obtained by comparing the different energies during the GRB. A detailed analysis of the afterglow light curves (Panaiteschu \\& Kumar, 01) and of the X-ray fluxes (Piran et al., 01)  have shown that $E_K$ the kinetic energy during the adiabatic afterglow phase is also narrowly distributed. Panaiteschu \\& Kumar (01) find  using $ E_{\\gamma }=E_{\\gamma iso}f_{b_{opt}} $ and their estimates of $f_{b_{opt}}$ that the average $E_\\gamma$ is larger than the average $E_K$: $\\bar E_\\gamma \\approx 3 \\bar E_K$. This result is somewhat puzzling (Piran, 01): A narrow $E_\\gamma$ distribution with $\\bar E_\\gamma > \\bar E_K$ implies (with no fine tuning) that  $E_{rel}$ (the total energy of relativistic ejecta) is narrowly distributed and $E_\\gamma \\approx E_{rel}$ (rather than $E_{rel} \\approx E_K$). The fact that $E_K<E_\\gamma$ is also narrowly distributed implies that $\\epsilon_{\\gamma}$, the conversion efficiency of relativistic kinetic energy to $\\gamma$-rays ($E_{\\gamma} \\equiv \\epsilon_{\\gamma} E_{rel}$), is close to unity and moreover, $\\epsilon_{\\gamma}$ itself should be rather narrowly distributed (between 70-80\\%). This last conclusion would be astonishing considering the dependence of $\\epsilon_{\\gamma}$ on the the distribution of energies and Lorentz factors of the different shells that produce the internal shocks. This puzzle can be resolved if according to the patchy shell model, the $\\gamma$-ray emitting regions are indeed narrower than the late time afterglow emitting regions. In this case the quantity: $E_{\\gamma iso}f_{b_{opt}} $ is an overestimate of $E_{\\gamma}$ by the relative beaming factor. Specifically, $E_{\\gamma}$ would have been overestimated by a factor of 3 if  $f_\\gamma \\approx f_{opt}/3$ (Piran, 01).  Before concluding we remark that our results also impose limits on the rate of  ``failed GRBs\".  We define here ``failed GRBs\" as events in which the source fails to produce GRB but ejects a relativistic ejecta (comparable in energy to GRBs' ejecta) that produces a detectable afterglow. From the point of view of this paper ``failed GRBs\" can be viewed as extreme cases of on-axis orphan afterglows in which the $\\gamma$-ray beam vanishes. In these events our ``on axis orphan afterglows\" would be true orphans having no GRB parent pointing towards any observer. The implied rates that we have found provide immediately limits on the ratio of ``failed GRBs\" to GRBs. In particular the ROSAT survey shows that the rate of ``failed GRBs\" with a regular X-ray afterglow cannot be more than a factor of eight (and possibly much lower) than the rate of GRBs. If the energy per solid angle is smaller, then as discussed earlier, the limit on the rate would be higher by the inverse factor.  Optical survey, of the type we discuss here, could limit the rate of ``failed GRBs\" with an optical afterglow. This would presumably correspond to events in which the matter is ejected with a lower Lorentz factor.   We thank Avishay Gal-Yam, Jonathan Granot, Jonathan Katz, Pawan Kumar and Luigi Piro for helpful discussions. This research was supported by the US-Israel BSF."
    },
    "0207/astro-ph0207636_arXiv.txt": {
        "abstract": "name{Samenvatting}% \\def\\bibname{Bibliografie}\\def\\chaptername{Hoofdstuk}% \\def\\appendixname{Bijlage}\\def\\contentsname{Inhoudsopgave}% \\def\\listfigurename{Lijst van figuren}\\def\\listtablename{Lijst van tabellen}% \\def\\indexname{Index}\\def\\figurename{Figuur}\\def\\tablename{Tabel}% \\def\\partname{Deel}\\def\\enclname{Bijlage(n)}\\def\\ccname{Ter attentie van}% \\def\\headtoname{Aan}\\def\\headpagename{Pagina}% \\def\\today{\\number\\day\\space\\ifcase\\month\\or januari\\or februari\\or maart\\or% april\\or mei\\or juni\\or juli\\or augustus\\or september\\or oktober\\or% november\\or december\\fi \\space\\number\\year}% \\typeout{ >>>>> use hlatex209 for Dutch hyphenation <<<<< }} \\hyphenation{Schrij-ver Krij-ger Kuij-pers Bal-le-gooij-en}   \\def\\warningoverprint #1{\\special{!userdict begin /bop-hook{gsave 100 600 translate -45 rotate /Times-Roman findfont 100 scalefont setfont 0 0 moveto 0.9 setgray (#1) show grestore}def end}}  \\newcounter{onefig} \\newcounter{fignumber} \\newcount\\nocaptions \\newcount\\nofigures \\newcount\\figwidth \\newcount\\viewgraphs \\def\\figsfile{} \\def\\figspath{} \\def\\paper{} \\def\\format{} \\def\\figlabel{} \\long\\def\\nextfig#1{\\setcounter{figure}{\\value{fignumber}} \\addtocounter{fignumber}{1} \\ifnum \\viewgraphs=1 \\newpage \\pagestyle{empty} \\fi \\ifnum\\value{onefig}=0 #1 \\fi \\ifnum\\value{onefig}=\\value{fignumber} #1 \\fi} \\def\\figwidths#1#2{\\ifnum \\nocaptions=1 #2mm \\else #1mm \\fi} \\def\\paper#1{}  % \\long\\def\\plotfig#1#2{\\ifnum \\nofigures=1 \\else #2 \\fi} \\long\\def\\captiontext#1{\\ifnum \\nofigures=1 \\raggedright \\fi \\ifnum \\nocaptions=1 \\paper \\ifnum \\viewgraphs=0 \\newline  \\mbox{}\\hrulefill\\mbox{} \\newline \\newline label:~\\{\\figlabel\\} \\fi \\else \\ifnum \\nofigures=0 \\fi #1 \\fi}  \\newcount\\panelwidth \\newcount\\panelheight \\newcount\\bxmin \\newcount\\bymin \\newcount\\bxmax \\newcount\\bymax \\newcount\\tbxmin \\newcount\\tbymin \\newcount\\tpanelwidth \\newcount\\tpanelheight \\newcount\\tpdif \\panelwidth=70 \\panelheight=70  % \\def\\panelsize #1,#2;{\\panelwidth=#1 \\panelheight=#2} \\def\\setbb #1,#2;#3,#4;#5,#6;{% \\tbxmin=#1 \\tbymin=#2    % \\bxmin=#3 \\bymin=#4      % \\bxmax=#5 \\bymax=#6}     % \\def\\barepanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,clip=,% width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}} \\def\\labelxpanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelxypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}}  \\newcommand{\\panellabel}[3]{\\makebox[0pt][r]{\\raisebox{#2} #3 \\hspace{#1}}}  \\def\\topfraction{0.99}                      % \\def\\dbltopfraction{0.99} \\def\\textfraction{0.01} \\def\\floatpagefraction{0.99}  \\def\\CC{\\par \\vspace*{-2ex} \\footnotesize \\baselineskip=8pt \\begin{verbatim}} \\def\\EE{\\par \\vspace*{-2ex} \\normalsize}  \\long\\def\\startignore #1\\stopignore{}   %   \\newenvironment{task}{% \\begin{list}{$\\bullet$}{\\itemsep=1ex \\parindent=3ex \\parskip=1ex \\topsep=0ex}} {\\end{list}}  \\def\\koprr{\\list{}{ \\itemindent=0pt \\listparindent=0pt \\itemsep=0pt \\leftmargin=5ex}\\item[]}                     % \\let\\endkoprr=\\endlist  \\def\\kolom{\\list{}{\\rightmargin=5cm      % \\topsep=0ex \\itemsep=0cm \\parsep=0ex \\parskip=0ex \\partopsep=0ex \\listparindent=3ex \\itemindent=0pt \\leftmargin=0pt} \\item[] \\vspace*{-0ex}}                        % \\let\\endkolom=\\endlist  \\renewcommand{\\labelitemi}{{\\bf --}}                % \\def\\setlistparams{ \\topsep=0.7ex                 % \\itemsep=0.7ex                % \\leftmargini=3ex}             % \\setlistparams                  %  \\newcounter{alistindex}  \\newcounter{alistvalue}     % \\newenvironment{alist}{% \\begin{list}{\\alph{alistindex})}{\\usecounter{alistindex}} \\itemsep=1.5ex \\parskip=1ex}{\\end{list}}  \\def\\spacing #1{\\small \\renewcommand{\\baselinestretch}{#1} \\normalsize \\setlistparams}             %  \\newenvironment{itemrr}{\\vspace*{-\\parskip} \\leftmargini=3ex \\topsep=0.6ex \\begin{itemize} \\itemsep=0.5ex \\parindent=0ex \\parsep=0ex}{ \\end{itemize} \\vspace*{-\\parskip}}  \\newenvironment{conditem}{ \\leftmargini=6ex \\topsep=0.2ex \\begin{itemize} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\newenvironment{enumerr}{ \\leftmargini=3ex \\topsep=1.5ex \\begin{enumerate} \\itemsep=0.7ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newenvironment{condenum}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{enumerate} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newcounter{romenumnr} \\newenvironment{romanenum}{ \\leftmargini=5ex \\topsep=1.5ex \\begin{list}{(\\roman{romenumnr}) --}{\\usecounter{romenumnr} \\itemsep=1ex \\listparindent=3ex \\parsep=0ex}}{ \\end{list}}  \\newcommand{\\posterpage}[2]{{\\Large \\sf \\newpage \\setlength{\\minipagewidth}{\\textwidth} \\addtolength{\\minipagewidth}{-2cm} \\fboxrule=0.3ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[\\textwidth]{ \\begin{minipage}[t]{\\minipagewidth} \\parskip=3ex \\vspace*{1ex} \\begin{center} \\LARGE \\bf #1 \\end{center} \\vspace{2ex} #2 \\vspace*{3ex} \\mbox{ } \\end{minipage}} \\hspace*{\\textwidth} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.2ex} \\framebox[\\textwidth]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{\\textwidth}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}}  \\newcommand{\\vwgraph}[3]{\\newpage \\begin{center} \\shadowframe{#1}{ \\begin{center} \\vspace*{2ex} {\\Large \\bf #2}\\\\[4ex] \\end{center} {\\Large \\sf #3}}   % \\end{center}} \\newenvironment{itemvg}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{itemize} \\vspace{-1ex} \\itemsep=0.5ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\long\\def\\lectvw#1#2{\\newpage \\vspace*{-15mm} \\begin{center} \\begin{center} {\\LARGE \\sc \\underline{#1}}\\\\[4ex] \\end{center} {\\Large \\sf #2}   % \\end{center}}  \\newlength{\\minipagewidth} \\newcommand{\\shadowframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{1.0ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{4.5ex} \\rule{#1}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newsavebox{\\boxcontent} \\newcommand{\\ovalhead}[1]{ \\unitlength=1cm \\sbox{\\boxcontent}{\\mbox{~~{#1}~~}} \\begin{center} \\ifdim\\wd\\boxcontent>6ex \\ifdim\\wd\\boxcontent<8cm \\begin{picture}(8,3) \\thicklines \\put(4.0,0.8){\\oval(8,1.6)} \\put(0.0,0.7){\\parbox{8cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\ifdim\\wd\\boxcontent<12cm \\begin{picture}(12,3) \\thicklines \\put(6.0,0.8){\\oval(12,1.6)} \\put(0.0,0.7){\\parbox{12cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\begin{picture}(16,3) \\thicklines \\put(8.0,0.8){\\oval(16,1.6)} \\put(0.0,0.7){\\parbox{16cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\fi \\fi \\fi \\end{center}}  \\newcommand{\\colloqframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\parskip=3ex \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{#1}{2ex} \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newcommand{\\soloval}[1]{\\fboxrule=0.5mm \\unitlength=1cm \\hspace*{5mm} \\sol{20mm}\\\\[-20mm] \\begin{picture}(15,3) \\thicklines \\put(7.5,2){\\oval(15.60,3.60)} \\put(7.5,2){\\oval(15.55,3.55)} \\put(7.5,2){\\oval(15.50,3.50)} \\put(7.5,2){\\oval(15.45,3.45)} \\put(7.5,2){\\oval(15.40,3.40)} \\put(7.5,2){\\oval(15.35,3.35)} \\put(7.5,2){\\oval(15.00,3.00)} \\put(2.0,1.9){\\parbox{13cm}{#1}} \\end{picture}}  \\setcounter{secnumdepth}{3} \\setcounter{tocdepth}{3}  \\def\\chapterrr#1{\\chapter{#1} \\label{chap:#1} \\vskip-\\parskip} \\def{ Vega, Sirius, $\\beta$ Leo, $\\alpha$ Car and $\\alpha$ Cen A belong to a  sample of twenty stellar sources used for the calibration of the detectors of the Short-Wavelength Spectrometer on board the Infrared Space Observatory (ISO-SWS). While general problems with the calibration and with the theoretical modelling of these stars are reported in \\citet{Decin2000b}, each of these stars is discussed individually in this paper. As demonstrated in \\citet{Decin2000b}, it is not possible to deduce the effective temperature, the gravity and the chemical composition from the ISO-SWS spectra of these stars. But since ISO-SWS is absolutely calibrated, the angular diameter (\\ad) of these stellar sources can be deduced from their ISO-SWS spectra, which consequently yields the stellar radius (R), the gravity-inferred mass (M$_g$) and the luminosity (L) for these stars. For Vega, we obtained \\ad$ = 3.35 \\pm 0.20$\\,mas, R$ = 2.79 \\pm 0.17$\\,\\Rsun, M$_g = 2.54 \\pm 1.21$\\,\\Msun\\ and L$ = 61 \\pm 9$\\,\\Lsun; for Sirius \\ad$ = 6.17 \\pm 0.38$\\,mas, R$ = 1.75 \\pm 0.11$\\,\\Rsun, M$_g = 2.22 \\pm 1.06$\\,\\Msun\\ and L$ = 29 \\pm 6$\\,\\Lsun; for $\\beta$ Leo \\ad$ = 1.47 \\pm 0.09$\\,mas, R$ = 1.75 \\pm 0.11$\\,\\Rsun, M$_g = 1.78 \\pm 0.46$\\,\\Msun\\ and L$ = 15 \\pm 2$\\,\\Lsun; for $\\alpha$ Car \\ad$ = 7.22 \\pm 0.42$\\,mas, R$ = 74.39 \\pm 5.76$\\,\\Rsun, M$_g = 12.80^{+24.95}_{-6.35}$\\,\\Msun\\ and L$ = 14573 \\pm 2268$\\,\\Lsun\\ and for $\\alpha$ Cen A \\ad$ = 8.80 \\pm 0.51$\\,mas, R$ = 1.27 \\pm 0.08$\\,\\Rsun, M$_g = 1.35 \\pm 0.22$\\,\\Msun\\ and L$ = 1.7 \\pm 0.2$\\,\\Lsun. These deduced parameters are confronted with other published values and the goodness-of-fit between observed ISO-SWS data and the corresponding synthetic spectrum is discussed. ",
        "introduction": "In the first two papers of this series \\citep[][ hereafter referred to as Paper~I and Paper~II respectively]{Decin2000A&A...364..137D, Decin2000b}, a method is described to analyze a sample of ISO-SWS spectra of standard stars in a consistent way. We did not only concentrate on the possibility to extract reliable stellar parameters from the medium-resolution ISO-SWS spectra, but have also demonstrated where problems in the computation of synthetic spectra --- based on the MARCS and Turbospectrum code \\citep{Gustafsson1975A&A....42..407G, Plez1992A&A...256..551P, Plez1993ApJ...418..812P}, version May 1998 --- and in the calibration of the ISO-SWS detectors destroy the goodness-of-fit between observed and synthetic spectra (Paper~II). These general results were based on a sample of 5 warm (\\teff $>$ T$_{\\mathrm{eff},\\odot}$) and 11 cool stars. In this paper, we will further analyse these 5 warm stars --- $\\alpha$ Cen A, $\\beta$ Leo, $\\alpha$ Car, Sirius and Vega ---  in order to extract relevant astrophysical data.  After a description of the general problems for these warm stars in Sect.\\ \\ref{sumprob} (as described in Paper~II), we will outline the method of analysis to deduce different stellar parameters in Sect.\\ \\ref{stelparameters} (based on the results of Paper~I and Paper~II). In the different subsections of Sect.\\ \\ref{stelparameters}, each star will be discussed individually. In order to assess the observed accuracy, some specific calibration details will be given. If available, different AOT01 observations{\\footnote{Each observation is determined uniquely by its observation number (8 digits), in which the first three digits represent the revolution number. The observing data can be calculated from the revolution number which is the number of days after 17 November 1995.}}  (i.e.\\ a full SWS scan at reduced spectral resolution, with four possible scan speeds) are compared with each other to demonstrate the calibration precision of ISO-SWS. With these remarks in mind, the synthetic spectrum based on assumed and deduced parameters is confronted with the ISO-SWS spectrum. Furthermore, we will discuss why we have assumed certain parameters and we will confront the deduced parameters from the ISO-SWS spectra with other literature values.   The appendix of this article is published electronically. Most of the grey-scale plots in the article are printed in colour in the appendix, in order to better distuingish the different spectra. ",
        "conclusions": "The five warmest stars in a sample of 16 stars --- used for the calibration of the detectors of ISO-SWS --- have been discussed spectroscopically. The absence of molecular features and the presence of atomic features whose oscillator strengths are not well-known rendered the determination of the effective temperature, gravity, microturbulent velocity, metallicity and the abundance of C, N, and O from the ISO-SWS data unfeasible. Good-quality published values were then used for the computation of the synthetic spectra. In general, no more discrepancies than the ones reported in \\citet{Decin2000b} have been detected. A comparison with other --- lower resolution --- ISO-SWS data revealed a rather good relative agreement ($\\sim 2$\\,\\%), but the absolute flux-level and so the deduced angular diameter could differ by up to 16\\,\\%. Nevertheless, the angular diameter, luminosity and stellar radius deduced from the ISO-SWS data are in good agreement with other published values deduced from other data and/or methods.  Since this research has shown clearly that the available oscillator strengths of atomic transitions in the infrared are at the moment still very inaccurate, one of us (J.\\ S.) has worked on a new atomic linelist by deducing new oscillator strengths from the high-resolution ATMOS spectrum of the Sun (625 -- 4800 cm$^{-1}$) \\citep{Sauval2000}. This new atomic linelist will be presented in Paper~V of this series."
    },
    "0207/astro-ph0207150_arXiv.txt": {
        "abstract": "{ The equations of state for degenerate electron and neutron gases are studied in the presence of magnetic fields. After including quantum effects to study the structural properties of these systems, it is found that some hypermagnetized stars can be unstable based on the criterium of stability of pressures. Highly magnetized white dwarfs should collapse producing a supernova type Ia, whilst superstrongly magnetized neutron stars cannot stand their own magnetic field and must implode, too. A comparison of our results with a set of the available observational data of some compact stars is also presented, and the agreement between this theory and observations is verified. ",
        "introduction": "Many stellar objects are known to be endowed with large magnetic fields~(Dryzek et al.~\\cite{dryzek02}, Heyl~\\cite{heyl99}, \\cite{heyl00}). For instance, white dwarfs with surface magnetic fields whose strengths range from $10^{5}$ to $10^{9}$~G have been discovered (see an earlier set of references in Kemp et al.~\\cite{kemp}, Putney~\\cite{putney95}, Schmidt \\& Smith~\\cite{schmidt95}, Reimers~\\cite{reimers96}, and the updated list of Suh \\& Mathews~\\cite{mathews}). Moreover, magnetic field strengths of the order of $10^{20}$~G have been suggested to exist in the core of neutron stars or pulsars. Recently, in Refs. Chaichian et al.~\\cite{chaichian} and P\\'{e}rez et al.~\\cite{perez}, electron and neutron gases in a strong magnetic field were considered with the aim to study the equation of state of white dwarf stars, neutron stars and its relation to supernovae. It was found a pure quantum effect: the appearance of a {\\it ferromagnetic configuration} in the neutron star interior, which is intrinsically related to the presence of the magnetic field. Such an effect opens the possibility of a quantum magnetic collapse of the gas under consideration. This effect is related to the density of the star and its magnetic field, and as such, it allows one to establish a criterium of stability taking into account these physical properties.  The presence of a magnetic field drives the loss of the rotational symmetry of the particle spectrum, which in turn manifests as an anisotropy in the thermodynamic properties of the system. This behavior can be seen through the energy-momentum tensor $\\mathcal{T}_{\\mu\\nu}$, if we recall that the external magnetic field $\\mathcal{H}$ induces a magnetization $\\mathcal{M}$ in the medium which is described through the relation $\\mathcal{H}=B-4\\pi \\mathcal{M}(B)$. Here $B$ is the microscopic magnetic field, which is assumed to point along the $x_{3}$ axis in what follows. Starting from the energy-momentum tensor, one can derive the expressions for the pressure components, longitudinal ($p_{3}$) and perpendicular ($p_{\\perp}$) to the magnetic field, \\begin{equation} \\mathcal{T}_{\\perp}=p_{\\perp}=-\\Omega-B\\mathcal{M}\\ ,\\quad\\mathcal{T} _{33}=p_{3}=-\\Omega\\ , \\end{equation} where $\\Omega$ is the thermodynamic potential.  In classical electrodynamics (Landau \\& Lifshitz~\\cite{Landau}), i.e., when the spin interactions are not taken into account, this anisotropy appears but the results are quite different. Since \\begin{equation} p_{\\perp}=p_{0}+\\frac{B^{2}}{4\\pi}\\ ,\\quad p_{3}=p_{0}-\\frac{B^{2}}{4\\pi}\\ , \\end{equation} where $p_{0}$ is the isotropic pressure, one obtains at the classical level \\begin{equation} p_{\\perp}>p_{3}\\ . \\end{equation}  This fact explains the oblateness of some astrophysical objects when one studies them in a classical way. Proper examples of this are provided by Shapiro \\& Teukolsky~\\cite{Shapiro} in the case of {\\it magnetic white dwarfs}, and by Cardall, Prakash \\& Lattimer~\\cite{CPL} for the highly magnetized neutron stars. Contrary to these well-known effects, the interplay of magnetic and quantum effects imply that some stars get a cigar-like shape along the $x_3$ axis, and some of them may even collapse.  The purpose of this paper is to discuss in more detail and to exploit some of the astrophysical consequences of previous works (Chaichian et al.~\\cite{chaichian}, P\\'{e}rez et al.~\\cite{perez}) regarding the appearance of hydrodynamic instabilities in strongly magnetized electron and neutron gases as realizations of the physics taking place inside white dwarfs (WDs) and neutron stars (NSs), respectively. In particular, these instabilities on the configuration of any super critically magnetized stellar object are shown to appear due to the action of quantum-mechanical effects, as the occupation of the particle  Landau ground state, which is driven by the cooperative particle spin-magnetic field coupling. These new effects allow us to introduce a new model to show that some WDs may become ultramagnetized and may collapse in a kind of SN type Ia, even without reaching the Chandrasekhar mass limit. As a matter of consistency, we test this theoretical argument in a wide and extensively studied set of astrophysical sources and show that the conclusions drawn from our analysis are in complete agreement with present observations of compact remnant stars, while forbid the existence of some exotic configurations of them.  The structure of the paper is as follows: In Section~\\ref{electrongas} we present the physics of a magnetized electron gas and apply it to the study of the stability conditions of white dwarfs. Based on this physics, a new model for SN type Ia events is presented, which not necessarily depends upon the WD to reach its critical mass. Section~\\ref{neutrongas} considers the case of a neutron gas, as a model for NSs. We show that the so-called magnetars cannot in principle be formed based on the instability criterium of anisotropic pressures, while canonical pulsars would. Finally, in Section~\\ref{conclusions} we present our conclusions. ",
        "conclusions": "\\label{conclusions}  If a hypermagnetized neutron star could somehow be formed in a supernova explosion, the abrupt amplification of its magnetic field will drive it into collapse. Since magnetic flux is dissipated during the implosion via a process analogous to the Sun's \\textit{coronal mass ejections} (as in Malheiro et al.~\\cite{malheiro04}), the most likely outcome of such a collapse may give as a remnant a strange star of canonical magnetic field, black hole or even an exotic black string. Once again, the violent ``consumption\" of the neutron star nuclear material may trigger the emission of a gamma-ray burst of overall energy $\\sim10^{52}$~erg, inasmuch as in the quark-nova model (Ouyed et al.~\\cite{ouyed02}).  In the case of neutron stars such as the suggested magnetars, the anisotropy of pressures must naturally develop, and the condition for the collapse in the direction perpendicular to the dipole magnetic field could be satisfied (Fig.~\\ref{fig2}) for the typical values of density and magnetic field strength routinely quoted for these stars.  To summarize, we have presented a consistent theory to discuss the stability of compact remnant stars whose structure is dictated by a combination of quantum and magnetic effects. We have shown that, in general,  the theory agrees quite well with current observational data for magnetic WDs (Shapiro \\& Teukolsky~\\cite{Shapiro}), canonical pulsars (Lorimer~\\cite{lorimer99}), and the realistic relativistic models of Cardall, Prakash \\& Lattimer~\\cite{CPL}. A major outcome is the fact that some special configurations of highly magnetized white dwarfs could collapse triggering explosions similar to a SNIa. Besides, some neutron stars endowed with superstrong magnetic fields are shown to be naturally unstable, and therefore should collapse, for their quoted surface magnetic fields. This would drive a powerful supernova explosion followed by a gamma-ray burst, too."
    },
    "0207/astro-ph0207366_arXiv.txt": {
        "abstract": "WIMP direct detection experiments are just reaching the sensitivity required to detect galactic dark matter in the form of neutralinos. Data from these experiments are usually analyzed under the simplifying assumption that the Milky Way halo is an isothermal sphere with maxwellian velocity distribution. Observations and numerical simulations indicate that galaxy halos are in fact triaxial and anisotropic. Furthermore, in the cold dark matter paradigm galactic halos form via the merger of smaller subhalos, and at least some residual substructure survives.  We examine the effect of halo modeling on WIMP exclusion limits, taking into account the detector response. Triaxial and anisotropic halo models, with parameters motivated by observations and numerical simulations, lead to significant changes which are different for different experiments, while if the local WIMP distribution is dominated by small scale clumps then the exclusion limits are changed dramatically. ",
        "introduction": "Arguably the best motivated non-baryonic dark matter candidate is the neutralino (the lightest supersymmetric particle), and current direct detection experiments are just reaching the sensitivity required to probe the relevant region of parameter space~\\cite{lars}.  The most stringent exclusion limits on Weakly Interacting Massive Particles (WIMPs) in general currently come from the Edelweiss~\\cite{edelnew} and Cryogenic Dark Matter Search (CDMS) experiments~\\cite{CDMS}, with competitive constraints also having been produced by Heidelberg-Moscow~\\cite{HM} and IGEX~\\cite{IGEX}. The sensitivity to WIMPs will be improved significantly in the short term future by the continued operation of Edelweiss, and CDMS moving in a low background environment at the Soudan mine~\\cite{CDMS2}, and in the longer term by, for instance, the planned GENIUS project~\\cite{genius}.  The direct detection event rate, and its energy distribution, depend crucially on the WIMP speed distribution. Data analyzes nearly always assume a standard smooth halo model with isotropic maxwellian velocity distribution. The change in the expected signal has been calculated for various non-standard halo models of varying degrees of sophistication~\\cite{anal,newevans,uk,sikdm}. For models which are effectively close to maxwellian, while there may be a significant change in the annual modulation and angular dependence of the signal, the change in the mean (averaged over time and recoil direction) differential event rate is typically small~\\cite{anal}. Models with triaxiality or velocity anisotropy may however produce a significant change even in the mean differential event rate~\\cite{newevans,uk}. Furthermore all of the non-standard halo models which have previously been considered are essentially smooth\\footnote{An exception is Sikivie's late infall model~\\cite{sik}, which contrary to the standard picture of halo formation in CDM cosmologies (see e.g. Refs.~\\cite{mooredm,hws}) assumes axial symmetry and cold collapse.}. N-body simulations, however, produce dark matter halos which as well as being triaxial with anisotropic velocity distributions~\\cite{mooredm,hws} also contain substructure~\\cite{Nbody1}. A number of groups have recently investigated the local dark matter distribution numerically~\\cite{swf,mooredm,hws}, using different methods and reaching, to some extent, different conclusions.  Triaxiality, anisotropy and clumping in the WIMP velocity distribution could potentially have a significant effect on the WIMP direct detection signal.  Constraints (and in the future possibly best fits) calculated assuming a standard maxwellian halo could be erroneous~\\cite{uk}. On the other hand, more optimistically, it might be possible to derive useful information about the local velocity distribution, and hence the formation of the galactic halo, if WIMPs were detected~\\cite{swf,hws}. Belli et. al.~\\cite{damare} have recently reanalyzed the DAMA collaborations annual modulation signal~\\cite{dama} for a range of halo models, finding that the allowed region of WIMP mass--cross-section parameter space is significantly enlarged. This illustrates that it is important to take into account uncertainties in halo modeling when comparing exclusion limits and/or allowed regions from different experiments.   Given the importance of the local dark matter distribution for direct detection experiments we devote Sec.~\\ref{halodis} to a detailed discussion of the global properties of real and simulated dark matter halos and recent work on the local dark matter distribution~\\cite{swf,mooredm,hws}.  In Sec.~\\ref{analysis} we examine the effect of realistic halo modeling on exclusion limits. We first investigate triaxial and anisotropic halos models, with parameter choices motivated by the observations and simulations, and then, more speculatively examine the possible effects of small subhalos. ",
        "conclusions": "Rapid progress is being made in the field of WIMP direct detection, with experiments closing in on the sensitivity required to detect neutralinos, if they constitute a non-negligible fraction (greater than $10^{-4}$) of the halo density~\\cite{dgg}. Data analyzes usually assume the simplest halo model: an isothermal sphere with maxwellian velocity distribution. There is no clear justification, either observational or theoretical, for this assumption apart from simplicity. In fact numerical simulations~\\cite{Nbody1,mooredm,hws} and observations~\\cite{sackett,om2,glob} suggest that galaxy halos are triaxial and anisotropic. The local density distribution may also be non-smooth, with late accreting subhalos leading to velocity clumping~\\cite{swf,hws}. More speculatively it is even possible that the dark matter could be mainly distributed in tiny dense clumps, so that the local density distribution could be dominated by a single clump, or could even be zero~\\cite{mooredm}. It is therefore crucial to examine the effect of realistic halo modeling on the WIMP direct detection signal.  In this paper we have investigated the change in exclusion limits due to triaxiality, velocity anisotropy and small scale clumping, taking into account detector performance and using parameter values motivated by numerical simulations and observations. Triaxiality and velocity anisotropy lead to non-negligible changes in the exclusion limits, even for detectors with relatively poor energy resolution. Furthermore the changes are different for different data sets and depend on how the anisotropy is modeled. If the local WIMP distribution is dominated by small scale clumps then the local density may be zero (making it impossible to detect WIMPs) or significantly enhanced (making it easier to detect WIMPs with a given cross-section), and the exclusion limits are changed dramatically. Clearly the survival of subhalos at the solar radius is a very important issue for WIMP direct detection.    Even if the local WIMP distribution is smooth, to derive reliable constraints on WIMP parameters and compare results for different experiments a framework needs to be developed for dealing with the uncertainty in the WIMP velocity distribution; either the development of a framework for parameterizing deviations from a baseline model, or the establishment of an agreed set of benchmark models, spanning the the range of plausible WIMP velocity distributions."
    },
    "0207/gr-qc0207098_arXiv.txt": {
        "abstract": "We present the first computations of quasiequilibrium binary neutron stars with different mass components in general relativity, within the Isenberg-Wilson-Mathews approximation. We consider both cases of synchronized rotation and irrotational motion. A polytropic equation of state is used with the adiabatic index $\\gamma=2$. The computations have been performed for the following combinations of stars: $(M/R)_{\\infty,\\rm ~star~1}$ vs. $(M/R)_{\\infty,\\rm ~star~2} = 0.12 ~{\\rm vs.}~ (0.12, ~0.13, ~0.14), ~0.14 ~{\\rm vs.}~ (0.14, ~0.15, ~0.16), ~0.16 ~{\\rm vs.}~ (0.16, ~0.17, ~0.18), ~{\\rm and} ~0.18 ~{\\rm vs.}~ 0.18$, where $(M/R)_{\\infty}$ denotes the compactness parameter of infinitely separated stars of the same baryon number. It is found that for identical mass binary systems there is no turning point of the binding energy (ADM mass) before the end point of the sequence (mass shedding point) in the irrotational case, while there is one before the end point of the sequence (contact point) in the synchronized case. On the other hand, in the different mass case, the sequence ends by the tidal disruption of the less massive star (mass shedding point). It is then more difficult to find a turning point in the ADM mass. Furthermore, we find that the deformation of each star depends mainly on the orbital separation and the mass ratio and very weakly on its compactness. On the other side, the decrease of the central energy density depends on the compactness of the star and not on that of the companion. ",
        "introduction": "\\label{s:intro}  Coalescing binary neutron stars are expected to be one of the most promising sources of gravitational waves that could be detected by the ground based, kilometer size laser interferometers such as the Laser Interferometric Gravitational wave Observatory (LIGO), VIRGO, GEO600, and TAMA300 \\footnote{For TAMA300, data taking has started in summer 1999, and the first results of the data analysis have been published \\cite{Tagoshi01}.}, and also are considered as one of candidates of gamma-ray burst source \\cite{NaraPP92}.  Due to the emission of gravitational radiation, binary neutron stars decrease their orbital separations and finally merge. When discussing the evolution of the system, it is convenient to separate it into three phases. The first one is the {\\em inspiraling phase} in which the orbital separation is much larger than the neutron star radius, and the post-Newtonian (PN) expansion constitutes an excellent approximation. Recently, two groups succeeded in deriving the 3PN equation of motion of point-mass binary systems \\cite{DamourJS,BlanchetG}, and the equivalence of the results between these groups is shown in \\cite{DamourJS01,AndradeBF01}. The second stage is the {\\em intermediate phase} in which the orbital separation becomes only a few times larger than the radius of a neutron star, so that hydrodynamics as well as general relativity play an important role. In this phase, since the shrinking time of the orbital radius due to the emission of gravitational waves is still larger than the orbital period, it is possible to approximate the state as quasiequilibrium. The final stage is the {\\em merging phase} in which the two stars coalesce dynamically. As in the intermediate phase, since hydrodynamics as well as general relativity play an important role, fully relativistic hydrodynamical treatments are required in this phase which therefore pertains to the field of numerical relativity. The first successful computations of the evolution of binary neutron stars from their innermost stable circular orbits (ISCO) to black hole or massive neutron star formation have been performed by Shibata and Uryu \\cite{Shiba99,ShibaU00}. Other efforts are presented in \\cite{OoharN99,FontGIMRSSST02}.  The present article belongs to the series of works \\cite{BonazGM99a}, \\cite{GourGTMB01} (hereafter Paper I), \\cite{TanigGB01} (hereafter Paper II), \\cite{TanigG02}, devoted to the intermediate phase. This stage is interesting because we may get informations about equation of state of neutron stars through the ISCO \\cite{FaberGRT02}. Furthermore, it is important from a numerical point of view since it provides initial data for the dynamical simulation in the merging phase \\cite{Shiba99,ShibaU00}. Then numerous theoretical efforts are devoted in this phase including (semi-)analytic Newtonian \\cite{LaiRS93,TanigN00} and post-Newtonian \\cite{LombaRS97,TanigS97,ShibaT97,Tanig99} approaches and numerical Newtonian \\cite{TanigGB01,TanigG02,HachiE84a,HachiE84b,UryuE98}, post-Newtonian \\cite{Shiba96}, and general relativistic \\cite{BonazGM99a,GourGTMB01,BaumgCSST97,MarroMW98,MarroMW99,UryuE00,UryuSE00,UsuiUE99,UsuiE02} ones. Among these studies, we concentrate on numerical computations in the general relativistic framework. Kochaneck \\cite{Kocha92} and Bildsten and Cutler \\cite{BildsC92} have shown that the gravitational-radiation driven evolution is too rapid for the viscous forces to synchronize the spin of each neutron star with the orbit as they do for ordinary stellar binaries. Rather, the viscosity is negligible and the fluid velocity circulation (with respect to some inertial frame) is conserved in these systems. Provided that the initial spins are not in the millisecond regime, this means that close binary configurations are well approximated by zero vorticity (i.e. {\\em irrotational}) states. Formulations for irrotational binaries in general relativity have been developed by several authors \\cite{BonazGM97,Asada98,Shiba98,Teuko98} (see Appendix~A of Paper~I for a comparison between them). In addition to the quasiequilibrium and irrotational assumptions, we use the approximation of a conformally flat spatial 3-metric, introduced by Isenberg \\cite{IsenbN80} and Wilson and Mathews \\cite{WilsoM89} (hereafter IWM approximation; see Sec. IV.C of \\cite{FriedUS02} and Sec. III.A of Paper~I for a discussion).  Until now, several groups have produced quasiequilibrium configurations of binary neutron stars, as listed above. Among them, there are results for synchronized \\cite{BaumgCSST97,MarroMW98} and irrotational \\cite{BonazGM99a,GourGTMB01,MarroMW99,UryuE00,UryuSE00} rotation states in the general relativistic framework, within the IWM approximation, and those for synchronized one within some axisymmetric approximation \\cite{UsuiUE99,UsuiE02}. However they all deal with {\\em identical} star binaries and calculations for {\\em different} mass binary system in general relativity have not been performed yet (in Newtonian theory they have been performed in the synchronized case \\cite{HachiE84b,TanigG02}, and recently in the irrotational one \\cite{TanigG02}).  We will present here the first numerical results of relativistic binary systems composed of different mass stars. In addition, we give results for systems of identical stars which extend those already presented in Paper~I. As a first step in the study of different mass systems, we restrict our computations to a polytropic equation of state, with the adiabatic index $\\gamma=2$.  The plan of the article is as follows. A brief overview of our method is given in Sec.~\\ref{s:method}. In Sec.~\\ref{s:tests}, new tests of the numerical code are shown, which were not given in Papers~I and II. We present the numerical results in Sec.~\\ref{s:results}, and discuss them in Sec.~\\ref{s:discussion}. Finally Sec.~\\ref{s:summary} is devoted to the summary.  Throughout this article, we adopt geometrical units: $G=c=1$, where $G$ and $c$ denote the gravitational constant and speed of light, respectively. ",
        "conclusions": "\\label{s:discussion}  In this section, we discuss about the behaviors of the axial ratios, the relative change in central energy density, and the turning point of the synchronized sequences.  In Figs. \\ref{fig:axial_ratio_eq} and \\ref{fig:axial_ratio_df}, the axial ratios are shown as a function of the cusp indicator $\\chi$. In Fig. \\ref{fig:axial_ratio_eq}, we show the axial ratios $a_2/a_1$ and $a_3/a_1$ for identical mass binaries with the compactness $M/R=0.12$ vs. 0.12 and 0.16 vs. 0.16. One can see from this figure that each axial ratio has almost the same track regardless of the difference of the compactness, in particular for synchronized binary systems. Moreover, such a coincidence of the sequence appears in the cases of different mass binary systems (see Fig. \\ref{fig:axial_ratio_df}). In those cases, the sequences almost coincide with each other regardless of the compactness of the companion star. Note that the sequences of the axial ratios for more massive stars end before $\\chi=0$, while those for less massive stars should reach the point $\\chi=0$. However we cannot treat a cuspy figure with our numerical method, because we adapt the innermost numerical domain to the surface of the star and the domain boundaries are assumed to be smooth (see discussion in Sec.~\\ref{s:results}). Therefore we had to stop the sequences at around $\\chi=0.4$ for the less massive star.  \\begin{figure}[htb] \\vspace{0.5cm} \\begin{center} \\includegraphics[width=8cm]{fig13.eps} \\end{center} \\caption[]{\\label{fig:dec_vs_chi} Relative change in central energy density as a function of the cusp indicator $\\chi$. Left (resp. right) panel is for synchronized (resp. irrotational) binaries. Solid lines with open circle, square, diamond, and triangle denote the cases of the identical mass binaries with the compactness $M/R=0.12$ vs. 0.12, 0.14 vs. 0.14, 0.16 vs. 0.16, and 0.18 vs. 0.18, respectively. Thin and thick dotted lines are for the compactness $M/R=0.12$ vs. 0.14 of the less massive and more massive stars, respectively. Thin and thick dashed lines are for the compactness $M/R=0.14$ vs. 0.16 of the less massive and more massive stars, respectively. Thin and thick long-dashed lines are for the compactness $M/R=0.16$ vs. 0.18 of the less massive and more massive stars, respectively. } \\end{figure}%  \\begin{figure}[htb] \\vspace{0.5cm} \\begin{center} \\includegraphics[width=8cm]{fig14.eps} \\end{center} \\caption[]{\\label{fig:turning} Cusp indicator $\\chi$ at the turning point of the ADM mass along a sequence of synchronized identical mass binaries, as a function of the compactness $M/R$. } \\end{figure}%  The relative change in central energy density is presented as a function of $\\chi$ in Fig.~\\ref{fig:dec_vs_chi}. It is clearly seen that the sequences for the synchronized case with the same compactness coincide with each other, while those with the different compactness split into curves with respect to each compactness. For example, the relative change for the identical binary system with the compactness $M/R=0.14$ coincides with that for the more massive star of the different mass binary with $M/R=0.12$ vs. 0.14 and that for the less massive star with $M/R=0.14$ vs. 0.16. However, the sequences for stars with $M/R=0.12$, 0.14, and 0.16 are different from each other.  Finally, let us comment on the behavior of the turning point of the ADM mass (binding energy) along a sequence. We never found any turning points for irrotational binary systems with the polytropic index $\\gamma=2$. On the other hand, it appeared clearly for the synchronized case. In Fig.~\\ref{fig:turning}, the cusp indicator $\\chi$ at the turning point of synchronized identical mass binary systems is plotted as a function of the compactness $M/R$. Since the accuracy of the determination of the ADM mass is order of $10^{-5}$ in our calculations, the relative errors on the orbital angular velocity and the quantity $\\chi$ at the turning point become a few percents. Even when one is taking this fact into account, it is possible to conclude that the line in Fig.~\\ref{fig:turning} increases proportionally to the compactness of the stars. This behavior agrees with the value of $\\chi_{\\rm turning}$ that we have found in Newtonian calculations ($M/R=0$), namely $\\chi_{\\rm turning}=0.2862$ (Table~I of Paper~II).     In the present paper, we have computed constant baryon number sequences of binary neutron stars in quasiequilibrium, in both cases of synchronized and irrotational motion and in both cases of identical and different mass systems. We have performed a general relativistic treatment, within the IWM approximation (conformally flat spatial metric). We have used a polytropic equation of state with an adiabatic index $\\gamma=2$.  The summary of our results is that (1) Among the quasiequilibrium sequences we have calculated in the present article, only the synchronized identical star binaries terminate by the contact between two stars, while all the other types of sequences (synchronized different star binaries, irrotational identical star and different star binaries) end at the mass shedding points. Note here that we cannot conclude about the end points of {\\em non-quasiequilibrium} sequences (i.e. dynamical sequences), such that binary systems without viscosity but with intrinsic non-aligned spins, those with slight viscosity and deviating from both the irrotational state and the synchronized one, those with infalling radial velocity, and so on. (2) For identical mass binary systems there is no turning point of the ADM mass in the irrotational case, while there is clearly one before the end point of the sequence in the synchronized case. (3) It is more difficult to see the turning points of the ADM mass for different mass binary systems than for identical ones. (4) The deformation of the star is determined by the orbital separation and the mass ratio and is not affected much by its compactness. (5) The decrease of the central energy density depends on the compactness of the star and not on that of its companion."
    },
    "0207/astro-ph0207016_arXiv.txt": {
        "abstract": "{The interaction of the wind from a pulsar (or more generally from a star) with the ambient medium gives rise to the formation of a bow-shock nebula. We present a model of adiabatic bow-shock nebulae, including the presence of a neutral component in the interstellar medium. As demonstrated in a previous paper (Bucciantini \\& Bandiera \\cite{bucciantini01}, hereafter Paper I), the velocity and the luminosity of the pulsar, as well as the characteristics of the interstellar medium (ISM), play an essential role in determining the properties of bow shock nebulae. In particular, the Balmer emission is strongly affected by the rates of hydrogen charge-exchange and collisional excitation. So far, only one-component models have been proposed, treating the neutrals as a small perturbation. The distribution of neutral hydrogen in the nebula, derived by our model, is then used to determine how the \\halpha \\ luminosity varies along the bow shock. We find that the luminosity trend is different for charge-exchange and collisional excitation, and that processes like diffusion or ionization may reduce the emission in the head of the nebula. ",
        "introduction": "Pulsars are known to be sources of ultrarelativistic magnetized winds (see e.g. Michel \\& Li \\cite{michel99}). If during the supernova explosion the newly formed neutron star acquires a high velocity, it may escape from the supernova remnant before its pulsar activity has ceased. In such a case the pulsar wind interacts with the interstellar medium (thereafter ISM), possibly forming a steady flow, with a bow shock in the direction of the stellar motion. So far only four such nebulae are known: PSR 1957+20 (Kulkarni \\& Hester \\cite{kulkarni88}), PSR 2224+65 (Cordes et al. \\cite{cordes93}, Chatterjee \\& Cordes \\cite{chatterjee02}), PSR J0437+4715 (Bell \\cite{bell95}), and PSR 0740-28 (Jones et al. \\cite{jones01}). Probably even the nebula associated with the neutron star RX J185635-3754 (van Kerkwijk \\& Kulkarni \\cite{vankerkwijk01}) may be included in the list, even if there is no agreement whether the nebula is a bow shock or the result of photoionization.  All the nebulae discovered have been detected in Balmer lines (mostly \\halpha ). A pure Balmer spectrum is the signature of a non-radiative shock moving through a partially neutral medium (Chevalier \\& Raymond \\cite{chevalier80}, Chatterjee \\& Cordes \\cite{chatterjee02}). When recombination times are long the relevant processes for optical emission are exciting charge-exchange and collisional excitation of neutral H atoms.  As mentioned in a previous paper of this series (Bucciantini \\cite{bucciantini02}, hereafter Paper II), the ISM is seen by the pulsar as a plane parallel flow with a constant density (at least on the length scale of the bow shock). If the ISM is partially neutral, and if the neutral atoms have mean free paths comparable with, or longer than the size of the bow shock a pure fluid treatment is invalid. The interaction of the relativistic magnetized pulsar wind with the ionized component of the ISM is mediated by the magnetic field advected by the pulsar wind and compressed in the head of the nebula. The neutral atoms interact with the ions (mainly protons) thermalized in the bow shock. The penetration thickness for the H atoms changes from pulsar to pulsar depending on their velocity, their luminosity as well as the density and the neutral fraction of the ISM; and the same holds for the diffusion and ionization (Paper I). Thus to have a more realistic picture of bow-shock nebulae in this regime, it is necessary to include the effect of mass loading from a neutral component as well as the diffusion and ionization of dragged atoms.  As a starting point we have used a two thin layers  bow shock model, for which an analytic solution may be derived (Comeron \\& Kaper \\cite{comeron98}). In Paper II we have verified, using a numerical hydrodynamic code, that this analytic model gives reasonable estimates for the values of surface density and tangential velocity in the external layer. These estimates describe the fluid structure in the external layer and have been used to model the interaction and the evolution of the neutral atoms.  Once the density of the neutral H atoms in the external layer has been derived, it is used to evaluate the \\halpha \\ luminosity of the nebula, and its variation along the bow-shock. The comparison of the observed luminosity with that obtained from a thin layer model may be used to determine the properties of the local ISM, as well as those of the pulsar wind. In Section 2 we present the equations we have used to derive the distribution of neutral atoms, and the solutions derived integrating these equations; in Section 3 we derive the \\halpha \\ luminosity of the nebula and discuss what informations may be obtained from observations. ",
        "conclusions": "We have modified the two thin layer solution (Comeron \\& Kaper \\cite{comeron98}) for bow shock produced by runaway stars (with the aim of deriving a model in the case of pulsar wind bow shock nebulae) with the inclusion of a neutral component. We find that the ISM neutrals lead only to minor effects to the shape of the nebulae (see Fig.~2), but plays an essential role in their luminosity. We have also introduced in our model the possibility that the H atoms, once stopped by the thermalized protons, may diffuse out of the layer or be ionized, as a  consequence of charge-exchange and collisional ionization. Both processes act depleting the neutral surface density in the external layer, but their efficiency is higher in the head of the nebula both as a consequence of the decrease of the rate of charge-exchange and ionization in the tail, and the increase of the geometrical thickness of the layer.  The distributions obtained have been used to evaluate the Balmer emission along the nebula. The study of such emission can be a good tool to understand the properties of local ISM as well as those of the pulsar.  Even if this is not a simple task, and requires very good measurements, which only for two objects are actually available, an accurate study may show which hypothesis can be applied (like, for example, the absence of preionization both from the pulsar surface radiation, or from the emission of shocked ISM). In Paper I, we stressed the point that the total \\halpha\\ flux depends on the velocity of the pulsar.  Here we evaluate its variations along the nebula.  The behaviour is different if the dominant process is exciting charge-exchange or collisional excitation, so not only the total flux but also its variation moving away from the stagnation point can be used. Concerning the emission we find that: 1- charge-excange and collisional excitation have different trends; 2- processes like diffusion or ionization may play an essential role changing the emission in the head of the nebula.  This not only allows one to put stronger constraints on the regime of such nebulae, but it may be used to verify the validity of the model, by comparison with the observed cases."
    },
    "0207/astro-ph0207220_arXiv.txt": {
        "abstract": "The case for small neutrino mass differences from atmospheric and solar neutrino oscillation experiments has become compelling, but leaves the overall neutrino mass scale $m_\\nu$ undetermined.  The most restrictive limit of $m_\\nu<0.8~{\\rm eV}$ arises from the 2dF galaxy redshift survey in conjunction with the standard theory of cosmological structure formation. A relation between the hot dark matter fraction and $m_\\nu$ depends on the cosmic number density $n_\\nu$ of neutrinos.  If solar neutrino oscillations indeed correspond to the favored large mixing angle MSW solution, then big-bang nucleosynthesis gives us a restrictive limit on all neutrino chemical potentials, removing the previous uncertainty of $n_\\nu$. Therefore, a possible future measurement of $m_\\nu$ will directly establish the cosmic neutrino mass fraction $\\Omega_\\nu$. Cosmological neutrinos with sub-eV masses can play an interesting role for producing the highest-energy cosmic rays (\\hbox{$Z$-burst} scenario).  Sub-eV masses also relate naturally to leptogenesis scenarios of the cosmic baryon asymmetry. Unfortunately, the time-of-flight dispersion of a galactic or local-group supernova neutrino burst is not sensitive in the sub-eV range. ",
        "introduction": "Atmospheric and solar neutrino experiments provide rather compelling evidence for the phenomenon of flavor oscillations.  The celebrated up-down-asymmetry of the atmospheric $\\nu_\\mu$ flux measured by Super-Kamiokande is consistently explained by $\\nu_\\mu\\to\\nu_\\tau$ oscillations \\citep{Fukuda2000} with the mixing parameters that are summarized in Table~\\ref{tab:osci}. The K2K long-baseline experiment provides a first laboratory confirmation, albeit in a pure disappearance mode \\citep{Nishikawa2002}. The recent results from SNO have largely established active-active flavor oscillations as a solution of the solar neutrino problem \\citep{Ahmad2002a,Ahmad2002b}. The LMA parameters are strongly favored, but the LOW case may still be viable (Table~\\ref{tab:osci}). Neutrino mass differences that are small compared to the eV scale seem to be established.  \\begin{table} \\caption{Experimental evidence for neutrino flavor oscillations.\\label{tab:osci}} \\smallskip \\begin{tabular}[4]{llll} \\hline\\noalign{\\vskip2pt}\\hline\\noalign{\\vskip2pt} Evidence&Channel&$\\Delta m^2$ [$\\,\\rm eV^2$] &$\\sin^22\\Theta$\\\\ \\noalign{\\vskip2pt}\\hline\\noalign{\\vskip2pt} Atmospheric&$\\nu_\\mu\\to\\nu_\\tau$&(1.6--3.9)${}\\times10^{-3}$&0.92--1\\\\ Solar\\\\ \\quad LMA &$\\nu_e\\to\\nu_{\\mu\\tau}$& (0.2--2)${}\\times10^{-4}$&0.2--0.6\\\\ \\quad LOW&$\\nu_e\\to\\nu_{\\mu\\tau}$ &$1.3\\times10^{-7}$ &0.92\\\\ LSND&$\\bar\\nu_\\mu\\to\\bar\\nu_e$&0.2--10&(0.2--$3)\\times10^{-2}$\\\\ \\noalign{\\vskip2pt} \\hline \\noalign{\\vskip2pt} \\hline \\end{tabular} \\bigskip \\bigskip \\end{table}  The only spanner in the works of this beautiful interpretation is the persistence of the unconfirmed evidence for flavor transformations from the LSND experiment~\\citep{Eitel2000}. If interpreted in terms of neutrino oscillations, the mixing parameters from all three sources of evidence are mutually inconsistent. Even including a putative sterile neutrino no longer provides a good global fit because all three sources of evidence prefer active-active over active-sterile oscillations \\citep{Strumia2002}. It is expected that MiniBooNE at Fermilab will confirm or refute LSND within the upcoming two years \\citep{Tayloe2002}.  As there is no straightforward global interpretation of all indications for neutrino oscillations I will follow the widespread assumption that something is wrong with the LSND signature.  If it is due to neutrino conversions after all, something fundamentally new is going on in the neutrino sector.  In that case much of the current thinking in this field will have to be revised.  This is certainly the attitude that Dennis Sciama would have taken. He eloquently advocated a cosmological scenario of radiatively decaying dark-matter neutrinos which solves several astrophysical problems, but requires several flavors of light sterile neutrinos and anomalously large electromagnetic transition moments \\citep{Adams1998,Sciama1998}.  No doubt he would have challenged the more conservative approach taken here. The final verdict on neutrino masses and mixings is not yet in, let alone on the more exotic possibilities.  The phenomenology of the neutrino sector may turn out to be far richer than envisaged in my presentation of a rather minimal scenario.  In what follows I will always assume that there are three neutrino mass eigenstates separated by the atmospheric and solar mass differences. In this scenario, a number of obvious questions remain open. The 12 and 23 mixing angles are large, the 13 mixing angle is small, but how small? Are there CP-violating phases in the mixing matrix? Are the neutrino masses of Dirac or Majorana nature?  Is the ordering of the masses ``normal'' with $m_2^2-m_1^2$ corresponding to the solar and $m_3^2-m_2^2$ to the atmospheric splitting, or is it inverted?  And finally, what is the overall neutrino mass scale? Are the masses hierarchical with $m_1\\ll m_2\\ll m_3\\approx 50~{\\rm meV}$ or degenerate with $m_1\\approx m_2\\approx m_3\\gg 50~{\\rm meV}$?  I will focus on these last questions and review the implications of neutrino masses in astrophysics and cosmology. Traditionally, cosmology has provided the most restrictive limits on neutrino masses, and this is again the case using large-scale galaxy redshift surveys in conjunction with the standard theory of structure formation. Conversely, if the solar LMA solution is indeed correct, the cosmic neutrino number density is well constrained by big-bang nucleosynthesis so that a laboratory measurement of the absolute neutrino masses, for example in the KATRIN tritium experiment \\citep{Osipowicz2001}, would directly establish the cosmic neutrino mass fraction. Moreover, neutrino masses can have a number of other interesting implications in astroparticle physics in the context of cosmic-ray physics, cosmological baryogenesis, and SN physics.  In Sec.~2 this review begins with neutrino dark matter and the latest $m_\\nu$ limits from large-scale redshift surveys.  Sec.~3 turns to the related question of how many neutrinos there are in the universe and how this issue connects with the solar neutrino problem.  Sec.~4 deals with $Z$-burst scenarios for producing the highest-energy cosmic rays, Sec.~5 with leptogenesis scenarios for producing the baryon asymmetry of the universe. Sec.~6 is devoted to the time-of-flight dispersion of supernova neutrinos caused by a non-vanishing $m_\\nu$.  Finally, Sec.~7 summarizes the status of neutrino masses in astroparticle physics. ",
        "conclusions": "The compelling detection of flavor oscillations in the solar and atmospheric neutrino data have triggered a new era in neutrino physics. In the laboratory one will proceed with precision experiments aimed at measuring the details of the mixing matrix. Future tritium decay experiments may well be able to probe the overall neutrino mass scale down to the 0.3~eV range, but if the absolute masses are smaller, it will be very difficult to measure them, and the overall mass scale may remain the most important unknown quantity in neutrino physics for a long time to come.  Unfortunately, it is unlikely that astrophysical time-of-flight methods will help much. Foreseeable SN neutrino detectors are sensitive to eV masses, but not the sub-eV range. On the bright side this means that the measured neutrino light-curve of a future galactic SN will faithfully represent the source without much modifications by neutrino dispersion.  Cosmological large-scale structure data at present provide the most restrictive limit on neutrino masses of $\\sum m_\\nu<2.5~{\\rm eV}$, corresponding to $m_\\nu<0.8~{\\rm eV}$ in a degenerate mass scenario. A rigorous relationship between the cosmic hot dark matter fraction $\\Omega_\\nu$ and $m_\\nu$ depends on the cosmic neutrino density~$n_\\nu$. If the solar LMA solution is correct, big-bang nucleosynthesis constrains $n_\\nu$ without further assumptions about the neutrino chemical potentials. In the LMA case neutrinos reach de-facto flavor equilibrium before the epoch of weak-interaction freeze out.  While neutrinos do not play a dominant role for dark matter or structure formation, the mass and mixing schemes suggested by the oscillation experiments are nicely consistent with leptogenesis scenarios for creating the cosmic baryon asymmetry. Therefore, massive neutrinos may be closely related to the baryons in the universe, not the dark matter.  If extremely-high energy neutrinos will be observed in future, the $Z$-burst scenario provides a handle on the cosmic background neutrinos and their mass through the observed cosmic rays near the GZK cutoff.  The great advance in our knowledge of neutrino properties together with the cosmological precision information has rendered the question of neutrino masses in astrophysics and cosmology far more subtle than it looked only a few years ago.  We will not know if the accumulation of new results would have persuaded Dennis Sciama that nature has used massive neutrinos perhaps in different ways for cosmology than he himself had imagined for so long.  Either way, the connection between neutrino properties and astroparticle physics remains of fundamental interest.  \\newpage"
    },
    "0207/nucl-th0207008_arXiv.txt": {
        "abstract": "The proton-proton ($pp$) fusion cross-section found at the heart of solar models is unconstrained experimentally and relies solely on theoretical calculations. Effective field theory provides an opportunity to constrain the $pp$ cross-section experimentally, however, this method is complicated by the appearance of two-nucleon effects in the form of an unknown parameter $L_{1,A}$.  We present a method to constrain $L_{1,A}$ using the Standard Solar Model and helioseismology.  Using this method, we determine a value of $L_{1,A}$ = 7.0 fm$^{3}$ with a range of 1.6 to 12.4 fm$^{3}$.  These results are consistent with theoretical estimates of $L_{1,A} \\approx$ 6 fm$^{3}$. ",
        "introduction": "The Standard Solar Model (SSM) is a model of the sun constructed by numerically solving the equations of stellar structure and evolution constrained by the observed composition, luminosity, and radius of the sun at the current age. This model yields acoustic oscillation frequencies ($p$-modes) that match the observed $p$-mode oscillation spectrum of the sun to better than 0.3\\%~\\cite{Gue97}. The uncertainties associated with the observed oscillation frequencies are an order of magnitude smaller.  The excellent agreement between the model and observed frequency spectra implies that the run of sound speed predicted by the model is close to that of the sun which, in turn, suggests that the physics used to construct the model of the sun is accurate.  The individual $p$-mode frequencies extend to distinct depths in the sun, and hence can be used to study the physics of specific regions.  Of interest to the work described here are the low-$l$ (where $l$ corresponds, in spherical harmonic nomenclature, to the number of nodes in azimuth) $p$-mode frequencies that penetrate deep into the core of the sun.  These modes, grouped in combinations that cancel out surface dependencies, can be used to test the important physics of the core.  One of the greatest sources of uncertainty in the physics of the solar model core is the $pp$-fusion cross-section.  The reaction is unconstrained by laboratory-based experiments, relying completely on precise calculations from theoretical nuclear physics.  Owing to the precise agreement between the observed and predicted oscillation frequencies that already exists, it is possible to use the oscillation frequencies to constrain the $pp$ cross-section, that is, to determine the range of cross-sections that yield oscillation frequencies within the known uncertainties of the solar model physics and constraints.  To date, this sort of constraint has been imposed only through simplifications of the standard solar model~\\cite{antia}, and here we use a fully-developed stellar evolution code, to be discussed later.  The cross-section for $pp\\rightarrow de^+\\nu_e$ has been studied extensively \\cite{Sal52,Bli65,Bah69,Gar72,Dau76,Bar79,Gou90,Car91,Kam94,Iva97,Sch98,Par98} and contains two components.  First, an `Impulse Approximation' (IA) contribution where the weak interaction takes place on a single proton, as shown in Fig.~\\ref{fig1}a).  This comprises more than 95\\% of the cross-section at energies of solar interest and is very well understood because it concerns only the weak properties of a proton and well-known proton-proton scattering physics.  The second component, and the remainder of the cross-section, is unconstrained by experiment and has led theorists to try to improve our understanding of this process. Conventional potential model calculations will introduce so-called meson-exchange currents; essentially two-nucleon effects where the weak interaction occurs while the two nucleons are interacting with each other.   An example of such a process is shown in Fig.~\\ref{fig1}b).  These two-nucleon effects introduce physics that cannot be constrained directly by $pp$ elastic scattering experiments, and the parameters of this extended physics must be constrained using other methods. Few other methods exist. \\begin{figure}[t] \\centerline{{\\epsfxsize=5.0in \\epsfbox{Figure1.eps}}} \\noindent \\caption{{\\it The process $pp\\rightarrow d e^+\\nu_e$. a) The Impulse Approximation (IA), where the weak interaction takes place on one proton while the other proton is just a spectator.  The final state, represented by the grey bubble, is the resulting neutron and proton being bound as a deuteron. b) A two-nucleon process where the protons are undergoing a strong interaction at the same time as the weak interaction (represented by the black circle).}} \\label{fig1} \\end{figure} Neutrino-deuteron scattering could, in principle, constrain the two-nucleon effects. Experiments involving neutrino-deuteron scattering are, however, quite difficult. Reactor antineutrino-deuteron experiments can be performed, where the experimental uncertainties are of order 10-20\\%. Nonetheless, they can be used to infer a constraint on the two-nucleon matrix elements~\\cite{bcv}. One interesting method is to use tritium $\\beta$-decay to constrain the two-nucleon matrix elements needed in $pp\\rightarrow de^+\\nu_e$~\\cite{Sch98}. The tritium half-life is well known, and the same two-nucleon effects occur here as in $pp$ fusion. However, there are three nucleons in tritium, hence we cannot be certain that this more complicated structure does not, in turn, complicate the weak interaction physics further.  If we were to rely purely on theory to predict the two-nucleon component of the $pp$ fusion cross-section, we still have one more problem.  How do we define the uncertainty in a theoretical calculation?  In quantum field theories, such as quantum electrodynamics, there is a perturbative expansion parameter that dictates the size of contributions at each order in perturbation theory and thus the size of any neglected remainder. The remainder represents the uncertainty in the theory.  However, conventional nuclear physics calculations use   Schr\\\"odinger's equation and nucleon-nucleon potentials to calculate initial and final-state wavefunctions and, in turn, calculate the matrix elements of relevant transition operators.  In this approach, no scheme exists for estimating the size of contributions from unknown or neglected physics. An alternative to using wavefunctions to study our problem is to use an effective field theory (EFT).  An EFT has many benefits, but there are drawbacks. If the EFT is perturbative (i.e., there is a small expansion parameter), then one can compute to a specific order in the parameter and then make direct statements of the scale of theoretical uncertainty of the problem as determined by the relative size of the next term in the expansion - just as one would do in quantum electrodynamics. This is particularly important given that EFTs typically have an infinite number of terms, and the expansion parameter is required to ensure that only a finite number of terms are needed for any given calculation.  In the particular case of nuclear EFTs, specifically ones where all interactions are of zero-range (i.e., no pions), the expansion parameter is denoted by $\\epsilon=Q/M_\\pi$, where $Q$ represents a generic momentum in the problem and $M_\\pi$ is the pion mass. An added benefit of this approach is that field theories pose no problems for conservation laws such as electromagnetic gauge invariance, which are often difficult to preserve when wavefunctions are used. However, EFTs cannot escape the problem that all parameters in an EFT must be constrained by experiment, including the various weak interaction parameters. ",
        "conclusions": "Using the relationship between effective field theory, nuclear cross-sections and the accuracy of helioseismology, the unknown counterterm, $L_{1,A}$, is determined to have a value of 7.0 fm$^3$ with a range of 1.6 to 12.4 fm$^3$.  This result for $L_{1,A}$ is consistent with the theoretical value determined by Butler $\\&$ Chen \\cite{BCK} using dimensional analysis of $L_{1,A}$ $\\approx$ 6 fm$^3$. The result and range is also consistent with other evaluations of $L_{1,A}$ which include 5.6$\\pm$2.0 fm$^3$ \\cite{NSGK} and 6.5$\\pm$2.4 fm$^3$ \\cite{Sch98}.  The standard solar model determined for this work shows remarkably excellent agreement with the latest BiSON oscillation frequency observations.  Theoretical total fluxes for $^8B$ neutrinos from the standard solar model result also show excellent agreement with total observed $^8B$ fluxes reported from the Sudbury Neutrino Observatory."
    },
    "0207/astro-ph0207121.txt": {
        "abstract": "{ This paper is a further step in the investigation of the morphology of the color-magnitude diagram of Galactic globular clusters, and the fine-tuning of theoretical models, made possible by the recent observational efforts to build homogeneous photometric databases.  In particular, we examine here the calibration of the morphological parameter $W_{\\rm HB}$ vs. metallicity, originally proposed by Brocato et al. (\\cite{brocatoEtal98}; B98), which essentially measures the color position of the red-giant branch. We show that the parameter can be used to have a first-order estimate of the cluster metallicity, since the dispersion around the mean trend with [Fe/H] is compatible with the measurement errors. The tight $W_{\\rm HB}$-[Fe/H] relation is then used to show that variations in helium content or age do not affect the parameter, whereas it is strongly influenced by the mixing-length parameter $\\alpha$ (as expected). This fact allows us, for the first time, to state that there is no trend of $\\alpha$ with the metal content of a cluster. A thorough examination of the interrelated questions of the $\\alpha$-elements enhancement and the color-$T_{\\rm eff}$ transformations, highlights that there is an urgent need for an independent assessment of which of the two presently accepted metallicity scales is the true indicator of a cluster's iron content. Whatever scenario is adopted, it also appears that a deep revision of the $V-I$-temperature relations is needed. ",
        "introduction": "\\begin{figure} {\\par\\centering \\resizebox*{0.9\\columnwidth}{!}{\\includegraphics{h3600f1.ps}} \\par} \\caption{M80 (filled squares) vs. M12 (open squares). An offset $\\delta (V-I)=-0.03$ and $\\delta V=+1.50$ was applied to the M12 diagram before overlapping it to that of M80 \\label{fig:m80vsm12}} \\end{figure}  The comparison of the observed color-magnitude diagrams (CMD) of a star cluster with the theoretical isochrones is the best tool we have to tune some fundamental parameters of the stellar evolutionary models. Only when we are sure that the input parameters and the input physics are correct, we can use the model to infer some properties of the stellar population (like age, helium content, etc.), not otherwise empirically measurable. Clearly, this is a complex job: on one side we want to use the cluster stars to tune the models, and on the other we need to use the models to infer properties of the cluster stellar population. Not surprisingly, any stellar models must adopt a number of assumptions, often not directly supported by observational evidence. One of the most uncertain parameters, strongly affecting the theoretical location of some branches of the CMD, like the red giant branch, is the mixing length parameter $\\alpha$.  This parameter, in the framework of the mixing-length theory (MLT; B\\\"ohm-Vitense \\cite{bohm-vitense58}), determines the efficiency of energy transport by convection in the outermost layers of a star. For a given stellar luminosity, it also determines the exact radius of the star, and hence its effective temperature and colors.  It is well known that, in order to reproduce the typical temperatures and colors of red giant stars, $\\alpha$ is required to have a value between 1.5 and 2.0. Also, a value close to 1.7 is required for reproducing the solar radius in non-diffusive solar models, whereas about 1.9 is favored when diffusion is taken into account. The fact that $\\alpha$ is similar for red giants and the Sun {\\changed (VandenBerg et al. \\cite{vandenberg00}; Alonso et al. \\cite{alonsoEtal00})} leads to the usual approach of calibrating $\\alpha$ by means of a solar model -- i.e. a model with 1~\\Msun\\ and solar composition, that is required to have the solar luminosity and radius at an age of $\\sim 4.5$~Gyr.  The same $\\alpha$ value is then used to model all stars, including red giants.  However, there is no theoretical justification that the same value of $\\alpha$ should apply for any star.    In this paper, we have attempted to investigate which is the best value of $\\alpha$ for low mass (globular cluster) stars and, overall, whether there is any dependence of $\\alpha$ with the cluster metallicity. Our results cannot be considered definitive; however this paper shows a possible approach to the problem, and enlightens all the uncertainties associated to the calibration of the mixing length parameter. Also, this work is limited to the context of the MLT, which is still a fundamental approach adopted in most computations of stellar models. However, one should keep in mind that the MLT is admittedly a very approximate theory, and that alternative approaches (e.g.\\ Canuto \\& Mazzitelli \\cite{canutomazzitelli91}; Spruit \\cite{spruit97}; Ludwig et al. \\cite{ludwigEtal99}) have been suggested.  This project has been stimulated by an investigation carried out by some of us a few years ago.  In the course of a photometric study of the Galactic globular cluster (GGC) M80 (Brocato et al. \\cite{brocatoEtal98}, hereafter B98), we compared the morphological characteristics of its color-magnitude diagram (CMD) with those of an ensemble of other globulars for which a CMD was available in the literature. That investigation left an open question concerning the relative position of the RGB with respect to the HB as a function of the metallicity.  More specifically, when the CMD of M80 was compared with that of GGCs with similar metallicity, we found a group of clusters (M3, M13, NGC~7492) whose members had CMDs overlapping that of M80, and another (M12, NGC~1904, NGC~5897) for which significant discrepancies were seen. The discrepancies were in the sense that while the HBs could be overlapped in a satisfactory way, the RGB fiducial lines showed a dispersion in color, M12 having the reddest branch (cf Fig. 6 and 7 in B98), as expected if M12 would be more metal rich than indicated in the literature. We also excluded that the discrepancy could be due to an age difference.  In order to have a more quantitative comparison among a larger sample of clusters, with different metal content, we devised a new morphological parameter with the aim of quantitatively measuring the distance in color of the RGB with respect to some fixed point on the HB.  We selected as a reference point the so called {}``HB turn-down{}''(HB-TD), since it has been demonstrated that the location of this point in the CMD is largely unaffected by the cluster age, metallicity, and primordial He content (cf. the detailed discussion in B98).  We fixed the HB-TD at $(B-V)_{0}=0$ and measured the RGB color at $0.5$ magnitudes brighter than the HB-TD level.  As the position of the HB- TD is fixed, $W_{\\rm HB}$ just shows the displacement of the RGB as a function of the metal content (cf. Fig.~10 of B98). We found that the trend with the metallicity is in the direction expected from the theoretical models, though with a large dispersion, particularly evident at intermediate metallicities ($-1.85<${[}Fe/H{]}$<-1.50$), larger than expected on the basis of the measurement errors.  We were not able to decide whether such an evidence might simply be the consequence of the errors in the photometric calibration, or a peculiar distribution in the global metallicity of the clusters, and deferred further discussion until direct measurements of $\\alpha$ elements and/or a database of CMDs with homogeneous calibration would have been available.  Such database is now available, thanks to the efforts of a number of the original investigators of B98, and we are now able to re-examine the question, taking also advantage of the new theoretical calculations that we have specifically performed for this project.   The paper is organized as follows. The datasets are presented in Sect.~\\ref{sec:datasets}. The measurement procedures and the corrections for differential reddening are described in Sect.~\\ref{sec:measurements} and \\ref{sec:diffredd-correction}, respectively. Sect.~\\ref{sec:thewhb-parameter} re-examines the original question of the trend and dispersion of $W_{\\rm HB}$ as a function of metallicity.  Sect.~\\ref{sec:compare-models} deals with the comparison of the VandenBerg et al. (\\cite{vandenberg00}; V00) and the Girardi et al. (\\cite{girardiEtal00}; G00) theoretical isochrones to the data. Due to the lower degree of sampling of the G00 isochrones, a more detailed description of the turn-down identification and its measurement is offered in Sect.~\\ref{sec:isoc-g00}.  The keys to the interpretation of the observed vs. computed trend of $W_{\\rm HB}$ are offered in Sect.~\\ref{sec:what-det-whb}. In particular, the dependence on metallicity, mixing length, age, helium content, $\\alpha$-enhancement, and the \\Teff-color transformations are examined in Sects.~\\ref{sec:what-met}, \\ref{sec:what-mixlen}, \\ref{sec:what-age}, \\ref{sec:what-helium}, \\ref{sec:what-enha}, and \\ref{sec:what-transf}.  Whether our data can be used  to constrain the mixing length parameter is examined in Sect.~\\ref{alpha-only}. Furthermore, we suggest in Sect.~\\ref{sec:metindex} that the parameter can be employed as a metallicity indicator, with accuracies comparable to other more widely used photometric indices. Our summary and conclusions are given in Sect.~\\ref{sec:sum-discussion}.  The nomenclature deserves a final note. We will be using several symbols for our parameter throughout this paper. When talking about the parameter in a general way, we will use the old symbol $W_{\\rm HB}$, but when referring to measurements specifically made for the $(B-V)$ or $(V-I)$ colors, we will use the symbols $W^{B-V}_{\\rm HB}$ and $W^{V-I}_{\\rm HB}$, respectively. ",
        "conclusions": "}  In this paper, we took advantage of the homogeneous photometric databases of Galactic globular clusters, presented in Rosenberg et al (\\cite{dutch00}, \\cite{jkt00}) and Piotto et al. (\\cite{piottoEtal02}), to complete another step in the characterization of the morphological properties of the CMD of GGCs, and in the fine-tuning of the theoretical models. To this aim, the $W_{\\rm HB}$ parameter, originally defined in B98, has been measured for all suitable clusters, and appropriate corrections to account for differential reddening effects have been applied. The quality of the new data sets has allowed to settle the original question posed in B98. The trend of $W_{\\rm HB}$ with metallicity showed a dispersion that was larger than formal measurement errors, a fact that can now be attributed to the inhomogeneous data sources employed.  The dispersion of $W_{\\rm HB}$ around the mean trend with metallicity, is now compatible with the error bar on the data points. This means that (a) one can reverse the argument and use $W_{\\rm HB}$ to have a first-order guess on a cluster metallicity (Sect.~\\ref{sec:metindex}), and (b) that whatever other variables influence the parameter, they must have a well-defined dependence on [Fe/H] (including zero dependence).  The second point was investigated in some detail by comparing the data to two independent sets of theoretical calculations. First, it was noted that the observed trend of $ W_{\\rm HB} $ with metallicity (which is stronger for the $B-V$ color) is well reproduced by both isochrone sets, although a zero-point problem in the $(V-I)$ color-temperature transformations seems to be present.  In order to test the dependence on all other variables influencing $W_{\\rm HB}$, their values were varied within reasonable limits. For any variable $x$, we checked the $W_{\\rm HB}$ vs. $x$ trend at several fixed metallicities, and noticed that these trends can be well approximated with linear relations. This means that we could use the slope $\\Delta W_{\\rm HB}/\\Delta x$ to rank the relative dependencies. We concluded that the influence of the helium content is negligible, and that that of the mixing-length parameter $\\alpha$ is much stronger than that of the age (see Table~\\ref{tab_derivate}). This fact was then exploited to investigate a long-standing theoretical problem of stellar evolution, whether $\\alpha$ should be varied with metallicity or not. From our comparisons, {\\em we find no trend of $\\alpha$ with the metal content of a cluster}.  With respect to the color-$T_{\\rm eff}$ transformations, a test was made by adopting either the Alonso et al. (\\cite{alonsoEtal99}) or the Kurucz (\\cite{kurukz92}) relations. This showed that, {\\em potentially}, the $B-V$ transformations could make a much larger difference than the $V-I$ transformations. However, since actually the theoretical $V-I$ colors show the greater problems, the agreement between the two transformations implies that a deep revision for them is needed.  Finally, we examined the question of whether models with enhanced $\\alpha$-elements better reproduce the CMD morphology of GGCs. The conclusion is that, unfortunately, the existence of two discrepant metallicity scales leaves this question open. The best results are obtained in the $\\alpha$-enhancement scenario, and adopting the ZW84 metallicity scale. However, one could well adopt the CG97 scale and claim that the $B-V$ color-transformations are wrong, since we have seen how much they change from author to author. Thus, we must conclude that there is an urgent need for an independent metallicity scale (which also involves the question of galactic chemical evolution models, see Saviane \\& Rosenberg \\cite{sr99})."
    },
    "0207/cond-mat0207086_arXiv.txt": {
        "abstract": "Electrostatic correlations play an important role in physics, chemistry and biology. In plasmas they result in thermodynamic instability similar to the liquid-gas phase transition of simple molecular fluids.  For charged colloidal suspensions the electrostatic correlations are responsible for screening and colloidal charge renormalization. In aqueous solutions containing multivalent counterions they can lead to charge inversion and flocculation. In biological systems the  correlations account for the organization of cytoskeleton and the compaction of genetic material. In spite of their ubiquity, the true importance of electrostatic correlations has become fully appreciated only quite recently. In this paper, I will review the thermodynamic consequences of electrostatic correlations in a variety of systems ranging from classical plasmas to molecular biology. ",
        "introduction": "Although the liquid state theorists have become quite accustomed to look at the correlation functions for simple and complex fluids, many of the thermodynamic consequences of correlations have not been fully appreciated.  To some extent this is the result of the success of mean-field theories for simple fluids.  The strategy of applying the  mean-field approximation to the many body problems goes all the way to the pioneering work of van der Waals on liquid-gas phase separation~\\cite{Wa73}.  The success of this theory and its physical transparency has set the stage for future applications of mean-field ideas.  These came in the form of the Curie-Weiss theory of magnetism~\\cite{We07} and the Gouy-Chapman~\\cite{Go10,Ch13} theory of diffuse ionic layers.  There are, however, some very familiar  systems in which the mean-field contribution to the free energy is identically  zero.  One such system is the classical two component plasma of positive and negative ions. In order for the thermodynamic limit to exist, charge neutrality constraint must be imposed.  However, for a bulk charge-neutral system the average electrostatic potential is zero, which means that the mean-field contribution to the total free energy vanishes.  Thus, the electrostatic free energy of a two component plasma is entirely due to positional correlations between the positive and negative ions. At low temperatures these correlations become so strong as to lead to  phase transition in which the plasma separates into two coexisting high and low density phases~\\cite{FiLe93,LeFi96}.  Polar fluids provide another example of a system in which the electrostatic correlations strongly affect the thermodynamics. Perhaps the simplest model of a polar fluid is a system of dipolar hard spheres $(DHS)$. The phase structure of $DHS$ is quite interesting and deserves a separate review~\\cite{GrDi94,GrDi96,GrDi97}. Here we shall confine our attention to the low density disordered fluid phase.  Since the average electric field inside a dipolar fluid is zero, it is evident that the mean-field contribution to the free energy also vanishes, and all the non-trivial thermodynamics is, once again, the result of electrostatic correlations.  The thermodynamics of $DHS$ is particularly tricky because of the unscreened long range interactions.  In fact the very existence of the thermodynamic limit for $DHS$ has been proven only recently~\\cite{BaGrWi98}.  Nevertheless, it has been taken for granted that if the temperature is sufficiently low, the $DHS$ will phase separate into coexisting liquid and gas phases.  Indeed, all the theories have been predicting exactly this kind of behavior~\\cite{GePi70,RuStHo73}. It came, therefore, as a great surprise when the simulations in the early $90's$ failed to locate the anticipated liquid-gas critical point~\\cite{WeLe93,Ca93,LeSm93}. Instead as the temperature was lowered, the simulations found chains of aligned dipoles. Formation of weakly interacting chains,  a consequence of strong positional and directional correlations between the dipolar particles, prevented the liquid-gas phase separation from taking place~\\cite{Le99,Se96,Ro96,TaTeOs97}.  Electrostatic correlations are also crucial in charged colloidal suspensions~\\cite{RoHa97,HaLo00}. In these systems the correlations come on two different levels. First, there are very strong positional correlations between the highly charged colloidal particles and their counterions. These correlations lead to charge renormalization~\\cite{AlChGr84} and to screening of Coulomb interactions between the colloidal particles. In water with monovalent counterions, charge renormalization stabilizes colloidal suspension against phase separation~\\cite{LeBaTa98,DiBaLe01}. In the presence of multivalent counterions, however, the layers of condensed counterions on different colloids can become strongly correlated, leading to a net attraction between the like-charged colloids~\\cite{Pa80,GuJoWe84} and to the phase separation~\\cite{LiLo99,Lo00}.  A similar kind of behavior was also observed in a number of important biological systems. Thus, it was noted that like-charged macromolecules can attract each other in solutions containing multivalent counterions. This attraction manifests itself in {\\it in vitro} formation of toroidal aggregates of concentrated DNA~\\cite{Bl91,Bl97}, similar to the one found in  bacteriophage heads~\\cite{Kl67}, and in  bundle formation  of  F-actin and tobacco mosaic virus~\\cite{Ta96}. A number of models have been suggested to explain these curious phenomena.  The fundamental ingredient in all of these models are the electrostatic correlations~\\cite{HaLi97,ArStLe99,Sh99}.   Although the phenomena described above are quite complex, we can get a long way towards understanding them by considering some surprisingly simple models and theories. In fact, we shall demonstrate that a lot of the  physics of electrostatic correlations is contained within the Debye-H\\\"uckel theory~\\cite{DeHu23} introduced $80$ years ago as a way of accounting for the unusual thermodynamic properties of strong electrolytes.  Consideration of only simple physical theories in this review is partially pedagogical, designed for a  broad audience not necessarily familiar with the complex machinery of correlation functions and field theories of the modern statistical mechanics. For Coulomb systems there is, however, an additional benefit. It is often found that the more sophisticated theories fail when applied to  strongly correlated charged fluids.  For example, the field theoretic calculations of Netz and Orland~\\cite{NeOr99} find that for charge asymmetric $(z:1)$ electrolytes the reduced critical temperature is a strongly increasing function of charge asymmetry. A completely  opposite behavior is observed in  computer simulations, the critical temperature decreases and the critical density increases with the charge asymmetry. In fact, the field theoretic predictions for the critical temperature of asymmetric electrolytes are so far-off that they have to be divided by a factor of six just to make them fit on the same graph (Fig. \\ref{Fig0}) with the results of simulations and of concurrent theories! The dramatic failure of field theoretic calculations can be attributed to their intrinsically perturbative nature. Similarly the integral equations, which have proven to be  very successful for  simple molecular fluids, fail to even converge for strongly asymmetric electrolytes.  Furthermore, it has  been known for a long time that the Hypernetted Chain Equation (HNC)  which is often used to study the Coulomb systems~\\cite{Be86a}, does not posses a true critical region, but only a ``no solution zone'' on the border of which  compressibility goes to zero with a square root singularity~\\cite{FiFi81,Be93}. This is a completely wrong behavior, since the compressibility must diverge at the critical point.   All these should be contrasted with the physically based Debye-H\\\"uckel-like theories, which are in qualitative and often in  quantitative agreement with the simulations and experiments, Fig. \\ref{Fig0}. \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig0a.ps} \\includegraphics[width=8cm]{fig0b.ps} \\end{center} \\caption{Estimates of the reduced critical temperature, $T_c^*$, and density, $\\rho_c^*$, for the $(z:1)$, charge-asymmetric (but equisized) primitive model showing, as labeled, the predictions of pure Debye-H\\\"uckel theory (without hard cores), (b) the MSA, (c) the SPB approximation~\\cite{SaBhOu98}, (d) the Netz-Orland field-theoretic treatment~\\cite{NeOr99} which, for $T_c^*$, have been divided by factors of $6$ and $12$ in order to bring them on to the plot (see labels NO$^\\prime$ and NO$^\\prime$$^\\prime$ respectively), and (e) the DHBjCI theory (following the Fisher-Levin approach~\\cite{FiLe93,LeFi96}), point $z=3$ is a preliminary result. The large open circles for $z=2$ and $3$ represent, the Monte Carlo simulations of Panagiotopoulos and Fisher~\\cite{PaFi02}, while the estimate for $T_c^*(4)$ follows from Camp and Patey~\\cite{CaPa99}. (After S. Banerjee and M. E. Fisher, to be published).} \\label{Fig0} \\end{figure} For some important many body systems, however, the integral equations provide the most accurate results.  For example,  predictions for the electrostatic free energy of the one component plasma obtained using the $HNC$ equation are in excellent agreement with the Monte Carlo simulations~\\cite{Ng74}. The integral equations were also the first to account for the correlation induced attraction between the like-charged macromolecules~\\cite{Pa80,KjMa84,KjMa86}. In general, as long as one stays away from the phase transitions, the integral equations provide one of the sharpest tools available to a statistical physicist or chemist. Unfortunately, the approximations involved in constructing the integral equations are not very clear.  There exists a great variety of closures to the Ornstein-Zernike equation, each working  well for specific kind of problems. Because of their complexity, I will not talk about the integral equations in this review, referring the interested reader to the literature~\\cite{HaMc76}. ",
        "conclusions": "We have explored the role of electrostatic correlations in systems ranging from classical plasmas to molecular biology. We saw how the positional correlations between the ions of an electrolyte can result in a thermodynamic instability.  We also saw how the strong correlations between the polyions and the counterions lead to colloidal charge renormalization which stabilizes deionized suspensions against a phase separation.  For two-dimensional plasmas electrostatic correlations are responsible for the metal-insulator transition. The critical behavior of superfluid $^4$He films, the roughening transition of crystal interfaces~\\cite{ChWe76}, the melting of two-dimensional solids, and the criticality of XY magnets~\\cite{Ne83} are all governed by the ``electrostatic'' interactions between the topological defects (charges). Nature has learned to take the full advantage of electrostatic correlations to efficiently package huge amounts of genetic material into tiny regions of space.  Throughout this review we have come to rely on some simple models in order to understand complex physical phenomena. While these models are often sufficient to grasp the underlying physics, it is quite easy to push the models too far. This is particularly the case when one deals with specific structural properties of biomolecules~\\cite{GiRaFi85,PeWaJo98,SpBa98}. As soon as the  length scales on the order of few angstroms become important, approximation of water as a uniform dielectric medium is no longer sufficient~\\cite{AuWe98,AuLoWe96}. Under these conditions reliance on simple models, which treat macromolecules and solvent as dielectrics, is probably no more than a wishful thinking. A careful path must be threaded between the simplification and the over-simplification."
    },
    "0207/astro-ph0207234_arXiv.txt": {
        "abstract": "{A statistical analysis of almost 50\\,000 soft X-ray (SXR) flares observed by GOES during the period 1976--2000 is presented. On the basis of this extensive data set, statistics on temporal properties of soft X-ray flares, such as duration, rise and decay times with regard to the SXR flare classes is presented. Correlations among distinct flare parameters, i.e. SXR peak flux, fluence and characteristic times, and frequency distributions of flare occurrence as function of the peak flux, the fluence and the duration are derived. We discuss the results of the analysis with respect to statistical flare models, the idea of coronal heating by nanoflares, and elaborate on implications of the obtained results on the Neupert effect in solar flares. ",
        "introduction": "With the availability of space-borne instrumentation, observations of solar flare phenomena in X-rays became possible in the 1960s. Disk-integrated soft X-ray emission measurements of the Sun have been collected more or less continuously since 1974 by the National Oceanic and Atmospheric Administration (NOAA), providing an almost unbroken record fully covering solar cycles~21 and~22 and the rising phase of cycle~23 (Garcia 2000).  Statistical investigations on temporal aspects of solar flares observed in various soft X-ray wavelengths have been carried out by Culhane \\& Phillips (1970), Drake (1971), Thomas \\& Teske (1971), Phillips (1972), Datlowe et al. (1974), and Pearce \\& Harrison (1988). In the meantime a wealth of data has accumulated, which makes worthwhile re-investigating the temporal characteristics of soft X-ray (SXR) flares on an extensive statistical basis. In the present analysis we make use of SXR flares observed by GOES during 1976--2000.  Moreover, on the basis of this comprehensive data set, we calculate frequency distributions of SXR flares as function of the peak flux, the fluence, i.e. the integrated flux from the start to the end of a flare, and the event duration. Frequency distributions of flare occurrence are related to the observational expectations from different flare models. Furthermore, they contain information about the possibility of coronal heating by nanoflares. Frequency distributions of solar flares from disk-integrated SXR measurements as function of the peak flux have been carried out by Drake (1971), Lee et al. (1995) and Feldman et al. (1997). However, only in the paper by Drake (1971) are frequency distributions of SXR fluence measurements also presented. Shimizu (1995) used spatially resolved observations from transient SXR brightenings and investigated frequency distributions as function of energy.  Soft X-ray measurements are an important counterpart to observations of flares in hard X-rays (HXR). From several observations it is reported that the SXR light curve has a similar shape as the time integral of the HXR curve. This led to the idea that there is a causal relationship between hard and soft X-ray emission of a flare, the so-called Neupert effect (Neupert 1968; Dennis \\& Zarro 1993). It supports a flare model, known as the thick-target model (Brown 1971), in which the HXR emission is electron-ion bremsstrahlung produced by energetic electrons as they reach the dense layers of the chromosphere. Only a small fraction of the electron beam energy is lost through radiation; most of the loss is due to Coulomb collisions, which serve to heat the ambient plasma. Due to the rapid deposition of energy by the particle beams, the energy cannot be radiated away at a sufficiently high rate and a strong pressure imbalance develops, causing the heated plasma to expand up into the corona (``chromospheric evaporation\"), where this hot dense plasma gives rise to the enhanced SXR emission (e.g., Antonucci et al. 1984; Fisher et al. 1985; Antonucci et al. 1999, and references therein). In this paper we make use of the flare frequency distributions and the correlations among distinct flare parameters to infer information about the validity of the Neupert effect in solar flares by statistical means.  The paper is structured as follows. The data set is described in Sect.~2. In Sect.~3.1 a statistical investigation of temporal properties of SXR flares is presented. Correlations among various flare parameters, such as characteristic times, peak flux and fluence, are analyzed in Sect.~3.2. In Sect~3.3 frequency distributions as function of the peak flux, the fluence and the duration are derived. In Sect.~4 we give a summary and discussion of the main results. Finally, the conclusions are drawn in Sect.~5. ",
        "conclusions": "Frequency distributions of SXR flare occurrence as function of the peak flux, the fluence and the event duration have been calculated. All distributions can be described by power-law functions over several decades. The distributions, derived separately for the times of minimum and maximum solar activity do not reveal any remarkable change in the power-law index, consistent with the predictions of avalanche flare models (e.g., Lu \\& Hamilton 1991; Lu et al. 1993).  Relating the SXR fluence measurements to the total radiated flare energy, the determined power-law index $\\alpha_{\\cal F}$ is also representative of the flare energy distribution. The obtained values of $\\alpha_{\\cal F}$ are 2.03 for the raw fluence data and 1.88 for the background subtracted fluence data. Both values are rather close to the critical value of 2, and no distinct conclusion can be drawn whether small-scale events provide a significant contribution to coronal heating or not. Moreover, deviations from the assumed proportionality between SXR fluence and flare energies would also cause deviations from the one-to-one correspondence of the respective power-law indices. If the slope of energy versus fluence in log-log space is less than one, then the power-law index of the energy distribution would be smaller than those of the fluence distribution, and vice versa.   The power-law index of SXR peak flux distributions (see Drake 1971; Lee et al. 1995; Feldman et al. 1997; this paper) is significantly larger than those reported for HXR fluence distributions (cf. Lee et al. 1993, and references therein), statistical evidence that the Neupert effect in its commonly stated form relating the X-ray emissions (Eq.~\\ref{EqNeup}) is not valid for the bulk of flares. However, this outcome does not necessarily mean that the Neupert effect does not work for the more fundamental relationship between the energies (Eq.~\\ref{NeupEn}). Depending on the validity of the Neupert effect for the bulk of solar flares, the differences in the power-law indices contain information upon additional energy sources for the SXR-emitting plasma, or upon the amount of SXR emission per HXR electron as a function of the flare size. From the present analysis, based only on SXR data, we cannot distinguish between these different possibilities. Moreover, a mixing of both cases might exist, making a distinction even more difficult. Therefore, a detailed statistical analysis of related SXR and HXR flares is in preparation in order to obtain deeper insight into the Neupert effect."
    },
    "0207/astro-ph0207002_arXiv.txt": {
        "abstract": "High-resolution ultraviolet observations of nearby bright and faint stars are required to evaluate changes in the past and future galactic environments of the Sun, and the possibly impact of these changes on the interplanetary environment at 1 AU (around the Earth).  The boundary conditions of the heliosphere and interplanetary environments are constrained by the characteristics of the surrounding interstellar material (ISM), which changes on timescales of 10$^3$--10$^5$ years. An increase in the density of the interstellar cloud surrounding the solar system to 10 cm$^{-3}$ decreases the heliosphere radius by about an order of magnitude.  UV observations of nearby stars at high spectral resolution ($>$300,000) and high signal-to-noise are required to evaluate future modifications to heliosphere properties by the ISM. ",
        "introduction": "High-resolution ultraviolet observations of nearby bright and faint stars are required to evaluate changes in the past and future galactic environments of the Sun, and the possibly impact of these changes on the interplanetary environment at 1 AU (around the Earth) and in the outer solar system. ",
        "conclusions": "The morphology of nearby ISM ($<$ 30 pc), and interstellar pressure variations affecting the heliosphere and interplanetary medium surrounding the Earth, {\\it can only be found} from high-resolution ($\\lambda$/$\\delta \\lambda$$>$300,000) observations of interstellar absorption lines in the ultraviolet wavelength band (912 --- 3000 Ang).  The relative simplicity of sightlines to nearby stars, combined with high-resolution high signal-to-noise UV data, provides the best opportunity for obtaining accurate abundances and cloud properties for interstellar clouds."
    },
    "0207/astro-ph0207372_arXiv.txt": {
        "abstract": "Although many methods of detecting extra-solar planets have been proposed and successful implementation of some of these methods enabled a rapidly increasing number of exoplanet detections, little has been discussed about the method of detecting satellites around exoplanets.  In this paper, we test the feasibility of detecting satellites of exoplanets via microlensing. For this purpose, we investigate the effect of satellites in the magnification pattern near the region of the planet-induced perturbations by performing realistic simulations of Galactic bulge microlensing events. From this investigation, we find that although satellites can often cause alterations of magnification patterns, detecting satellite signals in  lensing light curves will be very difficult because the signals are seriously smeared out by the severe finite source effect even for events involved with source stars with small angular radii. ",
        "introduction": "Various methods have been proposed to search for extrasolar planets (exoplanets).  These methods include the pulsar timing analysis, direct imaging, accurate measurement of astrometric displacements, radial velocity measurement, planetary transit, and gravitational microlensing [see the review of \\citet{perryman00}].  Since the first detection of an exoplanet around the pulsar PSR 1257+12 \\citep{wolszczan92}, nearly 100 exoplanets have been identified (http://exoplanets.org), mostly by the radial velocity method \\citep{mayor95}.   However, little has been discussed about the method of detecting satellites around exoplanets.  This is mainly because it is thought to be too premature to detect satellites given the difficulties of detecting exoplanets.  Currently, the only promising technique proposed to detect satellites is the transit method, where satellites are detected either by direct satellite transit or through perturbations in the transit timing of the satellite-hosting planet \\citep{sartoretti99}.   In this paper, we investigate the feasibility of detecting satellites of exoplanets via microlensing.  Detection of a low-mass companion by using microlensing is possible because the companion can induce noticeable anomalies in the resulting lensing light curves \\citep{mao91, gould92}. The microlensing method has an important advantage in detecting very low-mass companions over other methods because the strength of the companion's signal depends weakly on the companion/primary mass ratio although the duration of the signal becomes shorter with the decrease of the mass ratio.  Then, if lensing events are monitored with a high enough frequency, it may be possible to detect not only planets but also their satellites.  Such a frequent lensing monitoring program in space was recently proposed by \\citet{bennett00}.   The paper is organized as follows.  In \\S\\ 2, we describe the microlensing basics of multiple-lens systems.  In \\S\\ 3, we investigate the feasibility of satellite detections by carrying out realistic simulations of Galactic bulge microlensing events caused by an example lens system having a planet and a satellite.  We summarize the results and conclude in \\S\\ 4. ",
        "conclusions": "We have tested the feasibility of detecting satellites by using microlensing. For this purpose, we have investigated the effect of satellites on the magnification pattern near the region of planet-induced perturbations by carrying out realistic simulations of Galactic bulge microlensing events. From this investigation, we find that although satellites can often affect the maginification patterns, detecting satellite signals in the lensing light curves will be very difficult because the signals are seriously smeared out by the severe finite source effect.   \\bigskip \\bigskip We would like to thank J.\\ H.\\ An for making useful comments about the work. This work was supported by a grant (2001-DS0074) from Korea Research Foundation (KRF)."
    },
    "0207/astro-ph0207658_arXiv.txt": {
        "abstract": "A sophisticated approach to condensate opacity is required to properly model the atmospheres of L and T dwarfs.  Here we review different models for the treatment of condensates in brown dwarf atmospheres.  We conclude that models which include both particle sedimentation and upwards transport of condensate (both gas and particles) provide the best fit for the L dwarf colors.  While a globally uniform cloud model fits the L dwarf data, it turns to the blue in $J-K$ too slowly to fit the T dwarfs. Models which include local clearings in the global cloud deck, similar to Jupiter's prominent five-micron hot spots, better reproduce the available photometric data and also account for the observed resurgence of FeH absorption in early type T dwarfs. ",
        "introduction": "Long before the first discoveries of brown dwarfs, it was recognized that condensates would play a critical role in controlling their atmospheric opacity, at least in certain effective temperature ranges (Stevenson 1986; Lunine et al. 1989).  It was also evident that the correct choice for the vertical distribution of the condensates was not obvious (Lunine et al. 1989). Condensates might be well mixed in the atmosphere above their condensation level, or might coalesce into large particles, fall below their condensation level and be removed from the atmosphere.  Of course many intermediate cases are possible as well.  After the discovery of what came to be known as the L and T dwarfs, modelers initially focused on simple end cases.  Condensates were either assumed to have either completely settled from the atmosphere or else were mixed uniformly throughout the observable atmosphere. The prior approach works reasonably well for objects like Gl 229 B (Allard et al. 1996; Marley et al. 1996; Saumon et al. 2000; Tsuji et al. 1996), while the latter works for very late M and early L dwarfs like Kelu 1 (Ruiz, Leggett \\& Allard 1997).  Neither approach, however, could adequately reproduce the colors, let alone the spectra of the latest Ls or the `transition' late L/early T objects, like SDSS 1254 (Fig. 1). \\begin{figure} \\plotone{f1.eps} \\caption{Near-infrared color-magnitude diagram of M, L, and T dwarfs. The absolute $J$ magnitudes and $J-K_s$ infrared colors are shown for a sample of M (filled squares), L (filled triangles), and T (filled circles) dwarfs with known parallaxes. The positions of 2MASS 0559 and SDSS 1254 are indicated. The predicted colors and magnitudes for the DUSTY (dashed line), clear (left thin line), and cloudy (right thin line; $f_{\\rm sed} = 3$) (the variable formerly known as $f_{\\rm rain}$) atmosphere models are plotted as a function of $T_{\\rm eff}$ at constant gravity, $g = 10^5\\,\\rm cm\\, s^{-2}$ (typical for very low mass main-sequence stars and evolved brown dwarfs). Connecting the cloudy and clear tracks are the predicted fluxes for partly cloudy models at $T_{\\rm eff} = $800, 1000, 1200, 1400, 1600, and 1800 K. The circles indicate the cloud coverage fraction in steps of 20\\%. The apparent evolutionary track of brown dwarfs based on the empirical data is indicated by the thickened line, which crosses from the cloudy to clear track at $T_{\\rm eff} \\sim 1200\\,\\rm K$. Figure adapted from Burgasser et al. (2002).} \\end{figure} Models (Fig. 1) in which the condensates are absent from the atmosphere produce $J-K$ colors that are much bluer than the observed L dwarfs.  The blue color arises from water, pressure-induced $\\rm H_2$, and for lower effective temperatures ($T_{\\rm eff} < 1400\\,\\rm K$) $\\rm CH_4$ absorption in $K$ band.  Models with condensates distributed uniformly through the atmosphere match the early Ls in which the condensate column optical depth is small, but are much redder than the colors of the later L dwarfs.  The reason is that as the cloud deck falls progressively deeper in the atmosphere the column abundance of dust gets progressively larger. Since the cloud is forming at higher air densities the abundance of condensates to be mixed upwards from cloud base is larger with falling effective temperature.  Thus an outside observer sees an ever increasing dust optical depth until the observer is effectively looking at a dirt-filled atmosphere. Like blackbodies, such objects becomes progressively redder in $J-K$ as they cool. These models cannot reproduce the observed transition from red to blue in $J-K$ between the L and T dwarfs.  An obvious shortcoming of the first generation models is that none of them employed cloud decks like those seen throughout the solar system. Condensates in planetary atmospheres are generally neither completely settled out of the atmosphere nor distributed uniformly to the top of the atmosphere.  Rather they tend to exist in horizontally-extended cloud decks.  The $\\rm H_2SO_4$ clouds of Venus, stratiform water clouds on Earth, the ammonia cloud decks on Jupiter, and the methane clouds in Uranus and Neptune are all examples of this.  To demonstrate that a vertically constrained cloud deck would qualitatively be more in line with the available data, Marley (2000) constructed a simple model in which all clouds were 1 scale height thick.  He showed that such a model produced less red $J-K$ colors than well mixed models and also naturally explained the turnover in $J-K$ color from the red L dwarfs to the blue T dwarfs.  As the cloud forms progressively lower in the atmosphere as the object cools, the cloud disappears while the clear atmosphere above takes on the appearance of the cloud-free models.  Tsuji (2001) also used a simple model with a finite-thickness cloud to make the same point.  Given the apparent importance of correctly modeling cloud behavior, Ackerman \\& Marley (2001) developed a more rigorous cloud model for use in substellar atmosphere models.  The model attempts to capture the key `zeroth order' physics which influence the vertical condensate abundance and size profiles in a realistic (but 1-D) atmosphere.  This model was used by Marley et al. (2002) and Burgasser et al. (2002) to model the atmospheres of L and T dwarfs.  Other workers have recently turned to the work of Rossow (1978), originally developed to study the microphysics of clouds in planetary atmospheres, to model cloud behavior in substellar atmospheres.  Finally Tsuji (2002) has further developed his model in which the cloud top temperature is an adjustable parameter.  In Section 2 we review and compare these cloud models.  In Section 3 we summarize the results of Burgasser et al. in applying the Ackerman \\& Marley cloud model to L and T dwarfs and consider the role of dynamically-induced holes in the global cloud coverage. ",
        "conclusions": "It appears that a complete description of the behavior of L and T dwarfs will require a thorough understanding of the behavior of condensates in their atmospheres.  In doing so models must simultaneously and self-consistently account for a number of influences including particle nucleation, sedimentation, and vertical transport. Other influences, including large scale atmospheric dynamics which may be responsible for cloud patchiness, may also be important. Given this daunting task, it is worth remembering that even after half a century of dedicated effort, such key properties as the particle sizes and vertical structure of most of the cloud decks in the solar system are still poorly  known. The mechanisms responsible for those characteristics which are constrained are themselves only partly understood. Despite these challenges a coherent story for the behavior of L and T dwarf condensates is emerging, although our understanding is certainly still not complete. Clouds vary in time and in space and we should perhaps not be suprised that weather prediction is a challenging business."
    },
    "0207/astro-ph0207144_arXiv.txt": {
        "abstract": "A sample of bright contact binary stars (W~UMa-type or EW, and related: with $\\beta$~Lyr light curves, EB, and ellipsoidal, ELL -- in effect, all but the detached, EA), to the limit of $V^{max} = 7.5$ magnitude is deemed to include all discoverable short-period ($P < 1$ days) binaries with photometric variation larger than about 0.05 magnitude. Of the 32 systems in the final sample, 11 systems have been discovered by the Hipparcos satellite. The combined spatial density is evaluated at $(1.02 \\pm 0.24) \\times 10^{-5}$ pc$^{-3}$. The Relative Frequency of Occurrence (RFO), defined in relation to the Main Sequence stars, depends on the luminosity. An assumption of RFO $\\simeq$ 1/500 for $M_V > +1.5$ is consistent with the data, although the number statistics is poor with the resulting uncertainty in the spatial density and the RFO by a factor of about two. The RFO rapidly decreases for brighter binaries to a level of 1/5,000 for $M_V < +1.5$ and to 1/30,000 for $M_V < +0.5$. The high RFO of 1/130, previously determined from the deep OGLE--I sample of Disk Population W~UMa-type systems towards Baade's Window, is inconsistent with and unconfirmed by the new results. Possible reasons for the large discrepancy are discussed. They include several observational effects, but also a possibility of a genuine increase in the contact-binary density in the central parts of the Galaxy. ",
        "introduction": "\\label{intro}  Our views about the spatial density of contact, W~UMa-type binaries have meandered considerably during the recent half century. \\citet{sha48} was the first to point out that ``W~Ursae Majoris variables ... are not only the most numerous of eclipsing stars ... but ... more numerous than all other variable stars combined''. From the numbers given by Shapley, one can infer that he estimated that perhaps as many as about one percent of all solar-type stars are W~UMa-type variables.  Several attempts have been made to put the statement of Shapley in a quantitative way, with diverse results. \\citet{pop64} estimated the local density at $2 \\times 10^{-5}$ pc$^{-3}$, while \\citet{kra67} estimated it to be some 20 times lower, $10^{-6}$ pc$^{-3}$. Then \\citet{vnV75} found a much larger number, $11 \\times 10^{-5}$ pc$^{-3}$, which corresponds to about one percent of all stars being W~UMa-type binaries. Subsequently \\citet{due84} derived a value of $\\simeq 10^{-5}$ pc$^{-3}$, equivalent to the Relative Frequency of Occurrence (RFO) of $\\simeq 0.001$ (or 0.1 percent), and this value remained in popular use for some time. A similar, simplified, and thus rather convincing (based on naked eye objects) estimate of the frequency was given by \\citet{ruc93}, RFO~$\\simeq 0.0005 - 0.002$, that is one W~UMa binary per 500 -- 2000 ordinary stars.  The newest investigations, based on much larger statistical samples, suggested that the RFO is perhaps as high as suggested by Shapley, i.e.\\ some five to ten times higher than estimated by \\citet{due84}. During the last two decades, several EW systems have been discovered in old open clusters \\citep{KR93,RK94,ruc98b}. A combined approximate estimate of \\citet{ruc94} for a few old open clusters gave the RFO of $\\simeq 275 \\pm 75$ ordinary dwarfs per one EW system. Later, in a combined analysis of several open clusters, \\citet{ruc98b} showed that the RFO evolves with the age of the stellar system; the number of EW systems increases from the level of one per a thousand at the age of about 0.8 Gyr to the level of one per some 200 -- 300 dwarfs at the age of 5 -- 7 Gyr. The RFO was found also very high in globular clusters \\citet{ruc00}; however, Population~II contact binaries are of no relevance to the present study which concentrates on the Population~I objects, primarily in the solar neighborhood.  The highest RFO was estimated for the Galactic Disk, as probed by the narrow, deep pencil beam of the OGLE--I survey in the direction of Baade's Window \\citep{ruc95,ruc97a,ruc97b,ruc98a,ruc98b}. This finding was in a basic accordance with the view of the increased RFO from the age of the oldest open clusters to the age of the Galactic Disk Field of about 11 Gyr \\citep{bin00}. The great advantage of the OGLE--I result over all previous spatial density estimates was not only in the size and uniformity of the sample (a few hundred contact systems discovered in the same survey), but -- mainly -- in that the RFO estimate was based on the {\\it volume-limited samples\\/}, complete to $M_V \\simeq 5.5$ to the distance of 3 kpc, and (with better statistics for brighter systems) to $M_V \\simeq 4.2$ to the distance of 5 kpc. Large numbers of binaries in the OGLE--I sample permitted to evaluate the RFO on the per-$M_V$-bin way, in place of the previous estimates based on data averaged for all accessible spectral types, between late A to early K. Such comparisons through the luminosity function led to a statistically well established -- and high -- apparent frequency at the level of $\\simeq 1/130$ \\citep{ruc98b}. A clear drop in numbers for brighter, longer-period contact systems, with a sharp cut-off at $P \\simeq 1.3-1.5$ days, was also noted \\citep{ruc98a, ruc98b}.  The OGLE--I result would suggest a return to Shapley's value of the RFO at the level of one percent or higher. However, this would mean a rather large disagreement with the relatively firm result of \\citet{due84} for the solar neighborhood. Can we exclude a systematic error in the OGLE--I result? Even if based on the best {\\it number statistics\\/}, the OGLE--I may have {\\it systematic\\/} biases. After finding that the new 7.5 magnitude-limit sample indeed does not agree with the OGLE--I result, we attempt to give possible reasons for the discrepancy in Section~\\ref{discus}.  Sections~\\ref{criteria} -- \\ref{PC} describe construction and properties of the sample used in a new estimate of the spatial density. This estimate is through the luminosity function, as described in Section~\\ref{LF}. The related ``period function'' (the number of systems per unit of volume, per period interval) is described in Section~\\ref{PF}. The amplitude distribution for the sample is discussed in Section~\\ref{ampl}. Section~\\ref{discus} discusses the major discrepancy in the density spatial estimates between the current sample and the OGLE--I sample in the direction of Baade's Window \\citep{ruc98b}. Section~\\ref{future} looks into the future of the spatial density estimates for contact binaries. Section~\\ref{concl} summarizes the results of the paper. Appendix~\\ref{appendix} gives brief descriptions of the individual binaries, including those which have been excluded from the sample in the last stages of its definition. ",
        "conclusions": "\\label{concl}  The current sample of contact (EW) and related (EB and ELL) systems with orbital periods shorter than one day and selected from among bright stars to $V^{max}=7.5$ (with a large fraction detected by the Hipparcos satellite) is expected to be complete in the sense that it contains all binaries showing photometric light variation larger than about 0.05 magnitude. This level of photometric accuracy implies that -- according to predictions of \\citet{ruc01} (Section~3.3 and Figure~4 there) -- only about 10 -- 15 percent of variables remain undiscovered due to low orbital inclination angle. Therefore, even if current statistics do not take these undetectable systems into account, the spatial-density estimate of $\\rho = (1.0 \\pm 0.3) \\times 10^{-5}$ pc$^{-3}$ is expected to be sound and consistent with the Relative Frequency of Occurrence, RFO $\\simeq$ 1/500, when compared with the Main Sequence stars with $M_V > +1.5$. While both quantities, $\\rho$ and RFO, are still uncertain by a factor of about two times, mostly because of the small number statistics, the estimates should be relatively free of systematic errors. The spatial density estimate derived here is very close to that of \\citet{due84}, but does not agree with the high value estimated on the basis of the Disk Population Baade's Window sample based on the OGLE--I data. As discussed in Section~\\ref{discus}, while the discrepancy is most likely due to observational effects which produced the high value for the OGLE--I sample, we cannot exclude a possibility that the population of contact binaries observed towards the Center of the Galaxy is different from the one in the solar neighborhood."
    },
    "0207/astro-ph0207294.txt": {
        "abstract": "We report precise Doppler shift measurements of 55 Cancri (G8V) obtained from 1989 to 2002 at Lick Observatory.  The velocities reveal evidence for an outer planetary companion to 55 Cancri orbiting at 5.5 AU.  The velocities also confirm a second, inner planet at 0.11 AU. The outer planet is the first extrasolar planet found that orbits near or beyond the orbit of Jupiter. It was drawn from a sample of $\\sim$50 stars observed with sufficient duration and quality to detect a giant planet at 5 AU, implying that such planets are not rare.  The properties of this jupiter analog may be compared directly to those of the Jovian planets in our Solar System.  Its eccentricity is modest, $e$=0.16, compared with $e$=0.05 for both Jupiter and Saturn.  Its mass is at least 4.0 \\mjup (\\msini).  The two planets do not perturb each other significantly.  Moreover, a third planet of sub--Jupiter mass could easily survive in between these two known planets.  Indeed a third periodicity remains in the velocity measurements with P = 44.3 d and a semi--amplitude of 13 \\ms.  This periodicity is caused either by a third planet at $a$=0.24 AU or by inhomogeneities on the stellar surface that rotates with period 42 d. The planet interpretation is more likely, as the stellar surface is quiet, exhibiting $\\log(R'_{\\rm HK}) = -5.0$ and brightness variations less than 1 millimag, and any hypothetical surface inhomogeneity would have to persist in longitude for 14 yr.  Even with all three planets, an additional planet of terrestrial--mass could orbit stably at $\\sim$1 AU. The star 55 Cancri is apparently a normal, middle--aged main sequence star with a mass of 0.95 \\msune, rich in heavy elements ([Fe/H] = +0.27).  This high metallicity raises the issue of the relationship between its age, rotation, and chromosphere. ",
        "introduction": "\\label{intro}  Rarely in modern astrophysics does a nearby star attract intense scrutiny on three observational fronts.  The main sequence star 55 Cancri (= $\\rho^1$ Cnc = HD 75732 = HIP 43587 = HR 3522, G8V) has been examined for its extreme abundances of chemical elements, its close--in orbiting planet, and its controversial disk of dust.  These three putative properties are plausibly linked together by the formation and evolution of planetary systems making the system rich with implications.  The metal--rich nature of 55 Cnc was first noticed by H.Spinrad and B.Taylor who alerted Greenstein and Oinas (1968).  They all noted its unusually high abundance of iron and carbon relative to that in the Sun.  The iron lines and CN molecular absorption spectral feature were particularly prominent in blue photographic spectra.  These results were confirmed by Taylor (1970) and indeed, Bell and Branch (1976) reported that carbon was yet further enhanced over iron, [C/Fe]=+0.15 .  Later spectral analyses of 55 Cnc have confirmed its high metallicity (Cayrel de Strobel et al. 1992, 2001; Taylor 1996, Gonzalez \\& Vanture 1998, Feltzing and Gonzalez 2001) with estimates of (logarithmic) iron abundance relative to the Sun ranging from [Fe/H]=+0.1--0.5 .  Thus 55 Cnc is regarded as a rare ``super metal rich'' main sequence star, but confusion still remains about the interpretation of SMR stars (Taylor 2002, Reid 2002).  A planet was reported around 55 Cnc having an orbital period of 14.65 d, an implied orbital radius of 0.11 AU, and a minimum mass of, \\msini = 0.84 \\mjup (Butler et al.~1997).  It was the fourth extrasolar planet discovered, coming after the planets around 51 Peg, 70 Vir, and 47 UMa.  The velocity residuals to the orbital fit of 55 Cnc exhibited a monotonic increase of 90 \\ms from 1989--1995 followed by an apparent decrease in 1996.  Butler et al.~noted that these residuals constrained a possible second planet to have a period, $P>$8 yr, and a mass, \\msini $>$ 5 \\mjup.  The decrease in the velocity residuals continued during 1997 (Marcy \\& Butler 1998), supporting the planetary interpretation.  However, without a full orbital period nor a Keplerian velocity curve, the possibility of stellar activity as the cause of the residuals could not be excluded. This star joined 51 Peg (Mayor \\& Queloz 1995, Marcy et al.~1997) as members of a growing class of planet--bearing stars that have metallicity well above solar (Gonzalez 1998, Barnes 2001, Santos 2000, Butler et al.~2000).  A third issue arose for 55 Cnc when Dominik et al.~(1998) presented evidence for a Vega--like dust disk based on {\\it Infrared Space Observatory} (ISO) measurements between 25 $\\mu$m and 180 $\\mu$m. They detected the photosphere at 25 $\\mu$m and excesses at the higher wavelengths.  Trilling \\& Brown (1998) reported resolving the disk out to 3.2 arcsec (40 AU) with near--infrared coronographic images. Controversy over the disk detections arose when Jayawardhana et al. (2000) found the submillimeter emission to be lower by a factor of 100 than that expected from the disk reported by Trilling \\& Brown. Equally troubling were observations by the NICMOS near-infrared camera on the {\\it Hubble Space Telescope} (Schneider et al.~2001) which imposed an upper limit on the near--infrared flux that was 10 times lower than that reported by Trilling \\& Brown.  A possible resolution of the discrepancies was provided by Jayawardhana et al.~(2002) who found three faint sources of sub--mm emission that were located near but not centered on 55 Cnc, implying that past detections of IR flux might have come from background objects.  The NICMOS upper limit, the upper limit to the sub--mm flux, and the detection of background field sources suggest that no disk has been detected around 55 Cnc.  Indeed, Habing et al.~(2001) discuss the non--negligible probability of spurious detections of disks by ISO caused by fluctuations and by background field sources.  The star 55 Cnc is also a visual binary, with a common proper motion companion 7 magnitudes fainter (V=13, I=10.2), separated by 85 arcsec corresponding to 1100 AU projected on the sky (Hoffleit 1982).  We have measured the barycentric radial velocities for components A and B to be 27.3$\\pm$0.3 and 27.4$\\pm$0.3 \\kms respectively (Nidever et al.~2002).  Thus, the two common proper motion stars are indeed likely bound.  Their common space motion is similar to that of the Hyades supercluster (Eggen 1993).  This paper will be concerned only with component A that we will refer to as ``55 Cnc'' for which we report continued radial velocities measurements, extending from 1989 to 2002.4.  In section 2 we provide a update on the properties of the star, especially its mass, metallicity, and chromospheric activity level.  In section 3 we present all the radial velocity measurements and section 4 contains the orbital fit to two planets.  In the remaining sections we study the possibility of additional planets and the gravitational dynamics between the planets.  \\eject ",
        "conclusions": "All $\\sim$90 previously reported extrasolar planets reside in smaller orbits than that of Jupiter in our Solar System.  The duration of Doppler searches for extrasolar planets had not been long enough to capture an entire orbital period of 12 years for planets at 5 AU.  Indeed all previously known planets are known to have semimajor axes less than 4 AU, well within Jupiter's orbital distance.  Moreover, a large majority of the extrasolar planets reside in eccentric orbits.  Thus it has remained inappropriate to compare the extrasolar planets against Jupiter or Saturn in our Solar System.  The Lick Observatory Doppler planet search began in 1987 and thus now has the requisite duration to detect planets having orbital periods of over a decade.  The velocities of 55 Cnc during 13 years can be explained nearly adequately by two planets orbiting the star.  We had previously detected the inner planet to 55 Cnc (``55 Cnc b'', Butler et al. 1997). Its mass ($>$0.9 \\mjup) and circular orbit with a radius of 0.11 AU from the star represents a class of close--in extrasolar planets, sometimes called ``hot jupiters'', the first member of which was 51 Peg (Mayor and Queloz 1995).  The velocities of 55 Cnc now reveal strong evidence of an outer planet at 5.9 AU, previously suspected due to the additional wobble of 55 Cnc (Marcy \\& Butler 1998).  The reality of an outer planet with an orbital period of 13--15 yr and a minimum mass of 4 \\mjup is securely supported by the velocities.  Remaining velocity residuals with RMS of 13 \\ms are caused either by gas motions on the stellar surface or by additional orbiting bodies.  The best three--planet fits imply a third planet having \\msini = 0.25 \\mjup at 0.24 AU in an orbit with $e$=0.3. But the three--planet models only partially explain the discrepancies in the two--planet fit.  As the first extrasolar planet discovered that orbits farther than 4 AU from its host star, the outer planet to 55 Cnc (``55 Cnc c'') is the first one amenable to direct comparison with the Jovian planets in our Solar System.  The outer planet has an orbital eccentricity of 0.16$\\pm$0.06 to be compared with 0.048 and 0.054 for Jupiter and Saturn in our Solar System.  Thus 55 Cnc c has a modest orbital eccentricity corresponding to an orbital path that carries it as close as 5 AU from the star and as far as 6.8 AU.  At a typical angular separation from the star of 0.47 arcsec, the planet 55 Cnc c will induce an astrometric wobble in the host star with an amplitude of 1.8 milliarcsec/$\\sin i$ relative to the barycenter.  This gives some hope that Hipparcos, the {\\it Hubble Space Telescope FGS}, or some other ground--based astrometric program could detect the wobble.  We carried out an analysis of the Hipparcos and Tycho--2 catalog astrometry similar to that described by Pourbaix (2001) and Pourbaix \\& Arenou (2001).  In some analyses, we used the long term proper motion from Tycho--2 to search for a residual astrometric wobble in the Hipparcos astrometry of 55 Cnc.  In other analyses we searched for a self--consistent solution of all available astrometry from both Hipparcos and Tycho-2.  We found no significant wobble in 55 Cnc at a level of $\\sim$3 milliarcsec over the time scale of the life time of Hipparcos, 4 years.  However this time baseline is too short to place any constraints at all on the inclination of the orbit of 55 Cnc c.  The wobble of the star caused by it would be nearly linear during 4 yr, and hence would be absorbed into the solution of the proper motion of the star.  Similarly, no HST FGS astrometry has adequate duration.  We are unaware of any ground--based astrometry that has adequate precision to detect the companion.  Thus we are not able to place any constraints on the orbital inclination of 55 Cnc c, and hence cannot place an upper limit on its mass.  Both SIM and GAIA would carry adequate astrometric precision to detect the motion of the star due to 55 Cnc c.  But a mission lifetime of at least 7--10 years (nearly one orbital period) will be necessary to separate the proper motion from orbital parameters.  Astrometry having a precision of $\\sim$20 $\\mu$as, coupled with velocities, would constrain the mass of 55 Cnc c to within a few percent.  The velocity residuals to the two--planet model exhibit an RMS of 12 \\ms and a strong periodicity of 44.3 d (Fig 9).  These residuals are certainly not due to any instrumental effect, as we are monitoring 300 stars at Lick, including many stars with spectral type G5--K0. No periodicities between 40--50 d are seen among those other stars.  The proposition that the 44.3 d period is caused solely by surface effects on the star seems unlikely.  The stellar characteristics of 55 Cnc (section 2) suggest that it is a quiescent star of age 3--8 Gyr, showing very little variation in the usual surface diagnostics.  The photometric variation is no more than 1 mmag, and the level and activity in the CaII K-line emission reversal is small.  Such stars are well studied by precision Doppler programs and they exhibit velocity variations of 2--5 \\ms, presumably caused by turbulence and patchy magnetic regions located non--uniformly over the surface.  Thus we cannot support a model in which the velocity residuals with RMS of 12 \\ms are intrinsic to the stellar surface.  In contrast, our attempts to fit the velocities, notably the 44.3 d period, with a third planet yielded a significant improvement in the reduced $\\chi^2$ compared with that of a simple two--Keplerian fit. Neither a triple--Keplerian model nor a triple--planet Newtonian model succeeded in diminishing the value of $\\chi_{\\nu}^2$ to a value of 1.0--1.5, but instead left $\\sqrt{\\chi^2_{\\nu}}$=1.8 .  Moreover the near coincidence in periods between the 44.3 d velocity period and the 35--42 d rotation period leaves us uncomfortable about the interpretation of the 44.3 d period.  Nonetheless, the chromospheric and photometric quiescence of the star is not consistent with stellar surface effects as the cause of the 13 \\ms velocity variations. Indeed, we have never seen such large velocity amplitude and coherence in such a quiet star.  This issue is examined carefully by Henry et al. (2002).  Thus we favor the model that includes a third planet with that period.  Because the value of $\\chi^2$ remains too large, even with a model that contains three planets, one may consider alternative models. Perhaps 55 Cnc contains yet an additional planet located in the gap between 0.25 and 5 AU.  The simulations presented here show that a low--mass planet could persist stably there indefinitely.  Planets of sub--saturn mass located between 0.25 and 5 AU would be difficult to detect securely but would cause velocity variations of a few \\ms, as seen in our velocity measurements.  Another possibility is that rotational modulation of surface inhomogeneities is stronger in metal--rich stars than is seen in solar--metallicity stars.  In that case, the 44.3 d period could be caused by stellar rotation after all.  Such a hypothesis requires that surface effects both cause a stellar ``jitter'' of 13 \\ms and remain coherent in phase over time scales of years.  This could occur if one longitude maintains its inhomogeneity (spot, magnetic field) for a duration of years.   We also note that the lack of a dust disk (Jayawardhana et al.~2002, Schneider 2001) provides limits to the evolution of debris disks in the face of a giant planet at Jovian distances.  One wonders whether such Jupiter analogs tend to enhance to production of dust via enhanced collision rates between the comets and asteroids or instead promote the clearing of the dust due to gravitational perturbations during the lifetime of the star.  The separation of 0.47 arcsec between the star and the outer planet, makes this system a likely target for future coronographic imaging and interferometric nulling, especially from spaceborn telescopes. The outer planet, ``55 Cnc d'', subtends a fraction of the sky, $f = 1.6\\times10^{-9}$, as seen from the star. The wavelength--dependent albedo of giant planets in general is under active investigation (Seager \\& Sasselov 1998; Marley et al. 1999; Goukenleuque et al. 2000; Sudarsky, Burrows, \\& Pinto 2000).  The albedo at visible wavelengths is likely to be $\\sim$1/2 .  Thus, one expects the planet, 55 Cnc d, to be fainter than the host star by a factor of $0.8\\times10^{-9}$ at optical wavelengths. This implies a contrast of 22.7 mag at V band, and an apparent magnitude, $V$ = 28.7, for the planet.  With its high abundance of Fe, C, Si, and other heavy elements, along with its age of $\\sim$5 Gyr,  the 55 Cnc system makes an intriguing target for questions of organic chemistry and biology. A rocky planet at roughly 1 AU remains a viable prospect dynamically.  Moreover the inner two planets and the outer planet presumably formed from protoplanetary disk material. These extent planets beg the question of the final repository of the disk material that presumably existed between 0.3 and 5 AU.  The migration of the inner two planets would not have cleared the region at 1 AU.  Indeed, such migration could have occurred by virtue of the two planets delivering angular momentum and energy to the disk material outward of their orbits.  Indeed, the presence of two planets at 0.1 and 0.24 AU suggests that material existed between 0.3 and 5 AU, serving as the recipient of their orbital angular momentum.  We expect to detect, within the next 5 years, a sizable population of Jupiter--mass planets orbiting at 4--6 AU.  These planets may serve as signposts of planetary systems characterized by architectures similar to that of our Solar System: gas giants beyond 5 AU and rocky planets closer in.  Such jupiters are amenable to direct comparison with Jovian planets in our Solar System and will permit characterization of the properties of planetary systems in general.  \\eject \\eject"
    },
    "0207/astro-ph0207308_arXiv.txt": {
        "abstract": "s{ After introducing several aspects of the motivation for particle dark matter search, experimental principles and the present state of the main experiments are presented. Direct searches for WIMPs are explained in some detail; indirect WIMP searches and axion searches are presented more briefly.} ",
        "introduction": "All viable models in present-day cosmology require the existence of non-baryonic cold dark matter with a cosmological density dominating over all other forms of matter in the universe.  The generally accepted inflationary cosmological models require the universe to be flat, i.~e. the total matter and energy density $\\Omega = 1$ in units of the critical density separating a universe with positive from one with negative curvature of space-time, to avoid the ``fine-tuning problem'': any deviation of the cosmological density from unity would have increased by many orders of magnitude during inflation, contradicting observations \\cite{Boerner93}. Thus $\\Omega$ should be precisely equal to one from the beginning. The flatness of the universe is confirmed, $\\Omega = 1$ to within a few per cent, by measurements of the first acoustic oscillations in the angular power spectrum of the cosmic microwave background (CMB) with the BOOMERANG, MAXIMA, DASI and other mostly satellite or balloon borne instruments \\cite{Masi02,Stompor01,Pryke02}.  The observation of the brightness-to-distance relationship of type Ia supernovae reveals evidence for the presence of some form of Dark Energy $\\Lambda$ with a density $\\Omega_\\Lambda \\approx 0.7$ \\cite{Perl99}. A similar value (0.65 - 0.85) has been obtained by a combined analysis of CMB data and the power spectrum of the matter density distribution \\cite{Efs01}. This leaves an overall matter density $\\Omega_m \\approx 0.3$, in good agreement with the value inferred from the dynamics of galaxies within clusters. On the other hand, the theory of big bang nucleosynthesis combined with measurements of the cosmic abundances of helium and, more recently, deuterium relative to that of hydrogen limit the cosmic density of baryonic matter quite strictly to $\\Omega_b h^2 = 0.020 \\pm 0.002$, where $h$ denotes the Hubble parameter \\cite{Burles01}. Also this value agrees well with that resulting from CMB measurements.  Thus the missing matter must be in form of non-baryonic particles which interact only weakly with normal matter. Most particles in the universe have frozen out of thermal equilibrium with the primordial plasma at a temperature corresponding to their respective mass. Dark matter candidates that were relativistic at the time of freeze-out are referred to as hot dark matter (HDM); those which were non-relativistic as cold dark matter (CDM). This dark matter ``temperature'' can be ``measured'' by means of the power spectrum of the matter density distribution as it results from all-sky galaxy surveys. HDM would have, due to its relativistic streaming while structure formation in the early universe, wiped out small-scale structures, i. e. prohibited the early formation of individual galaxies. CDM instead, freezing out earlier than baryons due to its weak interaction, could have fallen into density fluctuations very early thus enhancing the formation of small scale structures. The observed structures are best described by the matter content being dominated by CDM, with some HDM being allowed. A small HDM component can readily be explained by low-mass neutrinos since neutrinos are now, after the intriguing results from Superkamiokande and SNO \\cite{Fukuda99,Ahmad01}, generally considered having a small mass: the present neutrino mass limits correspond to a contribution to the cosmic density in the range $0.001 < \\Omega_\\nu < 0.18$ \\cite{Ahmad01}.  This article will concentrate on CDM candidates and approaches to their detection, and within this frame it will be further constrained to those candidates that are well motivated in that they emerge from well-established theories in particle physics and can at the same time constitute the required matter density without arbitrary assumptions of their physical parameters: supersymmetric Weakly Interacting Massive Particles (WIMPs) and axions. ",
        "conclusions": ""
    },
    "0207/astro-ph0207622_arXiv.txt": {
        "abstract": "We present an analysis of two deep (75 ks) Chandra observations of the European Large Area ISO Survey (ELAIS) fields N1 and N2 as the first results from the ELAIS deep X-ray survey. This survey is being conducted in well studied regions with extensive multi-wavelength coverage. Here we present the Chandra source catalogues along with an analysis of source counts, hardness ratios and optical classifications. A total of 233 X-ray point sources are detected in addition to 2 soft extended sources, which are found to be associated with galaxy clusters. An over-density of sources is found in N1 with 30\\% more sources than N2, which we attribute to large-scale structure. A similar variance is seen between other deep Chandra surveys. The source count statistics reveal an increasing fraction of hard sources at fainter fluxes. The number of galaxy-like counterparts also increases dramatically towards fainter fluxes, consistent with the emergence of a large population of obscured sources. ",
        "introduction": "The results of recent deep X-ray surveys reveal that almost the entire X-ray background can be resolved into discrete sources. The ROSAT Deep Survey (Hasinger {\\it et al.} 1998) resolved 70 - 80\\% of the 0.5 - 2 keV background at a flux level of 1 $\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$. Observations with Chandra and XMM-Newton are now pushing the detection limits even further. In particular, the unprecedented resolution of Chandra allows extremely deep observations that are not limited by source confusion. This has been exploited in the Chandra Deep Fields (North, Brandt {\\it et al.} 2001, and South, Giacconi {\\it et al.} 2002). In the Chandra Deep Field-North 2 Msec of data has been accumulated reaching a flux limit of $\\sim 1.5 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ in the 0.5 - 2 keV band (Barger {\\it et al.} 2003). However, the greatest advances have been at higher energies where Chandra is now beginning to resolve the 2 - 8 keV background.  The majority of sources resolved by ROSAT were found to have spectra that were too steep to account for the flat spectrum of the hard X-ray background. However towards fainter fluxes a new population emerged in the ROSAT data with intrinsically harder X-ray spectra (Hasinger {\\it et al.} 1993, Almaini {\\it et al.} 1996). Chandra is now uncovering a large number of hard spectrum sources, and the majority of the 2 - 8 keV background has been resolved. Over the flux range 2 $\\times 10^{-16}$ to 10$^{-13}$ erg s$^{-1}$ cm$^{-2}$ the contribution of resolved sources to the 2 - 8 keV background is 1.1 $\\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (Cowie {\\it et al.} 2002). This translates to $\\sim$ 65 - 85 per cent of the background as measured by Vecchi {\\it et al.} (1999, Beppo-Sax) and Ueda {\\it et al.} (1999, ASCA) respectively.  Early spectroscopic observations are finding a majority of the sources with hard X-ray spectra to be type II AGN, indicated by the presence of narrow lines (Tozzi {\\it et al.} 2001, Barger {\\it et al.} 2001a, Hornschemeier {\\it et al.} 2001). Most of these are found at $z < 1$. However, a considerable fraction of the hard X-ray sources are optically faint, probably due to obscuration, and provide challenging targets for spectroscopic identification. Sources identified as type I AGN display softer X-ray spectra and are observed to have a higher median redshift.  There are still a number of unanswered questions relating to the properties of the hard X-ray populations at longer wavelengths. AGN with large X-ray absorbing columns do not always appear as type II AGN in the optical (e.g. Maiolino {\\it et al.} 2001, Willott {\\it et al.} 2002). The relationship between gas and dust absorption in AGN remains unclear. It is also uncertain where the absorbed radiation may be re-radiated. Approximately $\\sim$ 7 per cent of X-ray sources in the Chandra Deep Field North are sub-millimetre sources (Barger {\\it et al.} 2001b), however whether this is the result of reprocessed nuclear emission or due to a starburst component, is unknown. Almaini {\\it et al.} (2003) find evidence for a strong angular cross-correlation between the X-ray and sub-millimetre populations. They suggests there may be an evolutionary sequence in these galaxies between the major episode of star-formation (sub-millimetre sources) and the onset of quasar activity (X-ray sources). To more fully understand the nature of these sources will require in-depth multi-wavelength studies of the X-ray source population.  We are conducting a deep X-ray survey with Chandra and XMM in two of the European Large Area ISO Survey (ELAIS) fields, N1 and N2. These high latitude fields were chosen for their low cirrus emission, and have a wealth of multi-wavelength data available. Both fields have been observed with ISO at 7, 15, 90, and 175 $\\mu$m (Oliver {\\it et al.} 2000), with the VLA at 1.4 GHz (Ciliegi {\\it et al.} 1999, Ivison {\\it et al.} 2002), and have deep g$'$, r$'$, i$'$, H, and K imaging (Gonzalez-Solares {\\it et al.} 2003).  Region N2 has been mapped with SCUBA to 8 mJy at 850$\\mu$m (Fox {\\it et al.} 2001, Scott {\\it et al.} 2001). As well as the Chandra observations described here, XMM-Newton observations in region N1 ($5\\times 30$ ksec pointings) are awaiting scheduling.  In this paper we present the analysis of the Chandra X-ray data and the Chandra source catalogue. Paper II (Gonzalez-Solares {\\it et al.} 2003) will present details of the optical identifications. ",
        "conclusions": "We have presented the Chandra source catalogues for deep (75 ks) observations of the ELAIS fields N1 and N2. A total of 233 X-ray point sources are detected: 225 in the 0.5 - 8 keV band, 182 in the 0.5 - 2 keV band, and 124 in the 2 - 8 keV band. In addition, 2 extended sources are detected in N2 in the 0.5 - 2 keV band and are found to be associated with galaxy clusters.  An over-density of sources is found in N1 with 30\\% more sources than N2. This difference is present in both the soft and hard band number counts and may be attributed to large-scale structure. A similar variance is seen between other deep Chandra surveys.  Source count statistics reveal an increasing fraction of hard sources at fainter fluxes. The number of galaxy-like counterparts also increases dramatically towards fainter fluxes, consistent with the emergence of a large population of obscured sources. Additionally, objects with galaxy-like and faint optical counterparts exhibit harder X-ray spectra towards fainter fluxes, consistent with significant absorbing columns in this population.  \\bigskip  \\noindent The source catalogues and further information regarding the ELAIS deep X-ray survey can be found at this URL: http://www.roe.ac.uk/$\\sim$jcm/edxs"
    },
    "0207/astro-ph0207414_arXiv.txt": {
        "abstract": "We study on the self-consistency of the pulsar polar cap model, i.e., the problem of whether the field-aligned electric field is screened by electron-positron pairs that are injected beyond the pair production front. We solve the one-dimensional Poisson equation along a magnetic field line, both analytically and numerically, for a given current density incorporating effects of returning positrons, and we obtain the conditions for the electric-field screening. The formula which we obtained gives the screening distance and the return flux for given primary current density, field geometry and pair creation rate at the pair production front.  If the geometrical screening  is not possible, for instance, on field lines with a super-Goldreich-Julian current, then the electric field at the pair production front is constrained to be fairly small in comparison with values expected typically by the conventional polar cap models. This is because (1) positive space charge by pair polarization is limited to a small value, and (2) returning of positrons leave pair electrons behind. A previous belief that pair creation with a pair density higher than the Goldreich-Julian density immediately screens out the electric field is unjustified at least for for a super Goldreich-Julian current density. We suggest some possibilities to resolve this difficulty. ",
        "introduction": "One of the promising models for the pulsar action is the polar cap model where the field-aligned electric field accelerates charged particles up to TeV energies, and resultant curvature radiation generates an electromagnetic shower. The model can provide an environment for radio emission mechanism with a particle beam in electron-positron pair plasmas and may also explain the gamma-ray pulsation. The pairs will be a particle source of the pulsar wind. The polar cap potential drop is a part of the electromotive force produced by the rotating magnetic neutron star. The voltage is a few percent of the available voltage for young pulsars, while it becomes some important fraction for older pulsars. The localized potential drop is maintained by a pair of anode and cathode regions, the formation mechanism of which is the long-standing issue of the polar cap accelerator.  Because any deviation of the space charge density from the Goldreich-Julian (GJ) density $\\rho_{\\rm GJ}$ causes the field-aligned electric field within the light cylinder, how the GJ charge density changes along a given magnetic field line plays an essential role in formation of the field-aligned electric field. Note here that the GJ density is determined by the magnetic field geometry: $\\rho_{\\rm GJ} \\approx - \\Omega B_z / 2 \\pi c$, where $\\Omega$ is the angular velocity of the star and $B_z$ is the field strength along the rotation axis. For instance, the outer gap, another promising accelerator model, is thought to appear around the null surface on which the GJ density changes its sign. Due to the field geometry, the outer gap originally has a pair of the anode and cathode regions without current and pair plasma. The space-charge-limited flow for the polar cap model on field lines curving toward the rotation axis also has a similar GJ density distribution providing a pair of the anode and cathode regions (actually it is regarded as a version of outer gap with external current [Hirotani \\& Shibata 2001]).  Once pairs are created in the accelerator, since the pair {\\em number density} far exceeds the GJ value $|\\rho_{\\rm GJ}/e|$, where $-e$ is the electronic charge, dynamics of the created pairs makes crucial effects in localizing the field-aligned electric field.  For the polar cap models, the space-charge-limited flow is found to produce the field-aligned electric field in some cases where the current density and magnetic field geometry are assumed (Fawley, Arons \\& Scharlemann 1977, Scharlemann, Arons \\& Fawley 1978). Later on, effects of pair creation and general relativity are taken into account (Arons \\& Scharlemann 1979, Muslimov \\& Tsygan 1992). In these models, given a current density, the space charge density deviates from the GJ charge density, and as a result the field-aligned electric field develops. Since the work function of charged particles on the stellar surface is sufficiently small, either electrons or ions can freely escape from the surface, and therefore the electric field on the surface is assumed to be zero. Based on this scheme, primary electrons are shown to be accelerated to energies high enough to induce an electromagnetic shower.  The current density on each magnetic lines of force will not be determined by the local electrodynamics in the polar cap, but will be determined in a global manner. Most of the rotation power of the pulsar is conveyed by the pulsar wind, which is accelerated near and beyond the light cylinder. Although the acceleration mechanism of the wind is still mystery, it is certain that an energetically dominant process which determines the current distribution takes place near and beyond the light cylinder. When the wind is accelerated (or heated), the current runs across the field lines ($\\vec{E} \\cdot \\vec{j} >0$) as a simple result of the Poynting theorem, provided that the ideal-MHD condition is more or less valid. Therefore, some of the current coming up from the star crosses the magnetic field lines corresponding to the wind acceleration and returns back to the star. The rest of the current reaches large distances (e.g., nebula shock), and the same amount with opposite direction comes back to the star. In any case, the current closure is seriously controlled by the dynamics near and beyond the light cylinder. The current determined in such a way runs through the polar cap. Thus, the current density distribution across the  polar-cap magnetic flux should be linked to the outer magnetosphere: the current density distribution is most likely to be determined regardless of the field geometry of the inner magnetosphere. We, therefore, treat the current density as an adjustable free parameter in our series of paper (Shibata 1991, 1995, 1997, Shibata, Miyazaki \\& Takahara 1998): the current density is a free parameter for the polar cap model, and it is adjusted when the local model is linked with a model of the outer magnetosphere.  For a space-charge-limited flow, it is shown by Shibata (1997) that both toward and away geometries are possible to generate a large field-aligned electric field causing an electron-positron pair cascade as long as the current density is reasonably large, comparable to the GJ current density $\\rho_{\\rm GJ}c$. In his paper, the cases of super-GJ current densities are examined as well as the sub-GJ cases. Note that the GJ current density is simply the GJ current, $I_{\\rm GJ} \\approx \\mu \\Omega^2 /c$ divided by the polar cap area, where $\\mu$ is the magnetic moment of the star. The GJ current is derived by an order-of-magnitude argument, and the net current circulating in the magnetosphere will not be far different from this value. However, if the current is not distributed more or less uniformly over the polar cap, but is focused or fragmented in the magnetic flux tubes by analogy of planetary aurori, then the local current densities can be much higher than the GJ current density. Furthermore, the current density will change in a wide rage field lines by field line.  These points motivate us to examine various combinations of current density and magnetic field geometry for the local accelerator. Although deviation of the space charge is found to cause the accelerating electric field, it is not clear how the electric field is screened to complete the localized potential drop. Only the known case in which a finite potential drop is formed without pairs is the geometrical screening, i.e., the case of toward geometry or general relativistic geometry (Muslimov \\& Tsygan 1992) with a certain current density which is sub-GJ. For other cases, there remains in general an unscreened electric field in a region where pairs are created. Note that pair plasmas are copious and pair polarization may become efficient. In this paper, we derive the condition for the electric-field screening immediately behind the pair creation front where unscreened electric field is assumed to exist, in the presence of returning pair. The condition can be used for any current density (super- or sub-GJ) and for both away and toward field line curvatures. We will apply the condition to the cases for which geometrical effect is not efficient, i.e., away curvature and super-GJ current density. Another important application is the case where pairs are created before the geometrical screening completes.   In the previous paper (Shibata, Miyazaki \\& Takahara 1998,  Paper I), we demonstrated that the pair polarization does work when enough pairs are injected, where we assumed that pairs are created at a single point; the required multiplication factor is $\\sim 10^3$ which is the same order of the value inferred from observations of the Crab pulsar and nebula. However, we suggested that the required multiplicity has to be reached in a very small distance, and this will not be realized in an actual pulsar magnetosphere.  We have made some simplifications in Paper I in order to make the physics as clear as possible. They are, 1) pairs have the same initial Lorentz factor and outward momentum, 2) the pair creation occurs at a single point, and 3) no positrons return to the neutron star. In this paper, we relax the above assumptions 2) and 3) to study more realistic cases: electron-positron pairs are produced continuously along field lines for a finite length, and some particles may stop before reaching the screening point and return to the neutron star surface. As in Paper I, we assume that no frictional force works between the various components. Under these assumptions, we calculate the electric field structure of the pair creation zone and determine the maximum electric field which can be screened under a presence of returning particles. Since the thickness of this zone is geometrically thin, we adopt one-dimensional approximation in this paper.  In section 2, we formulate the problem, and in section 3 we solve the one-dimensional Poisson equation analytically assuming that the pair creation rate is constant as a function of the electrostatic potential. In section 4, we solve the problem numerically when the pair creation rate is spatially uniform to confirm that the analytic results of section 3 are sufficiently accurate. In section 5 we summarize our results and suggest a possible way out to resolve the screening problem. ",
        "conclusions": "In this paper, we considered the electric field screening by polarization of electron-positron pairs which are created beyond the pair production front as is generally supposed for the polar cap models. We allow the pair positrons to return back to the star as far as the condition for the acceleration is satisfied. The one-dimensional Poisson equation is solved together with particle motion for a steady state acceleration region. We obtain the formula (\\ref{27}) which give the electric field strength that can be screened for a given current density $j$, a given field geometry $j_0$ and a given pair creation rate $\\alpha_\\phi$. It is found that (1) pair polarization has little contribution for positive space charge, and (2) returning of positrons makes the screening difficult seriously because the pair electrons left behind the returning positrons produce negative space charge in the screening region where the positive space charge is required. As a result, the thickness of the screening is restricted to be as small as the braking distance $\\Delta s \\approx mc^2 / e | E_{\\parallel} |$ for which positrons become non-relativistic. We confirmed the previous result of Paper I that the electric field in the case where geometrical screening is not possible, i.e., $-j+j_0+jM_1 <0$ at the pair production front, we have, from (\\ref{30}), \\begin{equation} {E_\\parallel^2 \\over 8 \\pi } < mc^2 \\left( \\Omega B \\over 2 \\pi c e \\right) \\zeta^\\prime j \\Delta M_{\\rm screen} \\end{equation} where $\\Delta M_{\\rm screen}$ is the pair multiplication factor {\\em within} $\\Delta s$. If the primary current density is of order of the Goldreich-Julian (GJ) value, the required pair multiplication factor per one primary electron is enormously large and cannot be realized in the conventional pair creation models. A previous belief that pair creation with a pair density higher than the GJ density immediately screens out the electric field is unjustified.  Some mechanism to salvage this difficulty should be found. As already mentioned, if the last term of (\\ref{27}), $-(j-j_0 - j M_1) \\phi_{\\rm s}$, becomes positive screening is possible as far as $j<j_{\\rm cr} \\approx j_0$. Since $j_0 = B_{\\rm z}/B <1$, $j$ must be less than unity. It must be noted that the positive space charge exists before the pair production front in this case, so that screening takes place before the pair production front and $E_{\\parallel}$ is reduced. Even a reversal of $E_\\parallel$ is possible for a small $j$ (Shibata 1995). In any case, if geometrical screening works, a finite potential drop along field lines can complete. But the current density is restricted in a certain range.   The condition (\\ref{30}) implies that if the current density is much higher than the GJ value, the required multiplication factor can be reduced. Although the total current circulating in the magnetosphere should be of order of so-called GJ current, $I_{\\rm GJ} =\\mu \\Omega^2 / c$, where $\\mu$ is the magnetic moment of the star, current densities can be in principle focused: the GJ current density is just a value if $I_{\\rm GJ}$ is more or less uniformly distributed over the whole polar cap area, and therefore if the current distributed, say, in an annular region with fine structures such as seen in aurora, the focused current density will be much higher than the GJ current density. Thus, if $j$ is much larger than $j_0$, and the required multiplication factor will be reduced. (If the current is super-GJ, toward curvature will not helpful for screening because $-(j-j_0 - j M_1)$ is negative.) In any case, physics in focused field-aligned current, e.g., effects of inhomogeneity across the magnetic field lines, two stream instability in a high current density with unscreened electric field, has yet been properly studied.  Since $\\bar{\\gamma}$ is an order of 100, $\\phi_s$ turns out to be an order of 100 using (\\ref{e7}) and (\\ref{cnd2}). From (\\ref{alph}) we may safely approximate as $\\alpha_\\phi = M_2 /100$. From (\\ref{32}), $E_{\\rm f} \\le \\sqrt{M_2}/5$, taking $\\zeta^\\prime =2$ and $j=1$. As an example, let us assume that total pair production occurs with a multiplication factor $10^4$ over the distance $x_{\\rm p}=10^4$. Thus, $M_2 = 10^4 (x_{\\rm s}/x_{\\rm p})=x_{\\rm s}$. Noting that $E_{\\rm f} \\approx \\phi_{\\rm s} /x_{\\rm s}=100/x_{\\rm s}$, we obtain $x_{\\rm s} > 60$ and $E_{\\rm f } < 1.5$. Since the unit length $c/\\omega_{\\rm p}$ is typically 1 cm, this electric field corresponds to a voltage $mc^2/e$ per cm. Assuming that primary electrons are accelerated up to $\\gamma> 10^6$ with this weak field $E_{\\rm f} \\sim 1$, the required length between the neutron star surface and the pair production front is about $10^6$ cm in hight. Thus pair screening is possible if the pair production front is located high above the surface and if the field-aligned electric field is not so strong. However, the previous models predict a stronger electric field, so we would need an additional ingredient to keep a small field as long as a stellar radius. In any case, a weak electric field kept in a long distance is one possible solution to have a self-consistent polar cap model.    Another possible way out is to include frictional forces between various components of charged currents. Since friction on positrons pulls them outwards along with electrons, returning fraction of positrons will be reduced, which will make screening easier. Although there are some ideas for physical processes of friction (two stream instability, production of positroniums and others), it is not clear whether the frictional force can be strong enough to lift positrons and screen the electric field. Study on the frictional interaction under {\\it unscreened} electric field is strongly demanded.  \\subsection*"
    },
    "0207/astro-ph0207078_arXiv.txt": {
        "abstract": "The scientific goal of the SACY (Search for Associations Containing Young-stars) project is to identify eventual associations of stars younger than the Local Association, spread among the optical counterparts of the ROSAT X-ray bright sources. High-resolution spectra for possible optical counterpart later than G0 belonging to HIPPARCOS and/or TYCHO-2 catalogues were obtained in order to assess both the youth and the spatial motion of each target. The newly identified young stars present a patchy distribution in UVW space and in the sky as well revealing the existence of huge nearby young associations. Here we present the associations identified in the present sample. ",
        "introduction": "The detection of X-ray sources  by the ROSAT All-Sky Survey (RASS) associated with TTS outside star formation regions (Neuh\\\"{a}user 1997) gave a tool to find new young associations. In fact, Torres et al. (2000), and Zuckerman \\& Webb (2000), using these sources, found evidences for two young associations near the South Celestial Pole, the Horologium (HorA) and the Tucana (TucA) Associations. To examine the possibility that these associations are the same and to search for other ones we undertook a Search for Associations Containing Young-stars (SACY) (de la Reza et al. 2001; Torres et al. 2001; Quast et al. 2001). In the SACY we selected and observed all bright RASS sources that could be associated with TYCHO-2 or HIPPARCOS stars with (B-V) $>$ 0.6, excluding very well known RS CVn, W UMa, giants, etc in  SIMBAD.  We obtained high resolution spectra for the selected candidates with FEROS echelle spectrograph (Kaufer et al. 1999) (resolution of 50000; spectral coverage of 5000\\,\\AA) of the 1.52 m ESO telescope at La Silla or  with the coud\\'{e} spectrograph (resolution of 9000; spectral coverage of 450\\,\\AA, centered at 6500\\,\\AA) of  the 1.60\\,m telescope of the Observat\\'{o}rio do Pico dos Dias. For some stars we obtained radial velocities with  CORALIE at the Swiss Euler Telescope at ESO (Queloz et al. 2000).  >From  the collected spectra we have obtained spectral classifications, radial velocities and equivalent widths of Li\\,I lines. In particular, the Li\\,I line is important since it can provide a first age estimate (Jeffries 1995) for late type stars allowing us to select possible Post-T Tauri stars. For a given star, if its Li\\,I resonance line equivalent width is located near or above the Li\\,I line delimited by the members of the Local Associations clusters (Neuh\\\"{a}user 1997), it is  flagged as young star. ",
        "conclusions": ""
    },
    "0207/astro-ph0207552_arXiv.txt": {
        "abstract": "We explain why it is possible to detect directly X-ray emission from near the surface of the neutron star (NS) in SAX J1808.4-3658 but not in most other low-mass X-ray binaries (LMXBs), with the exception that emission from the surface can  be  seen during bursts events. We show that the X-ray emission from SAX J1808.4-3658 mostly originates in the Comptonization process in a relatively optical thin hot region (with an optical depth $\\tau_0$ around 4 and temperature is around 20 keV). Such a transparent region does not prevent us from detecting coherent X-ray pulsation due to hot spots on the NS surface. We give a precise model for the loss of modulation, such suppression of the QPO amplitude due to scattering can explain the disappearance of kHz QPOs with increasing QPO frequency. We also formulate general conditions under which the millisecond X-ray pulsation can be detected in LMXBs. We demonstrate that the observed soft phase lag of the pulsed emission is a result of the downscattering of the hard X-ray photons in the relatively cold material near the NS surface. In the framework of this downscattering model we propose a method to determine the atmosphere density in that  region from soft-lag measurements. ",
        "introduction": "Low-mass X-ray binaries (LMXB) presumably contain a  weakly magnetized neutron star (NS). Recently Titarchuk, Bradshaw \\& Wood (2001), hereafter TBW suggested a new method of estimating B-field strength  using the magnetoacoustic oscillation model and  found that the B-field strength for a number of NSs in  LMXBs is  $\\sim 10^8$ gauss. Near-coherent millisecond X-ray pulsations have been observed in 4U 1728-34 during thermonuclear (type I) X-ray bursts (e.g. Strohmayer et al. 1996). They are interpreted as X-ray intensity  modulated at period close to the spin period of the neutron star. Thus these B-field estimates in addition to the detection of the milisecond pulsation give strong arguments for LMXB neutron stars to be   progenitors of millisecond radio pulsars (MLP) (see review by Bhattacharya \\& van den Heuvel 1991). But there is still a question why these coherent pulsations are not found in persistent emission despite careful searches (Wood et al. 1991; Vaughan et al. 1994). The lack of coherent pulsations has been explained as modulation loss from gravitational lensing (Wood, Ftaclas \\& Kearney 1988; Meszaros, Riffert \\& Berthiaume 1988) or from to scattering (e.g. Bainerd \\& Lamb 1987; Kylafis \\& Klimmis 1987 ). The third explanation  for  the lack  of pulsations is presented by Cumming, Zweibel \\& Bildsten (2001) who argue that the surface field is weak because of magnetic screening  \\footnote{% The recent discoveries of MLP in XTE J1751-305 (Markwardt et al. 2002) and in XTE J0929-314 (Galloway et al. 2002 ) with extremely low mass transfer rates support the suggestion of the absence  of the magnetic screening  for the  low mass accretion rates.}. In this {\\it Letter} we put forth arguments for smearing out of the pulsar signal due to electron scattering in the optically thick environment typical of  most  observed LMXBs. In  \\S 2 we study the scattering effect and its relation with the observed timing and spectral characteristics  for various QPO sources. In \\S 3  we investigate the scattering effects in QPO sources. In \\S 4 we present results of the RXTE data analysis of spectral  properties of X-ray radiation in SAX 1808.4-3656. In \\S 5 we analyze the downscattering model and its application to the observed soft lag phenomenon detected in the coherent pulse signal from SAX 1808.4-3656. Conclusions follow in \\S 6. ",
        "conclusions": " (1)  the detection of the NS pulsations from SAX J1808.4-3658 was possible because of the relatively transparent Compton cloud covering in this source ($\\tau_0\\sim 4$). For the majority of the analyzed LMXBs the Compton cloud optical depth is at least twice as high as that in SAX J1808.4-3658. The high frequency pulsations with frequencies $\\gax 300$ Hz are strongly wiped out by scattering in the clouds with $\\tau_0>4$.  SAX J1808.4-3658 is a limiting case for this detection. There is a possibility of finding the NS pulsations in the sources that are much less luminous (than these  bright QPO LMXBs), because  in them the Compton cloud is more  transparent ($\\tau_0\\lax 4$) for the NS pulsed radiation. (2) The kHz QPOs  rms vs frequency  follows the resonance law $1/\\omega$ are weakened  by  scattering in the Compton cloud. (3) The soft-lag measurements along with the implementation of the downscattering model  can provide a tool for  density determination near  the NS surface.  We appreciate the fruitful discussions with   Paul Ray, Sergey Kuznetsov and particularly, with the referee."
    },
    "0207/astro-ph0207397_arXiv.txt": {
        "abstract": "{Rest-frame far-ultraviolet (FUV) luminosities form the `backbone' of our understanding of star formation at all cosmic epochs. FUV luminosities are typically corrected for dust by assuming that extinction indicators which have been calibrated for local starbursting galaxies apply to all star-forming galaxies. I present evidence that `normal' star-forming galaxies have systematically redder UV/optical colors than starbursting galaxies at a given FUV extinction.  This is attributed to differences in star/dust geometry, coupled with a small contribution from older stellar populations.  Folding in data for starbursts and ultra-luminous infrared galaxies, I conclude that SF rates from rest-frame UV and optical data alone are subject to large (factors of at least a few) systematic uncertainties because of dust, which cannot be reliably corrected for using only UV/optical diagnostics. }   \\addkeyword{Dust, Extinction} \\addkeyword{Galaxies: General} \\addkeyword{Galaxies: Stellar Content} \\addkeyword{Ultraviolet: Galaxies}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Understanding the star formation (SF) rates of galaxies, at a variety of cosmic epochs, is a topic of intense current interest (e.g., Yan et~al.\\@ 1999; Blain et~al.\\@ 1999; Haarsma et~al\\@ 2000). Many SF rates are derived from highly dust-sensitive rest frame far-ultraviolet (FUV) luminosities (e.g., Madau et~al.\\@ 1996; Steidel et~al\\@ 1999). In the local Universe, Calzetti et~al.\\@ (1994,1995) and Meurer et~al.\\@ (1999) found a tight correlation between ultraviolet (UV) spectral slope $\\beta$\\footnote{Defined by $F_{\\lambda} \\propto \\lambda^{\\beta}$, where $F_{\\lambda}$ is the flux per unit wavelength $\\lambda$.} and the attenuation\\footnote{Attenuation differs from extinction in that attenuation describes the amount of light lost because of dust at a given wavelength in systems with complex star/dust geometries where many classic methods for determining extinction, such as color excesses, may not apply.} in the FUV (\\afuvns), for a sample of in\\-hom\\-ogen\\-eous\\-ly-selected starburst galaxies. This correlation's low scatter requires a constant intrinsic value of $\\beta \\sim -2.5$ for young stellar populations (e.g., Leitherer et~al.\\@ 1999), coupled with some regularities in the distribution and extinction properties of dust (e.g., Gordon et~al.\\@ 1997). Assuming that this correlation holds for all galaxies at high redshift, this was used to correct the FUV flux for extinction in a statistical sense (see, e.g., Adelberger \\& Steidel 2000 and references therein).  However, recent work has called the universality of the $\\beta$--\\afuv correlation into question. Radiative transfer models predict a large scatter between $\\beta$ and \\afuv (Witt \\& Gordon 2000). Furthermore, both Large Magellanic Cloud (LMC) \\hii regions (Bell et~al.\\@ 2002) and ultra-luminous infrared galaxies (ULIRGs; Goldader et~al.\\@ 2002) do not obey the starburst correlation.  Tantalizingly, there are indications that `normal', quiescent star-forming galaxies have less UV extinction than predicted by the Calzetti et al.\\ relation (Buat et~al.\\@ 2002).  Taken together, these issues raise serious questions about the applicability of rest-frame UV-derived SF rates for non-starbursting galaxies.  Here, I investigate the relationship between $\\beta$ and \\afuv for quiescent, `normal' star-forming galaxies for the first time (to date, this correlation has been examined directly for starbursts, ULIRGs and \\hii regions only). For more details, see Bell (2002). ",
        "conclusions": "\\label{sec:conc}  Seven independent UV experiments demonstrate that quiescent, `normal' star-forming galaxies have substantially redder UV spectral slopes $\\beta$ at a given \\afuv than starbursting galaxies. Using spatially resolved data for the LMC, I argue that dust geometry and properties, coupled with a small contribution from older stellar populations, cause deviations from the starburst galaxy $\\beta$--\\afuv correlation. Neither rest frame UV-optical colors nor UV/\\ha significantly help to constrain the UV attenuation. Thus, SF rates estimated from rest-frame UV and optical data alone are subject to large (factors of at least a few) systematic uncertainties because of dust, which cannot be reliably corrected for using only UV/optical diagnostics.  However, SF rates for high $z$ galaxies derived from other wavelengths are also often subject to systematic errors of this magnitude.  For example, sub-mm fluxes for high-redshift star forming galaxies sample rest-frame $\\sim 200${\\micron}: they must be converted to total IR flux assuming some dust spectrum. Adelberger \\& Steidel (2000) derive values of $\\nu I_{\\nu}$ at 200{\\micron} vs.\\ $L_{\\rm FIR}$ of about 0.06 and 0.13 for ULIRGs and starbursts respectively.  Using Tuffs et~al.'s (2002) {\\it ISO} photometry of star-forming galaxies at 170{\\micron} as a constraint, I find that $\\nu I_{\\nu}$ at $\\sim$200{\\micron} vs.\\ $L_{\\rm FIR}$ is roughly 0.2 for galaxies with warm dust ($100/60{\\micron} \\sim 1$), growing to $\\sim 0.6$ for galaxies with cold dust ($100/60{\\micron} \\sim 6$). Thus, there is a factor of $\\sim 10$ systematic error because of dust temperature which affects submm-derived SF rates (see also Dunne \\& Eales 2001). A similar limitation\\adjustfinalcols affects radio-derived SF rates: because of the mismatch in cosmic ray propagation and SF timescales, an order of magnitude scatter between radio flux and SF rate is easily possible (Bressan et~al.\\@ 2002). Thus, when it comes to deriving SF rates for high-redshift galaxies from data at almost any wavelength, we are playing, at best, an order-of-magnitude game.  This work was supported by NASA grant NAG5-8426 and NSF grant AST-9900789."
    },
    "0207/astro-ph0207168_arXiv.txt": {
        "abstract": "The prevailing evidence suggests that most large-amplitude AGB variables follow the period luminosity (PL) relation that has been established for Miras in the LMC and galactic globular clusters. Hipparcos observations indicate that most Miras in the solar neighbourhood are consistent with such a relation. There are two groups of stars with luminosities that are apparently greater than the PL relation would predict: (1) in the LMC and SMC there are large amplitude variables, with long periods, $P> 420$ days, which are probably undergoing hot bottom burning, but which are very clearly more luminous than the PL relation (these are visually bright and are likely to be among the first stars discovered in more distant intermediate age populations); (2) in the solar neighbourhood there are short period, $P<235$ days, red stars which are probably more luminous than the PL relation. Similar short-period red stars, with high luminosities, have not been identified in the Magellanic Clouds. ",
        "introduction": "The Miras in globular clusters have always been key to calibrating the luminosity of the tip of the AGB. Unfortunately, because of their short lifetimes there are rather few Miras in globular clusters and fewer still that have been well studied. Let me remind you that the Miras are the most luminous stars found in the clusters; in fact they are the only stars with luminosities above the tip of the red giant branch (RGB). They are only found in metal-rich clusters ($\\rm [Fe/H]>-1$), and we presume that the AGB in metal-deficient systems terminates below the tip of the RGB. The pulsation period of a Mira is a function of the metallicity of its parent cluster (e.g.\\ Feast \\& Whitelock 2000b). In fact it is only for the Miras in clusters, and a very few in binary systems, that we can determine metallicities. The Miras in galactic globular clusters are all O-rich and there is no particular evidence to suggest that they have reached the thermally pulsing part of the AGB.  Feast et al.\\ (2002) recently reexamined the luminosities of the Miras in globular clusters using a new distance calibration based on Hipparcos parallaxes of sub-dwarfs and published photometry for 6 galactic globular clusters, together with new observations of NGC\\,121~v1, a short period low metallicity Mira in the SMC.  They demonstrated that these cluster Miras fit the same PL relation as do the LMC Miras, and derived a zero-point for the PL relation. There are many globular clusters, particularly near the galactic centre, which have not yet been properly surveyed for Miras. We are in the process of rectifying this situation using the Infrared Survey Facility in South Africa in collaboration with astronomers from the University of Tokyo. Any new cluster Miras will obvious improve our statistics, but we are particularly hopeful about finding some longer period stars in the metal-rich bulge clusters.  Once there are theoretical models which deal effectively with mass-loss, and allow us to predict accurately the AGB tip luminosity for different populations, we will be able to use them to calibrate extragalactic systems. But, until that level of theoretical understanding is reached, those who wish to study extragalactic systems will make deductions based on a comparison with globular clusters or with the galactic bulge. The recent literature contains numerous studies of AGB populations in local group galaxies and beyond and it is interesting to see comparisons being made with galactic globular clusters and very different conclusions being drawn by different authors from essentially the same data. Figure 1 shows a colour magnitude diagram for cluster variables, of the kind typically used for comparison with extragalactic systems. The most luminous cluster star illustrated here is NGC\\,6553~v4, for which the reddening is uncertain; it is plotted in the figure as two connected open circles for two different reddenings.  The lower mean luminosity seems more likely and this is one magnitude fainter than the $K=-8.5$ which some authors use.  \\begin{figure} \\hspace*{2cm}\\includegraphics[height=6cm]{paw_fig1.eps} \\vspace*{-5.5cm} \\narrowcaption{ A colour-magnitude diagram for the Miras in globular clusters; open circles represent less certain luminosities. A luminosity of $K= -8.5$ is often assumed for the brightest stars in globular clusters, but $K=-7.5$ is actually a better estimate (see Feast et al. 2002).} \\vspace*{0.5cm} \\end{figure}  It is also worth noting that comparing the luminosity of individual stars in extragalactic systems with those of AGB variables in the galactic bulge is even more fraught with uncertainty, because the shape of the bulge and the presence of significant numbers of foreground stars result in many stars having distances less than that of the centre and therefore luminosities that appear to be {\\it much} brighter than they really are.  Variability is also a factor in comparing one system with another. The short period Miras, found in globular clusters, typically have peak-to-peak $K$ amplitudes of around half a magnitude, so there is a high level of uncertainty associated with single measurements of the luminosity. Longer period stars have larger amplitudes, reaching over two magnitudes for the 1000 day variables discussed below. ",
        "conclusions": ""
    },
    "0207/astro-ph0207442_arXiv.txt": {
        "abstract": "{ We present an analysis of new Johnson and Str\\\"omgren photometric and medium-resolution spectroscopic observations of the $\\delta$ Scuti type variable star V784~Cassiopeae. The data were obtained in three consecutive years between 1999 and 2001. The period analysis of the light curve resulted in the detection of four frequencies ranging from 9.15 d$^{-1}$ to 15.90 d$^{-1}$, while there is a suggestion for more, unresolved frequency components, too. The mean Str\\\"omgren indices and Hipparcos parallax were combined to calculate the following physical parameters: $\\langle T_{\\rm eff} \\rangle$=7100$\\pm$100 K, ${\\rm log}~g$=3.8$\\pm$0.1, $M_{\\rm bol}$=1\\fm50$\\pm$0\\fm15. The position of the star in the HR diagram was used to derive evolutionary mass and age yielding to a consistent picture of an evolved $\\delta$ Scuti star with a mixture of radial plus non-radial modes. ",
        "introduction": "$\\delta$ Scuti-type variable stars are pulsating variables located in the lower part of the classical instability strip near or slightly above the main-sequence. The characteristic time-scale of the light variation is in the order of 0\\fd1, while  the observed light curves are usually multiperiodic due to the simultaneously excited radial and/or non-radial modes. Reliable mode identification requires long-term and/or multi-site observing campaigns, which are well-illustrated, e.g. by recent results from the Delta Scuti Network (Breger et al. 1998, Breger 2000), STEPHI program (Alvarez et al. 1998) or the Whole Earth Telescope (Handler et al. 1997). A handbook of reviews and discussion of the astrophysical importance of these variables has been published very recently (Breger \\& Montgomery 2000). The most complete catalogue of $\\delta$ Scuti stars has been tailored and analysed by Rodr\\'\\i guez et al. (2000) and Rodr\\'\\i guez \\& Breger (2001).  The light variation of the short period variable V784~Cas (=HD~13122=BD+59$^\\circ$422, P=0.1092 d, $\\Delta$V=0\\fm06, spectral type F5II, ESA 1997) was discovered by the Hipparcos astrometric satellite. The R00 catalogue (Rodr\\'\\i guez et al. 2000) includes this star listing the parameters derived from the Hipparcos observations. The star lies about 1$^\\circ$ NW of the open cluster Stock~2, but it is not associated with this strongly reddened cluster located at 316 pc (Krzemi\\'nski \\& Serkowski 1967). The Hipparcos parallax (9.81$\\pm$0.75 mas) supports the close proximity of the star (102$^{+8}_{-7}$ pc). A few radial velocity measurements can be found in the literature, they range from $-$6 km~s$^{-1}$ (De Medeiros \\& Mayor 1999) to $+20$ km~s$^{-1}$ (Duflot et al. 1995). $UBVRI$ photometry was given by Fernie (1983), while the star was included in the list of bright northern stars with interesting Str\\\"omgren indices by Olsen (1980). V784~Cas was also studied in a sample of bright giant stars by L\\`ebre \\& De Medeiros (1997), where no emission features, neither time variations or asymmetries of the H$\\alpha$ line profiles have been detected (this star was observed two times separated by ten months). The measured rotational velocity is 66 km~s$^{-1}$ (De Medeiros \\& Mayor 1999). Neither of the studies mentioned above dealt with the time-dependent phenomena, only the scatter of the velocity measurements (4 km~s$^{-1}$), as listed in De Medeiros \\& Mayor (1999), suggested the possible variability. Most recently, Gray et al. (2001) included the star in their large sample of late A-, F- and early G-type stars and determined its spectral type (F0-F2III). They also noted that V784~Cas is a mild A$m$ star, the lines of Sr II $\\lambda$4077 and $\\lambda$4216 are enhanced. There is no metallicity determination in the literature.  We started a long-term observational project of obtaining follow-up observations of bright, new variable stars discovered by the Hipparcos satellite. We have so far identified a candidate second overtone field RR~Lyrae variable (Kiss et al. 1999a), a new high-amplitude $\\delta$ Scuti star (Kiss et al. 1999b) and revealed the misclassification of a contact binary (Cs\\'ak et al. 2000). The main aim of this paper is to present an analysis of new photometric and spectroscopic observations of V784~Cas. The paper is organised as follows: the observations are described in Sect.\\ 2, Sect.\\ 3 deals with the period analysis, while radial velocities are discussed in Sect.\\ 4. Finally, the physical parameters are presented in Sect.\\ 5. ",
        "conclusions": "In this paper, we presented an analysis of photometric and spectroscopic observations of the recently discovered $\\delta$ Scuti variable V784~Cas. The $UBV$ photometry was carried out at Szeged Observatory (Hungary), while the simultaneous $uvby$ data were obtained at Sierra Nevada Observatory (Spain) in three consecutive years (1999-2001). Medium-resolution spectroscopy in the H$\\alpha$ region was carried out at the David Dunlap Observatory (Canada) in 1999 and 2001. These data were used to determine the frequency content and to estimate physical parameters of the star. The main results can be summarized as follows:  \\noindent 1. Multicolor data consisting of more than 3000 individual points were analyzed with the standard Fourier-analysis. The multiperiodic nature of the star is revealed unambiguously. Besides the dominant period listed also in the Hipparcos catalog, we could detect three more frequencies in the 9.46-15.9 d$^{-1}$ range. There is a suggestion for more, unresolved frequency components.  \\noindent 2. We have obtained almost 100 radial velocity measurements using the H$\\alpha$ line. The measured radial velocity curves also show the multiperiodic nature and a close correlation with the four-component light curve fit. Spectra obtained in 1999 covered a few weak metallic lines and the varying asymmetric line profiles suggest the presence of non-radial pulsation, too.  \\noindent 3. Physical parameters of the star are determined from the mean Str\\\"omgren indices and synthetic colour grids. The resulting parameters give a consistent picture of an evolved $\\delta$ Scuti star. Evolutionary mass ($1.89\\pm0.11 M_\\odot$) and age ($1.3\\pm0.3$ Gyr) is derived.  \\noindent 4. Possible mode identification was discussed based on the Str\\\"omgren photometric behaviour (amplitude and phase relations). We identify $f_1$ with the radial fundamental mode, while the remaining frequencies correspond to low-order ($l$=1 or 2) non-radial modes, although some ambiguity may arise from the moderate rotation of the star.  Further observations (photometric, as well as spectroscopic) of this variable star are expected to extend the data baseline yielding to a better resolution of the pulsational pattern, mode identification and detection of time-dependent phenomena (e.g. amplitude and/or frequency modulation)."
    },
    "0207/astro-ph0207218_arXiv.txt": {
        "abstract": "The well known Sunyaev-Zel'dovich (SZ) effect is reexamined using a Doppler shift type mechanism arising from the scattering of photons by electrons in an optically thin gas. The results are in excellent agreement with the observational data as well as with the results obtained using the diffusive pictures. A comparison of the results here obtained with other approaches is thoroughly discussed, as well as some important extensions of this method to other aspects of the SZ effect. ",
        "introduction": "}  As it is well known today, the cosmic microwave background radiation (CMBR) that fills the universe exhibits an almost perfect blackbody Planckian spectrum corresponding to a temperature of 2.726 K. A little over thirty years ago, it was predicted by the astrophysicists R.A Sunyaev and Ya. B. Zel'dovich that this spectrum could be distorted when photons from this radiation would penetrate large structures such as the hot intracluster gas now known to exist in the universe~\\cite{SZ1}. This distortion would arise from the interaction between the photons and electrons which constitute the plasma responsible for an important main contribution to the total mass contained in the cluster. Such distortion is only a very small effect changing the brightness of the spectrum by a figure of the order of 0.1 percent. This effect is now called the Sunyaev-Zel'dovich effect and its detection is at present a relatively feasible task due to the modern observational techniques available. Its main interest lies on the fact that it provides information to determine important cosmological parameters such as Hubble's constant and the baryonic density. Since all these and other important facts are readily available in the somewhat broad literature on the subject~\\cite{SZ1}-\\cite{Steen1}, we will not pursue any further details of its cosmological implications.  In this work we are mainly concerned with the nature of the various interpretations that have been provided to explain the effect and to propose another one which we believe, is much simpler to grasp than the others, and offers a much more direct way to account for the relativistic corrections, as well as the influence of $z$, the redshift factor, in the equations for the distorted spectrum.  The main idea behind the method we want to discuss is that photons incident on the hot electron gas of the plasma change their frequency by an absortion-emission process, so that the intensity of the spectral line is changed through the Doppler effect.  The overall effect is due to those photons that happen to be captured by an electron, an event which is in general unlikely to occur specially in an optically thin gas. Thus, an electron moving with a given thermal velocity emits (scatters) a photon with a certain incoming frequency $\\nu _o$ and outgoing frequency $\\nu $. The line breath of this process is readily calculated from elementary kinetic theory taking into account that the media in which the process takes place has an optical depth directly related to the usual Compton parameter $y$. When the resulting expression is convoluted with the incoming flux of photons obtained from Planck's distribution, one easily obtains the disturbed spectra. The results derived from the final equation are shown to be in agreement with those obtained from the observational data. To present these results we have divided the paper as follows. In section II a brief review of the previous methods developed to explain the SZ effect is given. This will allow a thorough comparison with our proposal. In Section III we develop our ideas and give the comparison between the mathematical results with both, those derived before and the observational ones. Section IV is left for some concluding remarks and future directions of this work. ",
        "conclusions": "Eq. (\\ref{trece}), the main result of this work hardly needs a more detailed explanation, it is a direct result of elementary statistical mechanics, except of course for the factor $\\tau $ which appears in the broadening of the spectral line. That this broadening must somehow depend on $\\tau $ is clear. As argued by many authors in the case of clusters of galaxies, the plasma is optically thin, so most of the photons are not scattered, and how many of them are must depend on the optical depth. Now, why in particular Eq. (\\ref{catorce}) is valid remains to be rigorously shown. In our case it came as a mere accident since in ref. \\cite{yo} the mass taken  in $w$ to perform the calculations was that of a proton, which turns out to be of the same order of magnitude as $\\frac{m_e}\\tau $ for $y=10^{-5}$. There is in fact one way of understanding this puzzle. In a pure diffusive (Fickian) process described by a Gaussian function one knows that the squared value of a single line  width grows as $t$. Looking at the original equation of Kompaneets and thinking only in its diffusive terms, one sees that the variable time is replaced by the Compton parameter \"$y$\", so that one should expect that, to a first approximation, the width of the curve grows as $\\sqrt{y}$. Hence, inside the gas, the effective frequency at which the photon propagates through the gas is, according to Eq. (\\ref{catorce}), proportional to $\\sqrt{y} \\times \\nu $. This argument, which is based on ideas pertinent to the diffusion mechanism must be extended to the Doppler picture, but this has not yet been accomplished.  The two main advantages of Eq. (\\ref{trece}) are, firstly, that its generalization to the relativistic case is straightforward and indeed, since in relativity theory  one has both the transverse and parallel Doppler effects~\\cite{Moller}, one can easily  study what their influence is, if any, in the full distorted spectrum. For the parallel effect the calculations carried out so far agrees reasonably well with those reported by Rephaeli  and other authors ~\\cite{rel} and will be the subject of a forthcoming publication. The case of the transverse effect is under study. Secondly, in all the work we have reported here, the cluster of galaxies, the source of the scattering processes of the CMBR and the electrons, has been considered a system at rest. One can now extend the calculations assuming that the cluster recedes to infinity according to Hubble's law and introduce the redshift factor into the formalism. The resulting prediction could be compared with the observational data and extract valuable information on such an important cosmological parameter. Work in this direction is also in progress. Concluding, we have the strong conviction that the Doppler effect approach to study the thermal SZ effect is not only very clear from the physical point of view, it also contains ingredients which are promising for further studies more advantageous than those offered by the diffusive model."
    },
    "0207/astro-ph0207504_arXiv.txt": {
        "abstract": "We describe the construction and study of an objectively-defined sample of early-type galaxies in low-density environments. The sample galaxies are selected from a recently-completed redshift survey using uniform and readily-quantified isolation criteria, and are drawn from a sky area of $\\sim$700 deg$^2$, to a depth of 7\\,000~\\kms\\/ and an apparent magnitude limit of $b_J \\le 16.1$. Their early-type (E/S0) morphologies are confirmed by subsequent CCD imaging. Five out of the nine sample galaxies show signs of morphological peculiarity such as tidal debris or blue circumnuclear rings. We confirm that E/S0 galaxies are rare in low-density regions, accounting for only $\\approx$8\\% of the total galaxy population in such environments.  We present spectroscopic observations of nine galaxies in the sample, which are used, in conjunction with updated stellar population models, to investigate star-formation histories.  Our line-strength analysis is conducted at the relatively high spectral resolution of 4.1~\\AA.  Environmental effects on early-type galaxy evolution are investigated by comparison with a sample of Fornax cluster E/S0s (identically analysed). Results from both samples are compared with predictions from semi-analytic galaxy formation models. From the strength of \\oii\\ emission we infer only a low level of ongoing star formation ($<0.15$~M$_\\odot$\\,yr$^{-1}$). Relative to the Fornax sample, a larger fraction of the galaxies exhibit \\oiiib\\/ nebular emission and, where present, these lines are slightly stronger than typical for cluster E/S0s.  The Mg--$\\sigma$ relation of E/S0s in low-density regions is shown to be indistinguishable from that of the Fornax sample.  Luminosity-weighted stellar ages and metallicities are determined by considering various combinations of line-indices; in particular the \\HgF\\/ {\\em vs}\\/ Fe5015 diagram cleanly resolves the age--metallicity degeneracy at the spectral resolution of our analysis. At a given luminosity, the E/S0 galaxies in low-density regions are younger than the E/S0s in clusters (by $\\sim$2--3~Gyr), and also more metal-rich (by $\\approx$0.2~dex). We infer that an anti-correlation of age and metallicity effects is responsible for maintaining the zero-point of the Mg--$\\sigma$ relation. The youngest galaxies in our sample show clear morphological signs of interaction. The lower mean age of our sample, relative to cluster samples, confirms, at least qualitatively, a robust prediction of hierarchical galaxy formation models. By contrast, the enhanced metallicity in the field is contrary to the predictions and highlights shortcomings in the detailed treatment of star-formation processes in current models. The [Mg/Fe] abundance ratio appears to span a similar, mostly super-solar, range both in low-density regions and in Fornax cluster galaxies. This result is quite unexpected in simple hierarchical models. ",
        "introduction": "\\label{sec:intro} Hierarchical galaxy formation models predict significantly different formation histories for early-type galaxies in cluster and low-density environments (Baugh, Cole \\& Frenk 1996, Kauffmann \\& Charlot 1998). In these models, present day clusters form from the highest peaks in the primordial density fluctuations where major mergers of dark matter halos, which harbour the first galaxies, rapidly produce bulge-dominated galaxies at high redshifts ($z\\ge2$).  The merging of galaxies and the infall of new cold gas cannot continue once the relative velocity dispersion of galaxies becomes large ($\\ge 500$~\\kms), \\ie\\/ the deep potential well of a cluster has been formed. Within this scenario it is possible to reconcile the characteristic ingredient of hierarchical galaxy formation, the merging process, with the observational finding that most stars in luminous cluster elliptical galaxies formed at $z \\ge 2$ (\\eg\\/ Aragon-Salamanca et~al. 1993, Ellis et~al. 1997, van Dokkum et~al. 1998).  In low-density regions, hierarchical models predict that (i) galaxies can accrete new cold gas and perhaps build up a stellar disk driving the morphology towards later types and (ii) that major mergers continue to take place at redshifts well below unity. As a result of this, early-type galaxies in low-density regions are able to incorporate stars formed at low redshifts and therefore should have, at the present day, younger luminosity-weighted ages than the equivalent cluster population \\cite{bau96,kau98,gov99,col00}. As shown in Figure~\\ref{fig:cdm_carlton} the models predict for cluster ellipticals a mean luminosity-weighted age of 9.6~Gyr (dashed line). The cluster S0s are predicted to be $\\sim$1~Gyr younger. Both ellipticals and lenticular galaxies in clusters show a weak trend in the sense that fainter galaxies are older. By contrast, the models predict that early-type galaxies in low-density regions should show a broad age distribution over the whole luminosity range with a mean age of 5--6~Gyr.  \\begin{figure} \\psfig{file=plot_sim_01.ps,width=8.5cm} \\caption[]{\\label{fig:cdm_carlton} Distributions of luminosity-weighted ages for early-type galaxies at $z=0$, predicted by hierarchical galaxy formation models \\cite{bau96,col00}. The four panels show predictions for faint and bright early-type galaxies in clusters (\\ie\\/ residing in dark matter halos with $>10^{14}$~M$_\\odot$) and low-density environments (\\ie\\/ in dark matter halos with $<10^{13}$~M$_\\odot$). Ellipticals and lenticular galaxies are represented by solid lines and dotted lines, respectively. The vertical dashed lines indicate the mean age of cluster ellipticals (\\ie\\/ 9.6~Gyr). For the predictions we adopt a cosmology of $H_0=72$~\\kmsM\\/ and $q_0=0.3$.} \\end{figure}  Recent observational efforts to investigate the formation of early-type galaxies in low-density regions, and test the model predictions discussed above, have largely focused on HST imaging of galaxies at redshifts $0.1 \\le z \\le 1$. When early-type galaxies (`spheroidals') are selected by morphology alone, no significant evidence of a decline in comoving number density with look-back time is found \\cite{men99}. At face value this suggests an early formation epoch.  However, some of these galaxies show colours which are too blue to be consistent with the predictions of a simple high-redshift monolithic-collapse model \\cite{men99}.  Furthermore, there is evidence that the dispersion in colour among field ellipticals is larger than for the equivalent cluster population at $z \\simeq 0.55$ (Schade et~al.  1999; see also Larson, Tinsley \\& Caldwell 1980 for a low redshift analogy). Schade et~al. also found spectroscopic evidence of ongoing star formation in the field, since about one-third of the field elliptical galaxies show \\oii\\/ lines with equivalent widths in excess of 15 \\AA. These results are further supported by the study of Menanteau, Abraham \\& Ellis (2001) who investigated the {\\em internal}\\/ colour variations of faint spheroidals in the HDFs. They find that at least one third of the galaxies show strong variations in internal colour, mostly showing centrally located blue cores, and conclude that at $z \\simeq 1$ approximately 50\\% of the field spheroidals experience episodes of star formation.  Recent fundamental plane studies of field early-type galaxies \\cite{vdok01,treu01} find, within their observational errors, no significant difference between cluster and field samples at $z \\le 0.5$. However, both samples provide evidence that the stars in field early-types are marginally younger than the equivalent cluster population, although most of the stars must have formed at high redshifts. van Dokkum et~al. conclude that their measurement of the evolution of the $M/L_B$ ratio with redshift is inconsistent with the predictions of semi-analytical models for galaxy formation \\cite{dia01}. While the models predict a systematic offset between field and cluster in $M/L_B$ at all redshifts (the field being brighter at a given mass), van Dokkum et~al. find no significant offset in their data.  There are a few investigations of the nearby population of early-type galaxies in low-density regions. de~Carvalho \\& Djorgovski (1992) investigate the properties of field and cluster early-type galaxies using a subset of the `7~Samurai' sample \\cite{fab87} and the data from Djorgovski \\& Davis (1987). They conclude that field ellipticals show more scatter in their parameters than cluster galaxies indicating the presence of younger stellar populations in the field.  Silva \\& Bothun (1998) investigate a sample of nearby early-type galaxies, specifically including galaxies with disturbed morphologies such as shells and tidal tails (see also Schweizer \\& Seitzer 1992). From their analysis of near-IR colours they conclude that there is little or no evidence for an intermediate age (1--3~Gyr) population of significant mass ($>$10\\%) in their sample, irrespective of morphological details. Colbert, Mulchaey \\& Zabludoff (2001) have undertaken an imaging survey of 23 nearby isolated early-type galaxies, finding morphological evidence for recent merging (\\eg\\/ shells and tidal features) in 41\\% of the galaxies, as compared to only 8\\% in their comparison sample of group members. Bernardi et~al. (1998) investigate the Mg$_2$--$\\sigma$ relation in a large sample of early-type galaxies drawn from the ENEAR survey \\cite{dcos00} and find that there is a small difference in the zero-point between cluster and field galaxies. They interpret this offset as an age difference, in the sense that field galaxies are younger by $\\sim$1~Gyr. They conclude however, that the stars in both field and cluster early-type galaxies formed mostly at high redshifts.  One of the main obstacles for studies of field galaxies is the exact treatment of the selection process. There are many possible definitions for the term `field', and it is critical to account for the different selection criteria when comparing published studies. For example, while Colbert et~al. find only 30 isolated early-type galaxies in the RC3 (de Vaucouleurs 1991, within $cz < 9\\,900$~\\kms), the field sample of Bernardi et~al. comprises more than two thirds of the entire ENEAR catalogue (631 out of 931 galaxies within $cz < 7\\,000$~\\kms). Clearly, the definitions of what is a field galaxy differ widely even for nearby galaxy samples. It is even more difficult to compare medium- or high-redshift samples, where redshift data is sparse, and selection criteria for the field often ill-defined.   The present paper presents a high-quality spectroscopic study of the stellar populations of early-type galaxies in low-density environments. In Section~\\ref{sec:sample} we describe the precise and reproducible selection criteria according to which our new sample is selected. The observations and basic data reduction processes are outlined in Section~\\ref{sec:data}, in which we also describe the verification and refinement of the galaxy sample. Our measurements of absorption (and emission) line strengths are presented in Section~\\ref{sec:linestrengths}. The principal results, detailed in Section~\\ref{sec:results}, derive from analysis of (i) the Mg--$\\sigma$ relation, (ii) the luminosity-weighted ages and metallicities and (iii) the [Mg/Fe] abundance ratios. We discuss these results in Section~\\ref{sec:discussion}, relating them to previous studies, and comparing with the expectations from hierarchical scenarios for galaxy formation. Our conclusions are presented in Section~\\ref{sec:conclusion}.  Throughout this paper, for model predictions as well as observations, luminosities and physical scales are computed for a Hubble constant $H_0=72$~\\kmsM. The adopted deceleration parameter is $q_0=0.3$; using a negative $q_0$, as preferred by SN~Ia data \\cite{riess98,perl99}, would have negligible effects on these calculations. ",
        "conclusions": "\\label{sec:conclusion} We have presented an analysis of a sample of nine nearby early-type galaxies (3 elliptical, 6 lenticular), selected from a redshift survey to reside in low-density regions. The sample is drawn from a sky-area of $\\sim$700 deg$^2$, to a redshift limit of 7\\,000~\\kms. Our stringent selection criteria allow only up to two neighbours within a search radius of 1.3~Mpc ($H_0 = 72$~\\kmsM, q$_0 = 0.3$) and $\\pm350$~\\kms. While the sample size is small, we emphasize that the multiple, well-defined selection criteria guarantee that these galaxies are of E/S0 morphology, and reside in large-scale environments of very low density. We have investigated the Mg$_2$--$\\sigma$ relation, and the luminosity-weighted age, metallicity and abundance ratio distribution. Our results have been compared to early-type galaxies in the Fornax cluster and with the predictions for hierarchical galaxy formation.  The principal conclusions of our study are as follows:  \\begin{itemize}  \\item Elliptical and lenticular galaxies are rare in the `field', and account for only $\\approx$8\\% of the galaxy population in the lowest-density environments. The existing small samples of early-type galaxies in low-density environments give no evidence for any significant departure from the luminosity function of E/S0 galaxies in denser environments.  \\item Five out of nine (56\\%) sample galaxies show disturbed morphologies (\\ie\\/ tidal tails or debris, blue circumnuclear rings), which is interpreted as evidence of late merger events in these galaxies. Furthermore seven galaxies show close neighbours of which five exhibit signs of ongoing interaction with the main galaxy.  \\item Compared to the cluster galaxies, our sample shows both a marginally higher incidence of \\oiii\\/ emission, and slightly stronger emission where present. However, we do not find galaxies with significant ongoing star-formation (\\eg\\/ \\oii\\/ $>10$~\\AA). Relative to studies of early-type galaxies at intermediate-redshift, this indicates a significant decline in star-formation activity compared to the field population at a redshift of $z=0.5$.  \\item The Mg--$\\sigma$ relation in low-density regions is indistinguishable from that of cluster E/S0s. However, Mg--$\\sigma$ alone is a poor diagnostic tool for detecting differences in star-formation history: intrinsic (anti-)correlations between age, metallicity and abundance ratios have degenerate effects which can conspire to maintain the scatter and zero-point of the relation.  \\item Early-type galaxies in low-density environments exhibit a broad distribution of luminosity-weighted ages, being on average younger than cluster ellipticals, while the distribution is more similar to that of lenticular galaxies in clusters. Taking the early-type galaxy population as a whole, the low-density regions harbour galaxies with 2--3~Gyr younger luminosity-weighted ages than their brethren in clusters. This result is robust against the specific selection of the comparison sample since imposing an upper $\\sigma$ limit does not change our result significantly while matching the luminosity distribution at the bright end increases the significance of our result. The younger ages of early-type galaxies in low-density regions is predicted by hierarchical galaxy formation, where the field population forms later and also experiences merger-induced star-formation episodes at lower redshifts.  We note that the youngest galaxies in our sample (LDR\\,08, 33, 34) all show clear signs of interaction, or show blue rings near the nucleus, the latter being suggestive of gaseous accretion.  \\item The luminosity-weighted metallicities of E/S0s in low-density environments are larger than for cluster members of similar luminosity. This effect ($\\Delta$[Fe/H]~$\\approx$0.2~dex) is not seen in semi-analytic models, which predict an offset in the opposite sense, at least for the luminosity range probed by our data. Furthermore, the current generation of semi-analytic models cannot explain the observed mass-metallicity relation of bright cluster members without compromising the reproduction of other observables, such as the luminosity function. These disagreements between the observations and semi-analytic model predictions highlight important shortcomings in the detailed treatment of the star-formation processes in present models.  \\item The E/S0 galaxies in our low-density sample exhibit mostly super-solar [Mg/Fe] ratios. The the non-solar [Mg/Fe] values rise with central velocity dispersion, following a trend similar to that of cluster members. By contrast, hierarchical galaxy formation models predict approximately solar abundance ratios, at least for the brightest galaxies in low-density regions.  \\end{itemize}  In summary, this study, though for a small sample, finds results consistent with one of the central predictions of hierarchical galaxy formation models: the formation of early-type galaxies continues to $z\\la1$ in low-density environments, while those in clusters formed most of their stars at $z\\ga2$. On the other hand, our results underscore the models' present failure to reproduce the observed luminosity--metallicity trends and their apparent dependence on environment. A future generation of models must also overcome the stiff challenge of generating super-solar [Mg/Fe] ratios, even in galaxies formed by late-time merging of potentially spiral galaxies with extended star-formation histories."
    },
    "0207/cond-mat0207707.txt": {
        "abstract": "\\noindent We discuss basic statistical properties of systems with multifractal structure. This is possible by extending the notion of the usual Gibbs--Shannon entropy into more general framework - R{\\'e}nyi's information entropy. We address the renormalization issue for R{\\'e}nyi's entropy on (multi)fractal sets and consequently show how R{\\'e}nyi's parameter is connected with multifractal singularity spectrum. The maximal entropy approach then provides a passage between R{\\'e}nyi's information entropy and thermodynamics of multifractals. Important issues such as R\\'{e}nyi's entropy versus Tsallis--Havrda--Charvat entropy and PDF reconstruction theorem are also studied. Finally, some further speculations on a possible relevance of our approach to cosmology are discussed. \\\\   \\vspace{3mm} \\noindent PACS: 65.40.Gr, 47.53.+n, 05.90.+m     \\\\ \\noindent {\\em Keywords}: R{\\'e}nyi's information entropy; multifractals; Tsallis--Havrda--Charvat entropy; MaxEnt \\draft ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The past two decades have witnessed an explosion of activity and progress in both equilibrium and non--equilibrium statistical physics. The catalyst has been the massive infusion of ideas from information theory, theory of  chaotic dynamical systems, theory of critical phenomena, and quantum field theory. These ideas include the generalized information measures, quasi--periodic and strange attractors, fully developed turbulence, percolation, renormalization of large--scale dynamics,  and attractive, albeit speculative, ideas about  quark--gluon plasma formation and dynamics. It is the purpose of this paper to proceed in this line of development. The issue at the stake is what modifications in statistical physics one should expect when dealing with systems with varied fractal dimension - multifractals. The view which we present here hinges on two mutually interrelated concepts, namely on R{\\'e}nyi's information entropy\\cite{Re1,Re2} and (multi)fractal geometry. In this connection we would like to stress that in order to exhibit the link between R\\'{e}nyi information entropies and (multi)fractal systems as generally as possible we do not put much emphasize on the important yet rather narrow class of (multi)fractal systems - chaotic dynamical systems.  \\vspace{3mm}  One of the fundamental observations of information theory is that the most general functional form for the mean transmitted information (i.e., information entropy) is that of R{\\'e}nyi. In Section II we briefly outline R{\\'e}nyi's proof and discuss some fundamentals from information theory which will show up to be indispensable in following sections. We show that with certain mathematical cautiousness Shannon's entropy can be viewed as a special example of R{\\'e}ny's entropy in case when R{\\'e}nyi's parameter $\\alpha \\rightarrow 1$. We also address the question of the status of Tsallis--Havrda--Charvat (THC) entropy\\cite{HaCh,Ts1} in the framework of information theory.  \\vspace{3mm}  Although R{\\'e}nyi's information measure offers very natural - and maybe conceptually the cleanest - setting for the entropy, it has not found so far as much applicability as Shannon's (or Gibbs's) entropy. The explanation, no doubt, lies in two facts; ambiguous renormalization of R{\\'e}nyi's entropy for non--discrete distributions and little insight into the meaning of R{\\'e}nyi's $\\alpha$ parameter. Surprisingly little work has been done towards understanding both of the former points. In Section III we aim to address the first one. We choose, in a sense, a minimal renormalization prescription conforming to the condition of additivity of independent information. R{\\'e}nyi's entropy thus obtained is then directly related to the information content (``negentropy\").  \\vspace{3mm}  To clarify the position of R{\\'e}nyi's entropy in physics, or in other word, to find the physical interpretation for $\\alpha$ parameter, we resort in Section IV to systems with a multifractal structure. Such systems are very important and highly diverse, including  the turbulent flow of fluids\\cite{Pa1,ArAr1}, percolations\\cite{Ah1}, diffusion--limited aggregation (DLA) systems\\cite{NST1}, DNA sequences\\cite{ZGU1}, finance\\cite{VVA1}, and string theory\\cite{MMW1}. Using the reconstruction theorem we argue that in order to obtain a ``full\" information about a (multi)fractal system  we need to know R{\\'e}nyi's entropies to all orders. Still, for discrete spaces and simple metric spaces (like ${\\mathbb{R}}^d$) we find that the contribution from Shannon's entropy dominates over all other R{\\'e}nyi entropies. We further show that from the maximal entropy (MaxEnt) point of view, extremizing the Shannon entropy on a multifractal is equivalent to extremizing directly Renyi's entropy without invoking the multifractal structure explicitly. Application of this result to a cosmic strings network will be presented elsewhere\\cite{AJS1}.  %for such systems it is natural to identify R{\\'e}nyi's parameter %with the corresponding Hausdorff dimension.  \\vspace{3mm}  %\\noindent In Section 4 we shall demonstrate that R{\\'e}nyi's entropy %can be naturally accommodated in thermodynamics provided we %introduce the notion of {\\em zooming energy} - i.e., extensive variable %conjugated to $\\alpha$ with respect to energy $U$. In this setting %we derive the generalized Maxwell thermodynamic relations.  %\\vspace{3mm}  We close with Section V where we present some speculations on the relevance of the outlined approach to string cosmology and quantum mechanics. % %Namely, we try to set cosmic strings thermodynamics. On such a %basis we argue that the critical (phase transition) temperature at %which strings tend to fragment into smallest allowed loops, while %large loops become exponentially suppressed - i.e., Hagedorn %temperature, is lower than  that predicted by the standard %thermodynamics. For reader's convenience we supplement the paper with eight appendices which clarify some finer mathematical manipulations.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0207/astro-ph0207262_arXiv.txt": {
        "abstract": "We study the transition of nuclear matter to strange quark matter (SQM) inside neutron stars (NSs). It is shown that the influence of the magnetic field expected to be present in NS interiors  has a dramatic effect on the propagation of a laminar deflagration (widely studied so far), generating a strong acceleration of the flame in the polar direction. This results in a strong asymmetry in the geometry of the just formed core of hot SQM which  resembles a cylinder orientated in the direction of the magnetic poles of the NS. This geometrical asymmetry gives rise to a bipolar emission of the thermal neutrino-antineutrino  pairs produced in the process of SQM formation. The $\\nu {\\bar \\nu}$  annihilate into $e^+ e^-$ pairs just above the polar caps of the NS giving  rise to a relativistic fireball, thus providing a suitable form of energy transport and conversion to $\\gamma$-emission that may be associated to short gamma ray bursts (GRBs). ",
        "introduction": "The transition to SQM is expected to occur in NSs if SQM has lower energy per baryon than ordinary nuclear matter \\cite{bodmer71,terazawa79,chin79,witten84}. This has strong consequences in the astrophysics of compact stars since, if SQM is the true ground state of strongly interacting matter rather than $^{56}$Fe,  compact objects could be strange stars (SS) instead of NSs. Some tentative strange star candidates are the compact objects associated with the X-ray bursters GRO J1744-28 \\cite{Cheng98}, SAX J1808.4-3658 \\cite{Li99a}, and  the  X-ray pulsar Her X-1  \\cite{Dey98}. In addition, the observed high and low frequency  quasi periodic oscillations in the atoll source 4U 1728-34 have been shown to be more consistent with a strange star nature \\cite{Li99b}. A number of different mechanisms has been proposed for  NM  to SQM conversion \\cite{aarhus} inside the star.  All them are based on the formation of a ``seed'' of SQM inside the NS. A possible mechanism is the so-called strangelet contamination, where a seed of SQM from the ISM enters a NS and converts it to a SS. In another possible scenario,  a seed of SQM forms in the core of a NS as a result of the increase of the central density above the critical density for deconfinement phase transition. A way to do this is through mass accretion onto the NS in a binary stellar system. In a third mechanism, a seed of SQM may naturally form inside a newly born neutron star from a core-collapse supernova explosion after a deleptonization timescale \\cite{mnras1,mnras2}. No matter which of these mechanisms are actually triggering the NS to SS transition, once the first seed of SQM is produced inside the NS, it will propagate as a combustion swallowing neutrons, protons and hyperons. The transition to SQM inside the burning front has actually two stages. Deconfinement driven by {\\it strong} interactions first liberates quarks confined inside hadrons. The just deconfined quark phase has a certain finite strangeness due to the presence of strange hadrons in NS matter. However, the composition of the just deconfined phase (with u,d and s quarks) is not in beta equilibrium and consequently chemical equilibrium is reached by {\\it weak} interactions. It is in this phase that a large amount of energy is released and then $\\nu$'s are copiously produced.    The type of diffusion-driven combustions studied so far by several authors \\cite{Slow1,Slow2} actually correspond to laminar deflagrations (slow combustions). Whether the conversion process remains forever as a deflagration, either laminar or turbulent; or jumps to the detonation regime, thus driving an explosive transient \\cite{Comb1,Comb3,Comb2} has been debated in the literature. The situation is closely analogue to the much more studied thermonuclear combustions leading to type Ia supernovae. In addition to the lack of detailed studies on the mode(s) of combustion (laminar/turbulent deflagration or detonation), there is (to the best of our knowledge) no calculation of the influence of ubiquitous magnetic fields expected to be promptly generated  or already present in the NS as a fossil. We shall consider hereafter an initially laminar deflagration in presence of a $B$-field. As stated, previous works  \\cite{Slow1,Slow2} have calculated the velocity of the laminar deflagration by considering the diffusion of $s$-quarks as the main agent for the progress of the conversion. Given that the combustion is idealized to happen near the $T = 0$ limit (which is small when  compared with the chemical potential of quarks), these diffusion-limited \"cold\" models are reasonable for this purpose. The general result, which also holds when temperature corrections and full non-linearity are considered, is that the laminar velocity of the front is relatively slow ($v_{lam} \\leq 10^{4} cm/s$) although this speed may be uncertain by several orders of magnitude (Olinto 1991). This velocity is a direct consequence of  both the timescale for weak decays that create $s$-quarks ($10^{-8} s$) and the physics of the diffusion process. ",
        "conclusions": ""
    },
    "0207/astro-ph0207054_arXiv.txt": {
        "abstract": "We present an improved uvby-metallicity relation calibrated for F,G, and early K dwarfs, and an analogous uvby-$T_{eff}$ relation, both derived using a Levenberg-Marquardt minimization scheme. Our calibrations are based on 1533 stars which appear in both the Cayrel de Strobel (2001) metallicity compilation, and in the Hauck-Mermilliod (1998) catalog of uvby photometry. We also examine the speculative possibility of using uvby photometry to produce a uvby-planeticity calibration. We conclude that while there is likely no strong photometric indicator of the presence or absence of short-period planets, the possibility of a spectroscopic indicator of planeticity is well worth examining. ",
        "introduction": "The number of known extrasolar planets is approaching 100, and more planets are being discovered every month.\\footnote {For the latest catalog of planets, see http://www.exoplanets.org} A remarkable empirical correlation has emerged from the aggregate of planets: The parent stars of the detected extrasolar planets appear to be significantly metal-rich in comparison to field FGK dwarfs in the solar neighborhood. This correlation was first extensively discussed by Gonzalez (1997), and has been further studied by a number of authors; a partial list might include Gimenez (2000), Laughlin (2000), Santos, Israelian \\& Mayor (2001), Murray \\& Chaboyer (2001), and Reid (2002).  These studies all agree that the fraction of stars containing readily detectable extrasolar planets (e.g. having $K>15$  m/s, and $P<5$ yr) increases substantially with increasing stellar metallicity. A search strategy geared toward the most metal-rich stars should thus detect short-period planets at a far greater rate than a conventional volume-limited survey.  One might argue that a metallicity-based planet search strategy will result in a haphazard accumulation of short-period planets, while simultaneously skewing the overall census of planets in the solar neighborhood. Indeed, it is important to continue with volume-limited surveys in order to round out the unbiased statistical distribution of planetary systems. The discovery of additional short-period planets, however, has important ramifications. For the following reasons, it is desirable to locate as many short-period planets as possible: ({\\bf 1}) The presence of a short-period planet is an excellent indicator that additional planets are present in the system. Currently, for example, 5 of 12 short-period planets observed at Lick Observatory show evidence for additional companions (Fischer et al 2001). Recently, it has been realized that some multiple planet systems such as GJ 876 allow for unambiguous determination of planet masses and orbital parameters (Laughlin \\& Chambers 2001, Rivera \\& Lissauer 2001). ({\\bf 2}) The current census of short-period extrasolar planets shows an interesting concentration of objects in the period range between roughly 2.98 and 3.52 days (8 out of 20 planets with periods $P<20$ d). While this pile-up is probably a consequence of disk-protoplanet tidal migration (Lin, Bodenheimer \\& Richardson 1996), the stopping mechanism is not at all well understood. It is thus worthwhile to find more short-period planets to better delineate the minimum planetary period. ({\\bf 3}) Geometric arguments indicate that 10\\% of Hot Jupiter-type planets transit the face of the parent star. A large increase in the rate of short-period planet detections thus translates into a significant increase in the number of transits discovered. Transiting planets permit accurate mass and radius determinations, and are extremely important data points within the overall theory of giant planets.  High-resolution spectroscopy provides the most precise method for flagging the metal-rich stars which have a high probability of bearing a detectable planet. Unfortunately, the number of field stars with spectroscopically determined metallicities is relatively small. The most recent compilation from the literature by Cayrel de Strobel, Soubiran, \\& Ralite (2001; hereafter CSR) contains 6534 metallicity determinations for 3356 stars. The majority of single, metal-rich, main sequence and subgiant FGK stars in the CSR compilation are already under radial velocity survelliance. A larger pool of high metallicity candidates is needed.  A study by Laughlin (2000) used the Hipparcos database (Perryman et al 1997), the Hauck-Mermilliod (1998) compilation of uvby photometry, and the Schuster \\& Nissen (1989; hereafter SN) uvby-[Fe/H] calibration to to identify 200 metal rich stars from a parent population of roughly 10,000 K,G, and late F-type main sequence stars (having $7.8<V<10$) in the immediate ($d<100{\\rm pc}$) solar neighborhood. Nineteen of these stars were added to the Keck Radial Velocity survey in the Summer of 2000. Three of these nineteen stars soon proved to be accompanied by planets (Vogt et al 2001, Butler et al 2002). Metallicity-based search strategies are therefore clearly optimizing the planet-detection efficiency.  The SN calibrations for F and G spectral types were based on 219 calibration stars spanning a range of metallicities from $-3.5 <  [{\\rm Fe/H}] < 0.4$. This very useful calibration, which has an overall standard deviation of 0.16 dex per star, is wholly empirical, and has the advantage of using only the standard uvby indices, rather than the differential indices, $\\delta m_{1}=m_{1}^{hyades}-m_{1}^{star}$, and $\\delta c_{1} = c_{1}^{hyades}-c_{1}^{star}$, which are used in the earlier calibrations of Crawford (1975) and Olsen (1984). The SN calibration has been widely used to investigate metallicity trends among solar neighborhood stars (e.g. Rocha-Pinto \\& Maciel (1996), or Haywood (2001). As recently emphasized by Twarog, Anthony-Twarog and Tanner (2002), however, the SN calibration systematically underestimates $[{\\rm Fe/H}]$ by $\\sim0.2$ dex for stars with spectroscopic metallicities $[{\\rm Fe/H}]>0.1$. The SN calibration is thus not optimally suited for compiling an accurately ranked catalog of high metallicity stars.  Our main goal in this letter is to make the uvby-metallicity calibration more discerning and accurate among the highest metallicity stars. We also use this forum to briefly introduce the concept of a ``planeticity calibration'', in which stellar properties are correlated directly with the detectable presence of an extrasolar planet. ",
        "conclusions": ""
    },
    "0207/astro-ph0207324_arXiv.txt": {
        "abstract": "It is shown that the commonly held view of a sharp Fe K edge must be modified if the decay pathways of the series of resonances converging to the K thresholds are adequately taken into account. These resonances display damped Lorentzian profiles of nearly constant widths that are rapidly smeared to impose continuity through the threshold. By modeling the effects of K damping on opacities, it is found that the broadening of the K edge grows with the ionization level of the plasma and that the appearance at high ionization of a localized absorption feature at 7.2 keV is identified as arising from the K$\\beta$ unresolved transition array. ",
        "introduction": "X-ray absorption and emission features arising from iron K-shell processes are of practical importance in high-energy astrophysics. This is primarily due to the iron cosmic abundance, but also to the relatively unconfused spectral region where they ubiquitously appear. Although the observational technology in X-ray astronomy is still evolving, many of such features are being resolved and consequently exploited in plasma diagnostics. In this respect, they are naturally grouped according to their origin, i.e. bound--bound or bound--free ionic transitions, and much of the interpretation of the latter has relied on atomic calculations \\citep{verner95,berrington97,donnelly00,berrington01} that predict a sharp increase of the photoabsorption cross section at the K-shell threshold. The purpose of this communication is to emphasize that this commonly held view is incorrect due to an oversimplified treatment of the decay pathways of the resonances converging to this limit, and that previous astrophysical inferences from K-edge structures should thus be revised. ",
        "conclusions": ""
    },
    "0207/astro-ph0207112_arXiv.txt": {
        "abstract": "We present 21-cm HI line and optical R-band observations for a sample of 26 edge-on galaxies. The HI observations were obtained with the Westerbork Synthesis Radio Telescope, and are part of the WHISP database (Westerbork HI Survey of Spiral and Irregular Galaxies). We present HI maps, optical images, and radial HI density profiles. We have also derived the rotation curves and studied the warping and lopsidedness of the HI disks. 20 out of the 26 galaxies of our sample are warped, confirming that warping of the HI disks is a very common phenomenon in disk galaxies. Indeed, we find that all galaxies that have an extended HI disk with respect to the optical are warped. The warping usually starts around the edge of the optical disk. The degree of warping varies considerably from galaxy to galaxy. Furthermore, many warps are asymmetric, as they show up in only one side of the disk or exhibit large differences in amplitude in the approaching and receding sides of the galaxy. These asymmetries are more pronounced in rich environments, which may indicate that tidal interactions are a source of warp asymmetry. A rich environment tends to produce larger warps as well. The presence of lopsidedness seems to be related to the presence of nearby companions. ",
        "introduction": "It has long been known that many galaxies have warped disks. The phenomenon occurs at large galactocentric radii, starting close to the edge of the optical disk. This makes 21-cm line observations essential for the study of warps: there are galaxies that have no optical warp and yet exhibit extraordinary warps in their outer HI layers.  The first indication that the HI distribution of disk galaxies was sometimes bent or warped came in 1957 from observations of our own Galaxy (Burke, 1957; Kerr, 1957). Early calculations showed that the tidal field due to the Magellanic Clouds was unable to account for the warp of the Galaxy (Burke, 1957; Kerr, 1957; Hunter \\& Toomre, 1969), and so it was regarded as a rare, uncommon phenome\\-non. Later on, Sancisi (1976) studied the HI layer of 5 edge-on galaxies and discovered 4 of them to have warps. This result indicated that warps are very common among galaxies. This discovery was not expected, because three of Sancisi's galaxies were quite isolated and thus presented a problem with the tidal-warp scenario. The high incidence of warps among spiral galaxies has later been confirmed with larger samples both in the optical \\cite{ss,res} and the HI \\cite{bosma}. The ubiquity of warps means that either warps are long-lasting phenomena or they are transient but easily and often excited. Neither of these has been proven and no satisfactory explanation has been found yet. A study by Briggs (1990) of a sample of 12 galaxies showed that: (1) warps start around R$_{25}$ - R$_{\\rm H_\\circ}$; (2) the line of nodes of the warp is straight inside R$_{25}$, and is a leading spiral outside.  The absence of a mechanism that would regularly generate warps and the fact that the line of nodes of the observed warps does not show a severe winding effect \\cite{briggs} suggest that warps are long-lived. The main issue in trying to explain a long lasting warp is differential precession, which in a few rotation periods will destroy the coherent warp pattern (this is similar to the winding problem of material spiral arms). The differential precession arises from the non-sphericity of the potential \\cite{avner}. However, the survival time of a coherent warp pattern can be larger if the halo in which the disk is embedded deviates only slightly from spherical symmetry \\cite{sanders}. Likewise, a halo with a flattening that decreases outwards in the right way would be able to maintain a coherent warped disk for a Hubble time \\cite{petrou}.  Extending the work of Hunter \\& Toomre (1969), Sparke \\& Casertano (1988) studied long-lived warping modes of a galactic disk inside an oblate halo potential, where no winding occurs. They found that the combined torque from a halo and a self-gravitating disk allows a configuration in which the precession frequency of the warp is the same at all radii, thus allowing the warp to maintain its shape unchanged for many rotation periods. One of their assumptions was that of a fixed halo potential that does not react to the potential of the disk.  The back-reaction of the precessing disk onto the halo, analogous to dynamical friction, was investigated by Dubinski \\& Kuijken (1995), Nelson \\& Tremaine (1995), and Binney et al. (1998) after an earlier investigation in terms of WKB density waves by Bertin and Mark (1980). It turns out that it has a very strong influence on the evolution of the warp, causing the disk and halo to reorient to a common plane of symmetry in a few orbital times, and making the warp decay.  \\begin{center} \\begin{table*} \\begin{center} \\begin{tabular}{cclcccrccrr} \\hline UGC & NGC & \\multicolumn{1}{c}{type} & R$_{25}$ & R$_{\\rm opt}$ & {\\it i} & \\multicolumn{1}{c}{PA} & B$_c$ & $\\Delta$B & \\multicolumn{1}{c}{d$_{\\rm V}$} & \\multicolumn{1}{c}{d$_{\\rm TF}$} \\\\ &  &  & $'$ & $'$ & $^\\circ$ & $^\\circ$ & & & Mpc  & Mpc \\\\ \\hline \\hline 1281 &    & Sc  &  2.23 &  2.59 &  90.0 &  38 &  10.95 &  0.11 &  5.4 &  5.1 \\\\ 2459 &    & Scd &  1.24 &  2.02 &  90.0 &  62 &  ----- &  -----&  36.3& 36.3 \\\\ 3137 &    & Sbc &  1.78 &  2.66 &  90.0 &  74 &  13.61 &  1.00 &  18.3& 33.8 \\\\ 3909 &    & SBc &  1.32 &  1.39 &  90.0 &  82 &  13.57 &  0.77 &  17.7& 24.5 \\\\ 4278 &    & SBc &  2.27 &  2.69 &  90.0 &  172&  11.09 &  0.23 &  10.4 & 8.1 \\\\ 4806 &2770& Sc  &  1.83 &  2.41 &  77.0 &  148&  11.62 &  0.07 &  29.6 &21.0 \\\\ 5452 &3118& Sc  &  1.26 &  1.41 &  90.0 &  41 &  12.88 &  1.00 &  22.1 &21.7 \\\\ 5459 &   &  SBc &  2.38 &  2.67 &  90.0 &  132&  11.58 &  0.11 &  19.6 &15.9 \\\\ 5986 &3432& SBd &  3.44 &  3.52 &  81.2 &  38 &  10.36 &  0.08 &  8.9 &  8.5 \\\\ 6126 &3510& SBcd&  2.03 &  2.13 &  90.0 &  163&  11.17 &  0.33 &  8.8 &  8.8 \\\\ 6283 &3600& Sab &  2.03 &  2.15 &  87.0 &  3 &  11.51 &  0.59 &  11.4 & 11.3 \\\\ 6964 &4010& SBcd&  2.10 &  2.13 &  85.5 &  66&  11.79 &  0.10 &  16.6 & 16.9 \\\\ 7089 &   &  Sc  &  1.61 &  2.73 &  87.2 &  36&  12.29 &  0.14 &  13.2 & 11.6 \\\\ 7090 &   4096  &  SBc    \t &  3.30 &  3.96 &  80.9 &  20 &  10.16 &  0.19 &  9.4 &  10.2 \\\\ 7125 &   &  SBd    \t &  2.35 &  2.26 &  90.0 &  85 &  12.44 &  0.92 &  19.6 &  12.6 \\\\ 7151 &   4144  &  SBc    \t &  3.06 &  3.34 &  81.6 &  104 &  10.72 &  0.16 &  4.3 &  6.0 \\\\ 7321 &   &  Sc     \t &  2.78 &  2.91 &  90.0 &  82 &  11.99 &  0.16 &  4.0 &  14.9 \\\\ 7483 &   4359  &  SBc    \t &  1.75 &  2.49 &  80.8 &  108 &  12.31 &  0.80 &  22.4 &  17.6 \\\\ 7774 &   &  Sc     \t &  1.75 &  1.90 &  90.0 &  102 &  12.98 &  0.86 &  7.3 &  20.6 \\\\ 8246 &   &  SBc    \t &  1.71 &  1.86 &  90.0 &  83 &  13.62 &  1.00 &  11.7 &  19.4 \\\\ 8286 &   5023  &  Sc     \t &  3.03 &  3.58 &  90.0 &  28 &  11.09 &  0.13 &  6.3 &  8.0 \\\\ 8396 &   5107  &  SBc    \t &  0.84 &  0.90 &  77.3 &  128 &  13.89 &  1.45 &  17.4 &  27.5 \\\\ 8550 &   5229  &  SBc    \t &  1.67 &  2.10 &  90.0 &  167 &  12.97 &  1.00 &  6.3 &  13.2 \\\\ 8709 &   5297  &  SBbc   \t &  2.72 &  3.02 &  83.6 &  148 &  11.13 &  0.21 &  37.3 &  19.8 \\\\ 8711 &   5301  &  SBc    \t &  2.02 &  2.49 &  87.2 &  151 &  11.98 &  0.22 &  25.8 &  22.5 \\\\ 9242 &   &  Sc     \t &  2.51 &  3.04 &  90.0 &  71 &  11.86 &  0.07 &  24.7 &  12.6 \\\\ \\hline \\end{tabular} \\caption{Optical properties of the galaxies in the sample: Galaxy name (UGC and NGC numbers), Hubble type, R$_{25}$, R$_{\\rm opt}$ (see Sect.~\\ref{ropt} for definition), inclination angle, position angle of the major axis, apparent b magnitude, corrected for inclination and galactic extinction, error in the magnitude, distance derived with a Virgocentric inflow model, and Tully-Fisher distance. All the data have been extracted from the LEDA database, except for the distances (see section \\ref{dist}). We have not used the corrected b magnitude of UGC 2459 because it is not reliable due to its low galactic latitude.} \\label{tab:opt0} \\end{center} \\end{table*} \\end{center}   More recently, several other mechanisms have been investigated. Debattista \\& Sellwood (1999) studied the warps generated by a disk embedded in a halo whose angular momentum vector is misaligned with that of the disk, and Ing-Guey \\& Binney (1999) built a model of a warp caused by a disk embedded in a halo which is constantly accreting material of angular momentum misaligned with that of the disk. Magnetic fields have also been proposed as a cause for warps. This hypothesis is mainly  based on the alignment of the warps of different galaxies that are close in space \\cite{bat}. Such alignment would be hard to explain in a warping model that is based only on the interaction of a disk and a halo of an individual galaxy without involving the surroundings.  In spite of the several models proposed, none is fully convincing, and many questions remain. Observationally the main problem lies in the fact that the current sample of warped galaxies is very small and inhomogeneous, and often the resolution of the data is poor. Furthermore, in many cases the warps are inferred by modeling the HI kinematics, using assumptions like axisymmetry and circularity of the orbits, which may be not completely correct.  Although most of the warps start outside the optical disk, optical warps do exist. In fact, Sanchez-Saavedra et al. (1990) claimed that the majority of galaxies have optical warps. A  study by Reshetnikov \\& Combes (1998) on optical warps showed the influence of the environment, in the sense that there are more optically warped galaxies in dense environments than in low-density regions. The relation between the optical and HI warps is not clear yet. There have been claims based on near infrared data (DIRBE) that the Galactic warp is more pronounced in the HI than in the old stellar population (Porcel, Battaner \\& Jim\\'enez-Vicente 1996), but an early cutoff of the old stellar disk \\cite{robin} can also explain the data. If magnetic fields have some influence in exciting and/or maintaining warps, the old stellar disk is expected to have a milder warp than the HI disk at every radius \\cite{jorge}.  Our goal in this paper is to present an analysis of a sample of late-type spiral galaxies to determine the occurrence rate of warps, and to study their properties (symmetry, dependence on environment, etc). We have selected a sample of edge-on galaxies that have been observed with the Westerbork Synthesis Radio Telescope (WSRT) for the WHISP project \\cite{whisp}. We have used optical R-band data as well \\cite{sm}, to explore possible links between the optical properties and the warps of our galaxies. These images also serve to study the environment of our sample galaxies and to identify possible satellite dwarf galaxies that may be orbiting them. ",
        "conclusions": "  \\begin{itemize}  \\item  We detect warps in 20 out of our 26 sample galaxies confirming that warping of the HI disks is a very common phenomenon in disk galaxies. In fact all galaxies that have an HI disk more extended than the optical are warped.  \\item  The amplitude of warps varies considerably from galaxy to galaxy. Also for a given galaxy there can be a large asymmetry in amplitude and shape between the two sides. A large number of warps in our sample are asymmetric. Most of the galaxies with both sides warped are antisymmetric (S shape warps). We only have two cases of U-type warps, and both galaxies are strongly interacting with nearby companions and are very disturbed.  \\item  The warping of the disks usually starts near the edge of the optical disk where the HI density drops down.  \\item  The connection between HI and optical warps is not clear. HI warps are found in general at larger radii than the optical ones, and as a consequence they probe a different region of the potential of the galaxy. A joint optical+HI study of warps could give important insights on the formation mechanism(s) of warps.  \\item  There seems to be a dependence of warps on environment in the sense that galaxies in rich environments tend to have larger and more asymmetric warps than galaxies in poor environments.  \\item  The presence of density lopsidedness (and in a weaker way that of ki\\-ne\\-matical lopsidedness) seems to be related to the presence of nearby companions.  \\end{itemize}"
    },
    "0207/astro-ph0207438_arXiv.txt": {
        "abstract": "We report the identification of cyclical changes in the orbital period of the eclipsing dwarf nova Z~Cha. We used times of mid-eclipse collected from the literature and our new eclipse timings to construct an observed-minus-calculated diagram covering 30 years of observations (1972-2002). The data present cyclical variations that can be fitted by a linear plus sinusoidal function with period $28\\pm 2$ yr and amplitude $1.0\\pm 0.2$ minute.  The statistical significance of this period by an F-test is larger than 99.9\\%. The derived fractional period change, $\\Delta P/P= 4.4 \\times 10^{-7}$, is comparable to that of other short-period cataclysmic variables (CVs), but is one order of magnitude smaller than those of the long-period CVs. Separate fits to the first and second half of the data lead to ephemerides with quite different cycle periods and amplitudes, indicating that the variation is not sinusoidal or, most probably, is not strictly periodic. The observed cyclical period change is possibly caused by a solar-type magnetic activity cycle in the secondary star. An incremental variation in the Roche lobe of the secondary star of $\\Delta R_{L2}/R_{L2} \\simeq 1.7 \\times 10^{-4}$ is required in order to explain both the observed period change and the modulation of the quiescent brightness previously reported by Ak, Ozkan \\& Mattei. ",
        "introduction": "Z~Cha is a short-period ($P_{orb}= 1.78$ hr) eclipsing cataclysmic variable (CV). In these binaries, a late-type star (the secondary) overfills its Roche lobe and transfers matter to a companion white dwarf (the primary). In most CVs the donor star has lower mass than the accreting star. Since conservative mass transfer in such situations would lead to an increase in the orbital separation (and therefore the cessation of mass transfer via Roche lobe overflow), the existence of CVs as mass-transfer binaries implies that they must continuously loose angular momentum in order to sustain the mass transfer process. As a consequence, the binary should evolve slowly towards shorter orbital periods (on time scales of $10^8-10^9$~yr). Possible mechanisms suggested for driving the continuous angular momentum loss are magnetic braking via the secondary star's wind (for $P_{orb}>3$ hr) and gravitational radiation (for $P_{orb}< 3$ hr) (King 1988). At very short periods, when the secondary star becomes fully degenerate ($M_2\\simlt 0.08\\;M_\\odot$), mass loss leads to an expansion of this star and reverses the secular trend, resulting thereafter in an increasing orbital period. However, the predicted mass transfer rate after this period minimum is low (\\.{M}$_2 \\simeq 10^{-12} \\; M_\\odot\\; yr^{-1}$) and few CVs are expected to be observed in such evolutionary stage (Warner 1995).  The secular evolution of the binary can in principle be detected by measuring the changes in the orbital period of eclipsing CVs. Eclipses  provide a fiducial mark in time and can usually be used to determine the orbital period (and its derivative) with high precision. However, attempts to measure the long-term orbital period decrease in CVs have been disappointing: none of the studied stars show a positive detection of an orbital period decrease (e.g., Beuermann \\& Pakull 1984). Instead, most of the well observed eclipsing CVs \\footnote{i.e., those with well-sampled observed-minus-calculated (O$-$C) eclipse timings diagram covering more than a decade of observations.} show cyclical period changes (e.g., Bond \\& Freeth 1988; Warner 1988; Robinson, Shetrone \\& Africano 1991; Baptista, Jablonski \\& Steiner 1992; Echevarria \\& Alvares 1993; Wolf et~al. 1993; Baptista et~al. 1995; Baptista, Catal\\'an \\& Costa 2000). The most promising explanation of this effect seems to be the existence of a solar-type (quasi- and/or multi-periodic) magnetic activity cycle in the secondary star modulating the radius of its Roche lobe and, via gravitational coupling, the orbital period on time scales of the order of a decade (Applegate 1992; Richman, Applegate \\& Patterson 1994). The relatively large amplitude of these cyclical period changes probably contributes to mask the low amplitude, secular period decrease.  Z~Cha seemed to be a remarkable exception in this scenario. The eclipse timings analysis of Robinson et~al. (1995) shows a conspicuous orbital period increase on a time scale of $P/|\\hbox{\\.{P}}|= 2 \\times 10^7$~yr, not only at a much faster rate than predicted ($\\simeq 10^9$~yr) but also with the opposite sign to the expected period decrease.  In this Letter we report new eclipse timings of Z~Cha which indicate a clear reversal of the period increase observed by Robinson et~al. (1995). The revised (O$-$C) diagram shows a cyclical period change similar to that observed in many other well studied eclipsing CVs. The observations and data analysis are presented in section \\ref{observa} and the results are discussed and summarized in section \\ref{discuss}. ",
        "conclusions": "\\label{discuss}  Our results reveal that the orbital period of Z~Cha is no longer increasing as previously found by Robinson et~al. (1995). Instead, the (O$-$C) diagram shows conspicuous cyclical, quasi-periodic changes of amplitude 1~min on a time-scale of about 28 yr.  Cyclical orbital period changes are seen in many eclipsing CVs (Warner 1995 and references therein). The cycle periods range from 4~yr in EX~Dra (Baptista et~al. 2000) to about 30~yr in UX~UMa (Rubenstein, Patterson \\& Africano 1991), whereas the amplitudes are in the range 0.1-2.5 min. Therefore, Z~Cha fits nicely in the overall picture drawn from the observations of orbital period changes in CVs.  If one is to seek for a common explanation for the cyclical period changes in CVs, then models involving apsidal motion or a third body in the system shall be discarded as these require that the orbital period change be strictly periodic, whereas the observations show that this is not the case (Richman et~al. 1994 and references therein). We may also discard explanations involving angular momentum exchange in the binary, as cyclical exchange of rotational and orbital angular momentum (Smak 1972; Biermann \\& Hall 1973) requires discs with masses far greater than those deduced by direct observations, and the time-scales required to allow the spin-orbit coupling of a secondary of variable radius are much shorter than the tidal synchronization scales for these systems ($\\sim 10^{4}$ yr, see Applegate \\& Patterson 1987).  The best current explanation for the observed cyclical period modulation is that it is the result of a solar-type magnetic activity cycle in the secondary star (Applegate \\& Patterson 1987; Warner 1988; Bianchini 1990). Richman et~al. (1994) proposed a model in which the Roche lobe radius of the secondary star $R_{L2}$ varies in response to changes in the distribution of angular momentum inside this star (caused by the magnetic activity cycle), leading to a change in the orbital separation and, therefore, in the orbital period. As a consequence of the change in the Roche lobe radius, the mass transfer rate \\.{M}$_2$ also changes. In this model, the orbital period is the shortest when the secondary star is the most oblate (i.e., its outer layers rotate faster), and is the longest when the outer layers of the secondary star are rotating the slowest.  The fractional period change $\\Delta P/P$ is related to the amplitude $\\Delta(O-C)$ and to the period $P_{mod}$ of the modulation by (Applegate 1992), \\begin{equation} \\frac{\\Delta P}{P}= 2\\pi\\; \\frac{\\Delta(O-C)}{P_{mod}}= 2\\pi\\; \\frac{A}{C} \\; . \\label{eq:pponto} \\end{equation} Using the values of $A$ and $C$ in Table~\\ref{zcha.efem}, we find $\\Delta P/P = 4.4 \\times 10^{-7}$. This fractional period change is comparable to those of other short-period CVs, but is one order of magnitude smaller than those of the CVs above the period gap ($\\Delta P/P \\simeq 2\\times 10^{-6}$) [Warner 1995].  The predicted changes in Roche lobe radius and mass transfer rate are related to the fractional period change by (Richman et~al. 1994), \\begin{equation} \\frac{\\Delta R_{L2}}{R_{L2}}= 39\\; \\left( \\frac{1+q}{q} \\right)^{2/3} \\left( \\frac{\\Delta\\Omega}{10^{-3}\\Omega} \\right)^{-1} \\frac{\\Delta P}{P} \\; , \\end{equation} and by, \\begin{equation} \\frac{\\Delta \\hbox{\\.{M}}_2}{\\hbox{\\.{M}}_2} = 1.22 \\times 10^5 \\; \\left( \\frac{1+q}{q} \\right)^{2/3} \\left( \\frac{\\Delta\\Omega}{10^{-3}\\Omega} \\right)^{-1} \\frac{\\Delta P}{P} \\; , \\end{equation} where $\\Delta\\Omega/\\Omega$ is the fractional change in the rotation rate of the outer shell of the secondary star involved in the cyclical exchange of angular momentum, $q\\; (= M_2/M_1)$ is the binary mass ratio, and the minus signs were dropped.  In the framework of the disc instability model (Smak 1984, Warner 1995 and references therein), the mass transferred from the secondary star will accumulate in the outer disc regions during the quiescent phase. Hence, changes in \\.{M}$_2$ will mainly affect the luminosity $L_{bs}$ of the bright spot where the infalling material hits the outer edge of the disc. The luminosity of the bright spot is $L_{bs} \\propto \\hbox{\\.{M}}_2$. The bright spot in Z~Cha contributes about 30 per cent of the total optical light (on an average over the orbital cycle) [Wood et~al. 1986]. Therefore, one expects that the changes in mass transfer rate lead to a modulation in the quiescent brightness of the system of $\\Delta m = (\\Delta L_{bs}/L) \\simeq 0.3 \\; (\\Delta$\\.{M}$_2$/\\.{M}$_2) \\simeq (0.06 - 0.12)$~mag for $\\Delta\\Omega/\\Omega = (1-2)\\times 10^{-3}$ (Richman et~al. 1994) and $q=0.15$ (Wood et~al. 1986).  Ak, Ozkan \\& Mattei (2001) reported the detection of cyclical modulations of the quiescent magnitudes and outburst intervals of a set of dwarf novae, which they interpreted as the manifestation of a magnetic activity cycle in their secondary stars. They found that the quiescent brightness of Z~Cha is modulated with an amplitude of $\\Delta m= 0.16$~mag on a time scale of 14.6~yr. Unfortunately, they do not list the epoch of maximum brightness.  The observed amplitude of the brightness modulation is larger than that predicted from the orbital period change with the assumption of $\\Delta\\Omega/\\Omega = (1-2)\\times 10^{-3}$. The model of Richman et~al. (1994) can account for both the measured period changes and the observed brightness modulation if $\\Delta\\Omega/\\Omega \\simeq 2.7 \\times 10^{-3}$. This yields $\\Delta R_{L2}/R_{L2} \\simeq 1.7 \\times 10^{-4}$.  The predicted change in the Roche lobe radius of the secondary star in Z~Cha is comparable to the observed change in the radius of the Sun as a consequence of its magnetic activity cycle (Gilliland 1981). The period of the quiescent brightness modulation is very close to half of the period of the observed $P_{mod}$ and may correspond to the first harmonic of a non-sinusoidal or non-strictly periodic period change.  The analysis of Ak et~al. (2001) covers only 18 years of observations. It would be interesting to see whether the 28~yr orbital period change also appear as a modulation in the quiescent magnitude in a dataset covering a larger time-interval. A simple and interesting test of the Richman et~al. (1994) model is to check its prediction that the maximum of the brightness modulation coincides with the minimum of the orbital period modulation, which occurred at $E \\simeq 5.2 \\times 10^4$ cycle (or about JD 2\\,444\\,150) for Z~Cha.  Finally, we remark that, if the magnetic activity cycle explanation is right, our confirmation that Z~Cha also shows cyclical period changes underscores the conclusion of Ak et~al. (2001) that even fully convective secondary stars possess magnetic activity cycles (and, therefore, magnetic fields)."
    },
    "0207/astro-ph0207330_arXiv.txt": {
        "abstract": "The eddy viscosity in the solar supergranulation layer is derived from the observed rotational shear by computing theoretical rotation laws for the outermost parts of the solar convection zone using the results from numerical simulations of rotating convection as input. By varying the eddy viscosity, the results can be tuned to match the observations. The value of $1.5 \\times 10^{13}$ cm$^2$/s found for the eddy viscosity is considerably larger than the eddy magnetic diffusivity derived from the sunspot decay. The results are checked by comparison of the horizontal cross correlations of the velocity fluctuations with the observed Ward profile. ",
        "introduction": "Sunspots rotate about 4\\% faster than the solar surface plasma  at all latitudes. Helioseismology reveals a corresponding maximum of the angular velocity rather close to the surface, as shown in Fig.~\\ref{fig1}. Such a clear subrotation of the outermost layer of the convection zone is easiest understood as the result  of angular momentum conservation of fluid elements with purely radial motions. In this domain of the solar convection zone, however, the gas motion is predominantly horizontal. Fluctuating fields with dominant horizontal intensity should simply produce superrotation rather than the observed subrotation (cf. R\\\"udiger 1989)  From the proper motions of sunspot groups, the faster of which tend to move toward the equator, Ward (1965) found the positive value of $\\approx 0.1 $ (deg / day)$^2 \\approx 2 \\times 10^7$ m$^2$/s$^2$ for the corresponding horizontal cross correlation on the northern hemisphere. More recent observations found smaller but always positive values (Nesme-Ribes, Ferreira, \\& Vince 1993).  \\begin{figure} \\includegraphics[width=8cm,height=5cm]{soho} \\caption{The internal solar rotation as found by helioseismology. Note the increase of the rotation rate with depth in the outermost layer. The maximum lies about 0.05 solar radii beneath the photosphere. (Graph courtesy NSF's National Solar Observatory).} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207040_arXiv.txt": {
        "abstract": "We calculate the coupled hydrostatic and ionization structures of spherically symmetric, pressure-supported gas clouds that are confined by gravitationally dominant dark-matter (DM) mini-halos and by an external bounding pressure provided by a hot medium. We focus on clouds that are photoionized and heated by the present-day background metagalactic field and determine the conditions for the formation of warm (WNM), and multi-phased (CNM/WNM) neutral atomic hydrogen (HI) cores in the DM-dominated clouds.   We consider $\\Lambda$CDM dark-matter halos with cuspy (NFW) and constant density (Burkert) cores. We compute models for a wide range of halo masses, total cloud gas masses, and external bounding pressures.  We present models for the pressure-supported HI structures observed in the Local Group dwarf irregular galaxies Leo A and Sag DIG. We find that the hydrogen gas becomes neutral for projected HI column densities exceeding $10^{19}$ cm$^{-2}$. We identify the HI cloud boundaries observed in Leo A and Sag DIG with the ionization fronts, and we derive an upper limit of $P_{\\rm HIM}/k \\lesssim 100$ cm$^{-3}$ K for the ambient pressure of the intergalactic medium in the Local Group. The observed HI gas scale heights in Leo A and Sag DIG imply characteristic DM densities of $1.2$ amu cm$^{-3}$ (or $0.03 M_\\odot$ pc$^{-3}$), consistent with the DM densities previously inferred via HI rotation curve studies of dwarf and low-surface brightness galaxies. Leo A and Sag DIG obey the scaling correlations that are expected for typical (median) DM halos in a $\\Lambda$CDM cosmology, provided the halos contain constant density cores, as suggested by Burkert.   We construct explicit ``mini-halo'' models for the multi-phased (and low-metallicity) compact high-velocity HI clouds (CHVCs). If the CHVC halos are drawn from the same family of halos that successfully reproduce the dwarf galaxy observations, then the CHVCs must be ``circumgalactic'' objects, with characteristic distances of 150 kpc. For such systems we find that multi-phased behavior occurs for peak WNM HI column densities between $2\\times 10^{19}$ and $1\\times 10^{20}$ cm$^{-2}$, consistent with observations. In contrast, if the CHVCs are ``extragalactic'' objects with distances $\\gtrsim 750$ kpc, then their associated halos must be very ``underconcentrated'', with characteristic DM densities $\\lesssim 0.08$ cm$^{-3}$, much smaller than expected for their mass, and significantly smaller than observed in the dwarf galaxies. Furthermore, multi-phased cores then require higher shielding columns. We favor the circumgalactic hypothesis. If the large population of CHVCs represent ``missing low-mass DM satellites'' of the Galaxy, then these HI clouds must be pressure-confined to keep the gas neutral. For an implied CHVC mini-halo scale velocity of $v_s=12$ km s$^{-1}$, the confining pressure must exceed $\\sim 50$ cm$^{-3}$ K. A hot ($\\sim 2\\times 10^6$ K) Galactic corona could provide the required pressure at 150 kpc.  Our static mini-halo models are able to account for many properties of the CHVCs, including their observed peak HI columns, core sizes, and multi-phased behavior. However, important difficulties remain, including the presence in some objects of extended low column density HI wings, and H$\\alpha$ emission line fluxes in several CHVCs that are significantly larger than expected. ",
        "introduction": "High velocity clouds (HVCs) are atomic hydrogen (HI) clouds with radial velocities inconsistent with gas in differential circular rotation in the Galactic disk. Since their discovery via 21 cm observations by Muller, Oort \\& Raimond (1963) the HVCs have been observed, surveyed, and catalogued in ever increasing detail (Bajaja et al. 1987; Hulsbosch \\& Wakker 1988; Wakker \\& van Woerden 1991; Hartmann \\& Burton 1997; Braun \\& Burton 2000; Br\\\"uns et al. 2000; Burton, Braun \\& Chengalur 2001; Putman et al. 2002; Lockman et al. 2002; de Heij, Braun \\& Burton 2002a). However, the nature and origin of the HVCs have remained the subject of considerable and unresolved debate (Oort 1966; Wakker \\& van Woerden 1997; Blitz et al. 1999; Wakker, van Woerden \\& Gibson 1999; Gibson et al.~2002; Sembach 2002).   The HVCs consist of several hundred compact, kinematically distinct, clouds with typical angular diameters of $\\sim 1-2^\\circ$ (Braun \\& Burton 2000; Putman et al. 2002), as well as several large (diffuse) and kinematically continuous complexes (e.g., the ``A, B, C, H, \\& M'' clouds) that extend over $\\sim 100$ square degrees. The HVCs appear to be devoid of any stellar counterparts. The radial velocities of the HVCs identified in the Wakker \\& van Woerden (1991) catalogue range from $-$464 to +297 km s$^{-1}$ relative to the local standard of rest. de Heij, Braun \\& Burton (2002b) identified 67 distinct compact high velocity clouds (CHVCs) in the single-dish ``Leiden/Dwingeloo'' HI survey (Hartmann \\& Burton 1997). Putman et al. (2002) identified an additional 179 CHVCs in the (more sensitive) southern-hemisphere ``HI Parkes All Sky Survey'' (HIPASS). Most of the CHVCs are infalling. The ``Leiden/Dwingeloo'' survey, carried out with angular and velocity resolutions of $0\\fdg5$ and 1 km s$^{-1}$ respectively, shows that the typical (one-dimensional) velocity dispersions of the CHVCs are $\\sim 14$ km s$^{-1}$ with little scatter about this value (Blitz et al. 1999). The velocity dispersions imply gas temperatures $\\sim 10^4$ K, consistent with a warm medium (WM) possibly consisting of a mixture of a warm neutral (WNM) and ionized (WIM) gas. Typically, N$_{\\rm HI}\\sim 5\\times 10^{18}$ cm$^{-2}$ averaged over the extents of the CHVCs. However, the CHVCs are barely resolved with the (25-meter) Dwingeloo and (64-meter) Parkes telescopes. More recently Burton et al. (2001) were able to resolve the spatial structures of ten CHVCs in higher resolution ($3.5 \\times 3.7$ arcmin) observations with the (305-meter) Arecibo telescope. They found that the typical observed (projected) $1/e$ exponential scale-lengths of the HI distributions equal 690 arcsec, with central HI columns ranging from $2\\times 10^{19}$ to $2\\times 10^{20}$ cm$^{-2}$. In some objects, gas was detected out to $\\sim 1^\\circ$ from the cloud centers down to column densities of $\\sim 2\\times 10^{17}$ cm$^{-2}$. The line widths are slightly narrower in the Arecibo data set compared to the Leiden/Dwingeloo observations, with corresponding velocity dispersions of $\\sim 11$ km s$^{-1}$. The velocity gradients across the clouds are relatively small, $\\sim 10$ km s$^{-1}$ degree$^{-1}$, and the WNM cloud dynamics are controlled primarily by thermal motions.  In another important development Braun \\& Burton (2000) and de Heij, Braun \\& Burton (2002c) carried out high resolution Westerbork interferometric observations (with $\\sim 1$ arcmin synthesized beams) and found that at least some CHVCs contain high column density (up to $\\sim 10^{21}$ cm$^{-2}$) cores, with typical radii $\\sim 10$ arcmin, and line profiles as narrow as 2 km s$^{-1}$. The narrow line widths imply HI gas in a cold neutral medium (CNM; $T\\lesssim 100$ K). Previous to this work CNM cores were known to exist only in the extended A, H, and M complexes (Wakker \\& Schwarz 1991). The recent single-dish and interferometric observations reveal that ``core/envelope'' CNM/WNM structures may be a common feature of the CHVCs.  The distance is the critical unknown for the HVCs. The large complexes are plausibly nearby objects, and stellar absorption-line observations in the large M and A clouds (Danly, Albert \\& Kuntz 1993; van Woerden et al. 1999) constrain their distances to within 5 and 10 kpc respectively. However, no direct distance determinations are yet available for the large population of CHVCs.  Many theories for the origin of the HVCs have been proposed over the years (Wakker \\& van Woerden 1997). Oort (e.g.~1966; 1970) recognized that if the HVCs are gravitationally bound objects with masses dominated mainly by the observed HI gas they must be distant, with distances $d\\gtrsim \\Delta v^2f/\\pi Gm_{\\rm H}N_{\\rm HI}\\theta$, where $\\Delta v$ is the observed velocity dispersion, $N_{\\rm HI}$ is the mean HI column within angular radius $\\theta$, and $f$ is the fraction of the total dynamical mass present as HI gas. If the clouds are bound, then the required value of $f$ is proportional to the assumed distance $d$, since the HI mass varies as $d^2$, whereas the dynamical mass $M_{\\rm dyn}\\equiv \\Delta v^2d\\theta/G \\propto d$. For $f=1$ the observations imply implausibly large ``binding distances'' $d\\gtrsim 10$ Mpc, and HI masses $M_{\\rm HI}\\equiv \\pi d^2\\theta^2m_{\\rm H} N_{\\rm HI}\\ \\gtrsim 10^9 M_\\odot$. Oort's preferred interpretation was that the HVCs are transient ``circumgalactic'' ($\\lesssim 100$ kpc) primordial gas clouds that have recently condensed out of the intergalactic medium, and are only now being accreted onto the disk in the end-stages of the Galaxy formation process. Shapiro \\& Field (1976) introduced the idea of the Galactic fountain (see also Bregman 1996) in which hot gas ejected from the disk cools and condenses into neutral clouds that fall back onto the Galaxy in a continuing cycle. Giovanelli (1981) suggested that the HVCs are scattered fragments of the Magellanic Stream.\\footnote{ The ``Magellanic Stream\" (Mathewson, Cleary \\& Murray 1974) is an example of extended HVC gas that is relatively well understood. It consists of a narrow trail of HI gas that emanates from the Magellanic Clouds, and stretches over 100$^\\circ$ across the sky.  The stream is almost certainly gas that has been tidally (Gardiner \\& Noguchi 1996) or ram-pressure stripped (Moore \\& Davis 1994) from the Magellanic Clouds as these satellite galaxies orbit the Milky Way.} A Local Group\\footnote{ We recall that the LG is a dynamically bound system of galaxies consisting of the Milky Way and its two largest neighbors M31 and M33, and 40 or more smaller and fainter dwarf galaxies (Mateo 1998). The LG is typical of the many small groups and clusters of galaxies situated along the outer boundary of the Virgo supercluster (van den Bergh 1999). The LG radius defined as the zero-velocity surface relative to the Hubble flow is $\\sim 1.5$ Mpc, and the characteristic velocity dispersion is $\\sim 60$ km s$^{-1}$ (Sandage 1986). The mass of the LG is dominated by the Milky Way and M31, and equals about $3.3\\times 10^{12} M_\\odot$.} (LG) origin for the HVCs was also proposed and discussed by various authors in the years immediately following their discovery (Oort 1966; Verschuur 1969; Einasto et al. 1974; Eichler 1976). In this picture the HVCs are genuine extragalactic objects associated with the LG system, and are at typical distances of $\\sim 1$ Mpc, rather than being a structural feature of the Galaxy and its immediate surroundings.  The distances to the clouds may be constrained via observations of optical H$\\alpha$ recombination line emission that provides a measure of the ionizing radiation field incident on the clouds (Weiner \\& Williams 1996; Bland-Hawthorn et al. 1998; Tufte et al. 1998; Weiner, Vogel \\& Williams 2001). At sufficiently large distances from the Galaxy, the dominant radiation field becomes the weak metagalactic background and the clouds should be weak H$\\alpha$ emitters. Nearby clouds might be ionized by more intense ``leakage radiation'' from the Galactic disk, and could be more intense H$\\alpha$ sources. Sensitive optical Fabry-Perot H$\\alpha$ observations have been reported and compiled by Weiner et al. (2001) for several HVC complexes and one isolated CHVC. Tufte et al. (2002) have reported H$\\alpha$ observations in five CHVCs. For the complexes Weiner et al. find that the measured H$\\alpha$ surface brightnesses range from 41 to 1680 milli-Rayleighs \\footnote{ 1 Rayleigh = $10^6/4\\pi$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$. For case-B recombination at a temperature of 10$^4$ K, a surface brightness of 1 Rayleigh is produced by a slab with emission measure $EM = 2.8$ cm$^{-6}$ pc, which for two-sided illumination in ionization equilibrium, is produced by a Lyman continuum photon intensity $J^* = 3.6\\times 10^5$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$.} (mR), implying photoionizing Lyman continuum (Lyc) intensities of 1.5$\\times 10^4$ to 6.0$\\times 10^5$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$. Weiner et al. detected an H$\\alpha$ intensity of 48 mR in CHVC 230+61+165 (an object included in the Burton et al. [2001] Arecibo data set). Tufte et al. (2002) report H$\\alpha$ intensities in the range of 20 to 140 mR in four CHVCs, and an upper limit of 20 mR in an additional CHVC. The detected fluxes are larger than the (2$\\sigma$) upper limit of 20 mR set by Vogel et al.~(1995) for the H$\\alpha$ intensity of the isolated extragalactic Giovanelli \\& Haynes (1989) cloud that Vogel et al. argue is probing the metagalactic field (see also Madsen et al. 2001; Weymann et al. 2001). Tufte et al. argue that the relatively large CHVC H$\\alpha$ fluxes imply that the CHVCs are located within the Galactic halo. However, the distance estimates are uncertain due to uncertainties in the actual strength of the metagalactic field, the Galactic radiation escape fraction, and the role of collisional ionization.  The metallicity and dust content are additional potentially important clues to the origin of the HVCs. Thermal ($\\sim 100$ $\\mu$m) emission signatures of dust grains have been searched for but not detected (Wakker \\& Boulanger 1986). However, a few metallicity measurements are available, mainly via ultraviolet absorption line studies toward background quasars and Seyfert galaxies. Significantly subsolar (though not purely primordial) metal abundances are indicated, primarily via observations of gas phase sulfur \\footnote{ Sulfur is generally considered to be robust probe of the intrinsic metallicity since S remains undepleted in diffuse gas (Jenkins et al. 1987).}. Low metallicities would appear to be inconsistent with a Galactic fountain, and favor models in which the gas originates outside the Galaxy. Lu et al. (1998) observed HVC 287+22+240 using the Goddard High Resolution Spectrograph (GHRS) on board {\\it Hubble Space Telescope} (HST) and detected several S{\\small II} and Fe{\\small II} absorption lines. Lu et al. derived S/H=$0.25\\pm 0.07$, Fe/H=$0.033\\pm 0.006$, and S/Fe=$7.6\\pm 2.2$ (relative to solar abundances). They interpreted the supersolar S/Fe ratio as being due to significant iron depletion and therefore evidence for the presence of dust grains. In view of similar abundances observed in the Magellanic Clouds, as well as the proximity of 287+22+240 to these objects, Lu et al. concluded that this HVC is a ``leading arm'' component of the Magellanic stream. Wakker et al. (1999) carried out GHRS spectroscopy of HVC gas in complex C and inferred S/H=$0.089\\pm 0.024$. They favor the Oort hypothesis, and interpret complex C as infalling primordial gas that has been ``contaminated'' by metals as it interacts with the Galactic halo. Braun \\& Burton (2000) derived a similar low metallicity for one of the CHVCs in their sample (125+41-207).  They combined their HI observations together with Mg{\\small II} absorption measurements reported by Bowen \\& Blades (1993), and inferred a Mg/H abundance in the range of 0.04 to 0.07 relative to solar.  Recently, Murphy et al. (2000) and Sembach et al. (2000) carried out {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) absorption-line studies of several HVCs. In complex C Murphy et al. found Fe/H $\\sim 0.5$ (relative to solar) that may indicate undepleted iron and hence a very low dust content. Sembach et al.~(2002) set a metallicity limit of $O/H < 0.46$ in CHVC 224-83-197 via the detection of O{\\small I} absorption. An important result of the FUSE observations was the detection of O{\\small VI} absorption tracing highly ionized gas in several HVCs. These observations complement earlier GHRS observations of C{\\small IV}, Si{\\small IV} and N{\\small V} absorption in high velocity gas that Sembach et al. (1999) suggested is produced in the ionized envelopes of neutral hydrogen HVCs (as the ionized clouds do not appear to be 21 cm sources). Sembach et al. (1999) argued that the observed C$^{+3}$, Si$^{+3}$ and N$^{+4}$ abundances are consistent with low pressure ($P/k \\sim 2$ cm$^{-3}$ K) gas photoionized by an assumed Lyc flux at a level plausibly consistent with the metagalactic background. They concluded that the C{\\small IV}, Si{\\small IV} and N{\\small V} absorbers are located at large distances and are Local Group objects. The O{\\small VI} abundances are too large to be compatible with photoionization, and may be produced collisionally at the interfaces between the HVCs and a tenuous Galactic corona or Local Group medium through which the they are moving (Sembach et al. 2000). We note that collisional ionization (in gas with pressures considerably larger than 2 cm$^{-3}$ K) may also contribute to the production of the C$^{+3}$, Si$^{+3}$ and N$^{+4}$ ions.  Evidence of an ambient medium in which the CHVCs may be embedded could also provide clues to their origin. Br\\\"uns, Kerp \\& Pagels (2001) mapped one of the CHVCs (125+41-207) in the Burton et al. sample with the 100m Effelsberg telescope. They found that the WNM component is asymmetric with a ``cometary'' appearance. The head-tail morphology suggests that gas is being stripped off of the main body of the cloud as it moves through an enveloping medium.  The Local Group hypothesis for the HVCs, particularly the CHVCs, has received renewed interest recently for additional observational and theoretical reasons.  Bajaja et al. (1987) noted that asymmetries in the ``position-velocity'' diagram of the HVC ensemble could be understood if the HVCs are moving within the LG as a whole.  More recently Blitz et al. (1999) and Braun \\& Burton (1999) demonstrated that the dispersion in CHVC velocities is minimized if the cloud velocities are measured relative to the LG barycenter (see however, Gibson et al. 2000). Furthermore, Blitz et al. presented a simulation of the LG dynamics that shows that ``test particles'' interacting with the Milky Way, M31, and the distant Virgo cluster, are (during a Hubble time) drawn into an elongated filamentary structure that reproduces the observed distribution and velocities of HVCs in the directions of the LG barycenter and anti-barycenter (see however Moore \\& Putman 2001; de Heij et al. 2002a). These kinematic considerations are strong evidence for a LG origin, although the cloud distances are not very well constrained. Blitz et al. concluded that the typical distances of the CHVCs are $d \\gtrsim 750$ kpc. For $d=750$ kpc, a characteristic radius of $\\sim 2.5$ kpc and a central gas density of $\\sim 6\\times 10^{-3}$ cm$^{-3}$ are implied by the Burton et al. (2001) observations. The HI mass in each cloud is $\\sim 1\\times 10^7 M_\\odot$, and the total HI mass in the system of CHVCs is $\\gtrsim 10^9 M_\\odot$. Such large masses may be in conflict with the apparent lack of similar objects around other galaxies and groups (Zwaan \\& Briggs 2000; Rosenberg \\& Schneider 2002). If the CHVCs are long-lived, gravitationally bound objects as conjectured by Blitz et al., then the measured velocity dispersions imply total dynamical masses $M_{\\rm dyn}\\approx  5\\times 10^8 M_\\odot$, comparable to dwarf galaxy masses.   In this picture the HI mass is a fraction $f=0.02$ of the total dynamical mass, and the 3.4 kpc radius corresponds to the ``scale height'' $\\sim GM_{\\rm dyn}/\\Delta v^2$ of the gravitationally confined gas.  Blitz et al. suggested that the dynamical mass is dominated by dark-matter and that the CHVCs are tracing individual dark-matter ``halos''.  The notion that the CHVCs represent dark-matter halos is motivated by the cosmological theory of hierarchical structure formation. In this theory, bound systems (``halos'') of non-baryonic, dark matter evolve in a hierarchical collapse, initiated by small gravitationally unstable primordial density fluctuations (Gunn \\& Gott 1972; Press \\& Schechter 1974; Blumenthal et al. 1984; Navarro, Frenk \\& White 1997; Bullock et al. 2001). Low-mass halos virialize first, and then merge into progressively more massive systems. Galaxies form as gas accumulates, cools, and collapses inside the virialized DM halos (White \\& Rees 1978). Detailed numerical simulations of the evolving collapse of dissipationless cold (non-relativistic) dark matter (CDM) have been carried out by several investigators (e.g.~Navarro et al. 1997; hereafter NFW). The simulations predict the existence of many more low-mass halos than observed at the faint-end (i.e. dwarf galaxy scale) of the galaxy luminosity function. Various suggestions have been made to account for this discrepancy, including: blow out of gas by winds and supernovae produced in small initial starbursts (Dekel \\& Silk 1986; Tegmark, Silk \\& Evrard 1993); photoionization heating of the baryonic component by the metagalactic radiation field, that inhibits the collapse of gas into the low-mass halos (Efstathiou 1992; Thoul \\& Weinberg 1996; Kepner, Babul \\& Spergel 1997; Kitayama \\& Ikeuchi 2000; Bullock Kravtsov \\& Weinberg 2000) and expulsion of gas trapped within the low-mass halos during the epoch of ``reionization'' (Barkana \\& Loeb 1999).  In all of these scenarios the low-mass halos (or ``mini-halos'' as coined by Rees 1986) continue to exist, but because star-formation has been inhibited in them, they are not observable as galaxies.  In two recent studies, Klypin et al. (1999) and Moore et al. (1999) presented high-resolution simulations of the evolution of dark matter halos at the mass and size scales of the Milky Way and M31 galaxies that dominate the LG.  They found that many low-mass sub-halos survive and continue to accrete into each of the dominating galaxy halos to the present day. In particular, the predicted number of low-mass satellite halos appears to exceed the number of observed satellite dwarf galaxies within $\\sim 300$ kpc of the Milky Way.  Klypin et al. suggested that the CHVCs are in fact the ``missing DM satellites'' that appear in their simulations. In this picture the CHVCs are ``circumgalactic'' mini-halos associated with the Milky Way, as opposed to ``extragalactic'' objects associated with the LG as proposed by Blitz et al. The presence of DM substructure may be detectable via the gravitational lensing of background quasars (Mao \\& Schneider 1998; Metcalf \\& Madau 2001; Dalal \\& Kochanek 2002).  The CNM/WNM structures observed in the CHVCs are in many ways similar to the multi-phased HI distributions observed in gas-rich dwarf irregular (dIrrs) galaxies \\footnote{ Blitz \\& Robishaw (2000) have recently shown that some dwarf spheroidals (dSphs) may also be gas-rich, with the gas being present at large offsets from the optical stellar light.}. As summarized by Mateo et al. (1998), the observed HI masses in the dIrrs range from 0.1 to 0.5 of the total galaxy mass (including the dark matter). The HI gas is generally much more extended than the ``Holmberg radii\" (0.1$-$1 kpc) of the stellar components as traced in the optical. Young \\& Lo (1996; 1997) carried out detailed {\\it Very Large Array} (VLA) interferometric studies of the dIrrs Leo A and Sag DIG and found that the observed HI line profiles can be separated into broad and narrow components. The broad (WNM) components appear as ``envelopes'' throughout the mapped regions of the galaxies, whereas the narrow (CNM) components are associated with small clumps or cores located near star-forming regions. The WNM envelopes appear to be pressure supported as opposed to rotationally supported.  In these respects the HI structures in the dwarfs are similar to those observed in the CHVCs. However, there are important differences (in addition to the fact that star forming regions are absent in the CHVCs). First, the peak HI column densities of the WNM gas in the dwarfs (up to a few 10$^{21}$ cm$^{-2}$) are significantly larger than in the CHVCs (up to $\\sim 10^{20}$ cm$^{-2}$). Second, the angular scale sizes of the HI distributions are smaller in the dwarfs, $\\sim 70$ arcsec in Sag DIG, and $\\sim 160$ in Leo A, compared to the CHVCs where the characteristic size is $\\sim 690$ arcsec. These facts will be important in our analysis of the CHVCs and dwarf galaxies as possibly related objects.  Another key difference is that the distances to the dwarfs are well determined (via stellar photometry) whereas the distances to the CHVCs are unknown.  In this paper we quantitatively examine the hypothesis that the CHVCs are stable HI clouds confined by the gravitational potentials of dark-matter halos.  We wish to determine whether viable ``mini-halo'' models for the CHVCs can be constructed as either ``circumgalactic'' ($d\\sim 100$ kpc) or ``extragalactic'' ($d\\sim 1$ Mpc) LG objects. Such clouds will be exposed to the background metagalactic radiation field that ionizes and heats the gas. The dark-matter confined clouds may also be subjected to bounding pressures, provided for example, by a hot circumgalactic Galactic corona, or by an intergalactic medium filling the Local Group.  We address several questions. Given the observed HI column densities and distributions what are the halo virial masses and characteristic dark matter densities required to confine the gas clouds? Are the required halo properties consistent with those expected from cosmological structure formation simulations? Under the assumption of photoionization by the metagalactic field, what are the minimum gas masses and column densities required for the formation of neutral hydrogen cores within the dark-matter halos? Are the observed HI distributions consistent with photoionization by the metagalactic field?  Under what conditions do the cloud cores become multi-phased? Can halo models be constructed for the pressure supported HI clouds in the Local Group dwarf galaxies? If the CHVCs are the starless ``cousins'' of the dwarf galaxies, and the CHVC and dwarf galaxy halos are similar, what are the implied distances to the CHVCs?  We construct general purpose models in which we compute the gas density distributions and ionization structures of pressure supported hydrostatic hydrogen gas clouds that are trapped in the potential wells of DM halos, are simultaneously heated and photoionized by an external radiation field, and are subjected to specified bounding pressures. Computations along these lines have been presented in the literature (with varying degrees of sophistication) mainly in the context of the mini-halo model for the intergalactic Ly$\\alpha$ clouds (e.g.~Rees 1986; Murakami \\& Ikeuchi 1990; Miralda-Escude \\& Rees 1993; Kepner et al. 1997). Of most relevance here are the computations presented by Kepner et al. (see also Corbelli \\& Salpeter 1993) who studied the phase transitions (ionized/neutral/molecular) of hydrogen gas clouds in DM halos. Their main focus was to show that low-mass halos form shielded neutral gas cores (a precondition for star-formation) only at late ($z\\lesssim 1$) cosmic times when the metagalactic photoionizing radiation field has become sufficiently diluted (Efstathiou 1992).  However, Kepner et al. made several simplifying assumptions that make it difficult to apply their models to either the CHVCs or the LG dwarf galaxies. These included assuming a fixed relation between the halo mass and the characteristic DM density (see \\S 2), adopting an arbitrary functional relationship between the total gas mass and halo mass, and the neglect of external pressure. In our analysis we relax all of these assumptions.  In addition, we determine the conditions required for the formation of thermally unstable multi-phased cores. To do this we employ and extend the methods presented by Wolfire et al. (1995a) in their study of the multi-phased Galactic disk, and calculate the thermal equilibrium properties of low metallicity HI gas heated by the metagalactic field. We then incorporate the results of these computations into our dark-matter halo models. We note that Wolfire et al. (1995b) (see also Ferrara \\& Field 1994) carried out an explicit analysis of the HVC complexes observed by Wakker \\& Schwarz (1991). In particular, the cloud distances were constrained by arguing that the CNM detected by Wakker \\& Schwarz is possible only when the ambient gas pressure, exceeds a well defined minimum $P_{\\rm min}$. In this analysis the gas pressure is provided by a Galactic X-ray corona. However, Wolfire et al. did not consider dark-matter dominated clouds. In such systems the central cloud pressures can become much larger than the ambient pressures. The transition from the WNM to the CNM phase then becomes possible in the cloud cores even when the ambient pressure is much smaller than $P_{\\rm min}$.  We first discuss in \\S 2 the basic parameters and properties of the dark-matter halos we consider in our study, including also a discussion of the relevant cosmological scaling relations. In \\S 3 we then describe our model clouds and discuss the methods we use to compute the cloud hydrostatics, radiative transfer, and HI phase structure.  In \\S 4 we discuss our representation of the metagalactic background field. In \\S 5 we derive several useful analytic formulae for the gas distributions in DM halos. In \\S 6 we present the results of our numerical computations. In \\S 7 we model the HI gas structures in the dwarf galaxies Leo A and Sag DIG. In \\S 8 we consider models of the CHVCs as either circumgalactic or extragalactic mini-halos. A summary is presented in \\S 9. ",
        "conclusions": "In this paper we examine the hypothesis that the compact high velocity HI clouds (CHVCs) trace sub-structure within the dark-matter halo of the Galaxy, or within the entire Local Group system.  For this purpose we carry out detailed computations of the coupled hydrostatic and ionization structures of pressure supported clouds that are confined by gravitationally dominant dark-matter (DM) ``mini-halos''. We focus on low-metallicity systems that are ionized and heated by the metagalactic background field. We provide a fit for this field (from the near-IR to the hard X-ray regime) that is based on a combination of observations and theoretical estimates. We explicitly consider the effects of external bounding pressures on the mini-halo cloud structure. We consider bounding pressures provided by a hot ionized Galactic corona or an intergalactic medium in the Local Group.  We consider dark-matter halos with either cuspy (NFW) or constant density (Burkert) cores. We adopt a $\\Lambda$CDM cosmological model to parameterize the expected correlations between the various halo scale parameters, as well as the dispersion in the correlations. We focus on low-mass halos with virial masses $\\sim 10^8$ to $10^{10}M_\\odot$, (or halo scale velocities between $\\sim 10$ to 50 km s$^{-1}$) containing warm ($10^4$ K) clouds of neutral plus photoionized gas with masses between $\\sim 10^4$ and 10$^8M_\\odot$.   We determine how the cloud sizes and hydrogen gas distributions, as well as the gas phase-states -- ionized (WIM), neutral (WNM), or multi-phased (WNM/CNM) --  depend on the halo parameters, gas masses, and bounding pressures. We consider both gravitationally confined systems, where the gas is effectively restricted to within the halo scale radii (or halo cores), as well as pressure-confined clouds in which the gas can extend to large halo radii. We find that gravitationally confined systems are those for which $v_s/c_g \\gtrsim 1.5$, where $v_s$ is the halo scale velocity and $c_g$ is the gas sound speed. We determine the conditions required for the transition from ionized to neutral clouds, and for the formation of thermally unstable multi-phased (WNM/CNM) cores within neutral clouds. We also compute the emission measures and H$\\alpha$ recombination line surface brightnesses that are produced in the WIM components of the halo clouds. For optically thick clouds ionized by the metagalactic field, we find that the H$\\alpha$ surface brightness ranges from 4 mR at the cloud centers to 7 mR at the limb-brightened edges.  As a step in our computations we construct pressure vs.~density phase diagrams for low metallicity ($Z=0.1 Z_\\odot$) gas that is heated by the metagalactic field for a wide range ($10^{18}-10^{21}$ cm$^{-2}$) of assumed HI shielding columns, and we compute the critical phase-transition pressures $P_{\\rm min}$ and $P_{\\rm max}$ as functions of the hydrogen ionization rate. We use the results of these computations to determine the thermal phase states of the neutral gas in our halo cloud models.  We present mini-halo models for the pressure supported HI structures observed in the Local Group dwarf irregular galaxies Leo A and Sag DIG. We include the stellar contribution to the gravitational potentials. We identify the HI cloud boundaries observed in Leo A and Sag DIG with the ionization fronts, and we derive an upper limit of $P_{\\rm HIM}/k \\lesssim 100$ cm$^{-3}$ K for the ambient pressure of the intergalactic medium in the Local Group. The distances to Leo A and Sag DIG are well established, and the observed HI gas scale heights of $0.5$ kpc in these objects imply characteristic DM densities of $1.2$ amu cm$^{-3}$ (or $0.03M_\\odot$ pc$^{-3}$) for the DM halos. These densities are consistent with those previously found via rotation curve studies of rotationally supported dwarfs and low-surface brightness galaxies. Leo A and Sag DIG obey the halo correlations that are expected for typical (``median'') DM halos in a $\\Lambda$CDM cosmology, provided the halos contain constant density cores.  NFW halos would have to be extremely underconcentrated (by $-4\\sigma$ relative to median halos) given the observed gas scale heights.  For the family of median Burkert (or $-4\\sigma$ NFW) halos that obey the scalings and correlations as fixed by the dwarf galaxies, we find that the transition from ionized to neutral clouds occurs for central WNM HI column densities equal to $3\\times 10^{18}$ cm$^{-2}$, independent of the halo virial mass or total gas mass. The gas is fully neutral within projected radii where HI column exceeds $10^{19}$ cm$^{-2}$. A photoionization cut-off is predicted near $10^{19}$ cm$^{-2}$ (similar to what is found for flattened, rotationally supported, disk systems) independent of the halo mass. We find that for clouds heated by the metagalactic field, multi-phased cores occur for peak WNM columns between $2\\times 10^{19}$ and $2\\times 10^{20}$ cm$^{-2}$. The peak WNM columns of $\\sim 10^{21}$ cm$^{-2}$ observed in the dwarf galaxies require additional (possibly internal) heat sources.  We construct explicit ``mini-halo'' models for the multi-phased (and low-metallicity) compact high-velocity HI clouds. We consider the CHVCs as either ``circumgalactic'' objects associated with the Milky Way halo, with implied characteristic distances of 150 kpc, or as more distant $\\gtrsim 750$ kpc ``extragalactic'' Local Group objects. If the CHVC halos are drawn from the same family of halos that successfully reproduce the dwarf galaxy observations, then the CHVCs must be circumgalactic objects. The observed $1/e$ gas scale heights then correspond to sizes of 0.5 kpc (as in the dwarfs). The typical HI mass for an individual CHVC is then $\\sim 3\\times 10^5M_\\odot$, and the total HI mass in the entire population of CHVCs is $\\sim 6\\times 10^7M_\\odot$. The predicted multi-phased behavior that occurs for peak WNM columns between $2\\times 10^{19}$ and $2\\times 10^{20}$ cm$^{-2}$ is consistent with the observed multi-phased behavior and range of peak WNM columns in the CHVCs. If the large population of CHVCs represent ``missing low-mass DM satellites'' of the Galaxy, then these HI clouds must be pressure-confined to keep the gas neutral within the weak DM potentials. The observed number of CHVCs imply typical CHVC mini-halo scale velocities of $v_s=12$ km s$^{-1}$. For such objects the confining pressure must exceed $\\sim 50$ cm$^{-3}$ K to keep the gas bound. We construct a simple model for a Galactic corona, and show that a hot $2\\times 10^6$ K Galactic corona (in pressure equilibrium with the Galactic disk) could provide the required pressure at 150 kpc.  If the CHVCs are ``extragalactic'' objects with distances $\\gtrsim 750$ kpc, then their associated halos must be very ``underconcentrated'', with characteristic DM densities $\\lesssim 0.08$ cm$^{-3}$, much smaller than expected for their mass, and significantly smaller than observed in the dwarf galaxies. Multi-Phased cores are possible, but require shielding columns that are generally higher than observed. For $d=750$ kpc, the typical CHVC HI mass is $\\sim 1\\times 10^7M_\\odot$, and the total HI mass in the system of CHVCs is $\\sim 2\\times 10^9M_\\odot$.  Our analysis favors the circumgalactic hypothesis for the location of the CHVCs.  In this picture, the CHVCs represent pressure-confined clouds that are associated with tidally truncated dark-matter sub-halos that have survived the hierarchical formation process of the Galactic halo.  The CHVC mini-halos are visible as neutral 21 cm sources due the compression provided by an ambient Galactic corona.  Our analysis also appears to favor Burkert halos with constant density dark-matter cores, as opposed to NFW halos with diverging central DM densities. While Burkert halos are consistent with the observed scale heights at circumgalactic distances of 150 kpc, NFW halos would have to be significantly underconcentrated to yield the observed cloud sizes for such distances.  We compute the maximum masses of WNM gas that can be maintained by the metagalactic field within the DM halos. Our analysis suggests that in the mini-halo scenario for the CHVCs, these objects are ``failed galaxies'' and did not form stars, because most of the neutral gas was always maintained as WNM by the metagalactic field.  In contrast the halos associated the dwarf galaxies in the Local Group (and elsewhere) contained sufficiently large masses of gas such that most of the neutral component was converted to CNM, a precondition for star-formation.  Our hydrostatic mini-halo cloud models are able to account for many properties of the CHVCs, including their observed peak HI columns, core sizes, and multi-phased behavior. However, important theoretical and observational difficulties remain. Theoretically, a question remains as to the origin of the gas in the CHVC mini-halos. The gas is unlikely to be purely primordial, since such gas is easily lost from low-mass mini-halos as the bounding pressure becomes low. The presence in some objects of extended low column density HI wings, and H$\\alpha$ emission line fluxes in several CHVCs that are significantly larger than expected, represent significant observational challenges to the mini-halo models we have presented in this paper. Additional high resolution HI mapping observations and sensitive H$\\alpha$ line measurements are required to determine the possible contributions of gas stripping and perhaps collisional ionization to the CHVC cloud structure.  Such observations will help establish, or possibly refute, the hypothesis that the CHVCs trace dark-matter substructure.           \\vspace{0.8cm}  We thank Leo Blitz, Stu Bowyer, James Bullock, Butler Burton, Marc Davis, Orly Gnat, David Hollenbach, Tsafrir Kollat, and David Spergel for many helpful discussions during the course of this work. We thank James Bullock for providing us with his CDM halo code. We also thank the referee for comments and suggestions that improved this paper. C.F.M. was supported in part by NSF grant AST-0098365. M.G.W. was supported in part by a NASA LTSA grant NAG5-9271, and by NSF grant AST 95-29167.  \\appendix"
    },
    "0207/astro-ph0207089_arXiv.txt": {
        "abstract": "We describe \\texttt{Brute}, a code for generating synthetic spectra from 3D supernova models.  It is based on a Monte Carlo implementation of the Sobolev approximation.  The current version is highly parameterized, but in the future it will be extended for full NLTE spectrum modelling of supernovae. ",
        "introduction": "The amount of evidence that the envelopes of supernovae (SNe) can deviate from spherical symmetry is increasing.  Net intrinsic polarization measurements from some SNe are consistent with ellipsoidal envelopes (Howell et al.\\ 2001; Leonard et al.\\ 2001; Leonard et al.\\ 2000; Wang et al.\\ 2001; Wang et al.\\ 1997; see Wheeler 2000 for a list of measurements).  In some flux spectra, a clumpy or macroscopically mixed ejecta distribution (or at least a nonspherical excitation structure) has been invoked to explain phenomena such as the ``Bochum event'' in SN 1987A (Phillips \\& Heathcote 1989; Utrobin, Chugai \\& Andronova 1995).  Most conspicuously, the interesting morphologies of SN remnants may indicate macroscopic mixing in the initial SN envelopes (Decourchelle et al.\\ 2001; Fesen \\& Gunderson 1996; Hwang, Holt \\& Petre 2000; Tsunemi, Miyama \\& Aschenbach 1995).  These findings have inspired new 3D explosion models (Khokhlov 2000; Reinecke, Hillebrandt \\& Niemeyer 2002; Kifonidis et al.\\ 2000). Testing new models as they arise will require detailed NLTE radiative transfer calculations in 3D.  Nonetheless, considering and deducing general constraints on SN geometry without such detailed calculations can be fruitful (Thomas et al.\\ 2002).  We have developed a parameterized synthetic spectrum code for 3D models of SNe called \\texttt{Brute} which we describe here.  In \\S 2 we describe the assumptions and implementation of \\texttt{Brute}, and in \\S 3 we outline future work to bring the code toward a detailed analysis code. ",
        "conclusions": ""
    },
    "0207/astro-ph0207510_arXiv.txt": {
        "abstract": "We present long-slit spectra of three irregular galaxies from which we determine the stellar kinematics in two of the galaxies (NGC 1156 and NGC 4449) and ionized-gas kinematics in all three (including NGC 2366). We compare this to the optical morphology and to the HI kinematics of the galaxies. In the ionized gas, we see a linear velocity gradient in all three galaxies.  In NGC 1156 we also detect a weak linear velocity gradient in the stars of (5$\\pm$1/$\\sin i$) \\kms\\ kpc$^{-1}$ to a radius of 1.6 kpc. The stars and gas are rotating about the same axis, but this is different from the major axis of the stellar bar which dominates the optical light of the galaxy.  In NGC 4449 we do not detect organized rotation of the stars and place an upper limit of (3/$\\sin i$) \\kms\\ kpc$^{-1}$ to a radius of 1.2 kpc. For NGC 4449, which has signs of a past interaction with another galaxy, we develop a model to fit the observed kinematics of the stars and gas. In this model the stellar component is in a rotating disk seen nearly face-on while the gas is in a tilted disk with orbits whose planes precess in the gravitational potential. This model reproduces the apparent counter-rotation of the inner gas of the galaxy. The peculiar orbits of the gas are presumed due to acquisition of gas in the past interaction. ",
        "introduction": "We have learned that irregular galaxies are not just small versions of spiral disks, but we do not know the intrinsic characteristics of these systems, the most common type of galaxy in the universe. Indeed, the fundamental structure of irregular galaxies is still in doubt. It has been generally believed that irregular galaxies are disks like spirals, and like spirals, surface brightness profiles in most irregulars are most often adequately fit with an exponential law (Patterson \\& Thuan 1996; van Zee 2000). However, studies of the distributions of projected minor-to-major axis ratios $b/a$ suggest that irregulars are thicker than spirals, perhaps having intrinsic flattening ratios $(b/a)_0$ of 0.3--0.4 rather than the 0.2 value generally adopted for spirals (Hodge \\& Hitchcock 1966, van den Bergh 1988, Binggeli \\& Popescu 1995). Staveley-Smith, Davies, \\& Kinman (1992), on the other hand, argue that the low ratio of rotation velocity to velocity dispersion often seen in irregulars must imply a thick disk with an intrinsic flattening ratio as high as 0.6. More recently Sung \\et\\ (1998) have also analyzed ellipticity distributions and concluded that dwarf irregulars and Blue Compact Dwarfs (BCDs) are more triaxial than disk-shaped, having axis ratios 1:0.7:0.5 and being only a little less spherical than dwarf ellipticals. What then is the true shape of irregulars?  Because of these outstanding issues, we have begun a program aimed at understanding the intrinsic structure of normal irregular galaxies through determining their stellar kinematics. In this paper we present results from long-slit spectral observations of stellar absorption features in two irregular galaxies of high surface brightness. While observations of giant galaxies are often detailed enough to justify sophisticated analysis of their structure, (see, for example, Statler \\& Fry 1994; Statler, Smecker-Hane, \\& Cecil 1996; Statler \\& Smecker-Hane 1999), observations of the stellar kinematics of irregular galaxies have been harder to obtain (but see Swaters 1999). Most irregular galaxies are much lower in surface brightness than giant galaxies, and they rotate much more slowly, so a higher velocity resolution is needed. Thus, the signal-to-noise that is possible with current instrumentation limits what can be deduced for irregular galaxies. Nevertheless, we can begin to learn about the highest surface brightness irregulars now, and expect to extend to lower surface brightness objects as long-slit spectrographs or field integrators with moderate spectral resolution become available on very large telescopes.  An understanding of the shape of Im galaxies is important to understanding other aspects of these systems. A colleague of ours saw a broad-band optical picture of the irregular galaxy Sextans A (Hunter 1997) on a bulletin board and asked why the galaxy appears square and has such sharp edges. Sextans A is not alone; other irregular galaxies also defy our expectations of what galaxies should look like. The rectangular shape of these galaxies is often taken to mean that the galaxies are barred. But, if that is the case, their bar structures are very different from those in spirals. In spirals the bar length is typically $\\leq$0.3 of the size of the stellar disk (Elmegreen \\& Elmegreen 1985). In the rectangular irregulars the bar is a significant fraction of the entire optical galaxy. For example, in NGC 4449 the bar has a length that is 90\\% of D$_{25}$, the diameter measured to a B-band surface brightness level of 25 magnitudes arcsec$^{-2}$ (Hunter, van Woerden, \\& Gallagher 1999). What does it mean for most of the galaxy to be a bar?  There are also often peculiarities in the gas kinematics and distributions that are not evident in the optical morphology. In Sextans A the velocity field of the neutral gas is not aligned with the optical axis of symmetry and there is some evidence that gas is moving on elliptical orbits (Skillman \\et\\ 1988). In NGC 4449 the inner gas is counter-rotating with respect to the outer gas and neither gas system is aligned with the isophotes of the optical galaxy (Hunter \\et\\ 1999). However, because of its large cross-section, the gas of many irregulars could have been perturbed by outside forces over the galaxy's lifetime (for example, NGC 4449: Hunter \\et\\ 1998). What are the kinematics of the stellar system, which has a smaller cross-section, and how do they compare to that of the gas?     In this paper we report results of our first observations of stellar kinematics determined from absorption spectra in the irregulars NGC 1156 and NGC 4449. We also discuss the kinematics of the ionized gas in these two galaxies as well as in NGC 2366. ",
        "conclusions": "In NGC 2366 we have measured the rotation of the ionized gas to 114\\arcsec\\ (1.9 kpc). Rotation with a major axis PA of 30\\arcdeg, that of the optical galaxy, or of 45\\arcdeg, that of the HI kinematics, fit the observations equally well. The rotation gradient is linear (can be fit with a straight line) with a gradient of ($30\\pm2/\\sin i$) \\kms\\ kpc$^{-1}$. Peculiar velocities southwest of the galaxy center, where there are two supergiant HII regions, mask the rotation pattern there.  In NGC 1156 the ionized gas shows a rotation gradient that can be fit with a straight line to 42\\arcsec\\ (1.6 kpc). The major axis is 84\\arcdeg---129\\arcdeg. The HI kinematics have a PA of order 130\\arcdeg\\ in the region the \\ha\\ is being measured, but a PA of order 83\\arcdeg\\ overall (Swaters 1999). However, the optical major axis has a PA of 39\\arcdeg, so the morphological and kinematic axes are misaligned. The gradient of the ionized gas is ($13\\pm1/\\sin i$) \\kms\\ kpc$^{-1}$. Rotation of the stars is weak, but linear to 43\\arcsec\\ (1.6 kpc). The most probable kinematic major axis of the stars is also 84\\arcdeg. The gradient is ($5\\pm1/\\sin i$) \\kms\\ kpc$^{-1}$.  In NGC 4449 we found that the ionized gas rotates in a manner consistent with a major axis of 46\\arcdeg, the same as the morphological major axis of the inner galaxy. The gradient is ($12\\pm1/\\sin i$) \\kms\\ kpc$^{-1}$ to 110\\arcsec\\ (2.1 kpc). The stars, on the other hand, show no ordered rotation, with an upper limit of (3/$\\sin i$) \\kms\\ kpc$^{-1}$ to 65\\arcsec\\ (1.2 kpc).  In all three galaxies we have found that the ionized and neutral gas velocities are similar, but often with important differences. First, the rise of the rotation speeds with radius is steeper in the ionized gas than in the HI. This is easily understood as due to the large difference in resolution of the two data sets. The HI observations typically were made with a beam-size of order 30\\arcsec\\ while the optical spectra used a slit of 3\\arcsec. This beam-smearing effect can make the rotation curve deduced from the HI shallower (see, for example, Bosma 1981; Rubin \\et\\ 1989; Blais-Ouellette \\et\\ 1999; Swaters, Madore, \\& Trewhella 2000).  Second, the major axes inferred from the kinematics of the different galactic components are not always the same, nor are they the same as the major axis of the optical morphology. The various PAs are collected in Table \\ref{tabpa}. In NGC 1156 the various gas and stellar kinematical axes are most likely the same, but this axis differs from the morphological major axis. In NGC 2366 the morphological and HI kinematical axes are different, and the ionized gas is equally consistent with either of these axes. In NGC 4449 the kinematical major axis of the ionized gas is coincident with the morphological major axis of the inner optical galaxy but differs from the HI kinematic axis.  In all three galaxies we can identify regions in which the ionized gas rotation velocity gradients can be approximated by a straight line. However, in many regions the true rotation is disguised by large superposed non-circular motions. The velocity gradient of the stars in NGC 1156 is also linear over the region of the galaxy that we measured rotation velocities. The gradients in the stars are lower than those in the ionized gas in both galaxies."
    },
    "0207/astro-ph0207456_arXiv.txt": {
        "abstract": "The prompt emission of gamma-ray bursts probably comes from a highly relativistic wind which converts part of its kinetic energy into radiation via the formation of shocks within the wind itself. Such \"internal shocks\" can occur if the wind is generated with a highly non uniform distribution of the Lorentz factor. We estimate the expected photospheric emission of such a relativistic wind when it becomes transparent. We compare this thermal emission (temporal profile + spectrum) to the non-thermal emission produced by the internal shocks. In most cases, we predict a rather bright thermal emission that should already have been detected. This favors acceleration mechanisms for the wind where the initial energy input is under magnetic rather than thermal form. Such scenarios can produce thermal X-ray precursors comparable to those observed by GINGA and WATCH/GRANAT. ",
        "introduction": "The cosmological origin of long duration gamma-ray bursts (hereafter GRBs) has been firmly established since the discovery of their optical counterparts in 1997 \\citep{vanparadijs:97}. These late and fading counterparts, the so called afterglows, have now been detected in many bursts, and in different spectral ranges\\,: X-rays, optical and radio bands. The redshift has been mesured for about 20 GRBs from $z=0.43$ to $z=4.5$. The corresponding isotropic equivalent energy radiated by these GRBs in the gamma-ray range goes from $5\\ 10^{51}\\ \\mathrm{erg}$ to $2\\ 10^{54}\\ \\mathrm{erg}$. The beaming factor that has to be taken into account to obtain the real amount of radiated energy can be deduced from afterglow observations (achromatic break in the lightcurve, \\citet{rhoads:99}). Current estimates lead to a total energy radiated in gamma-rays of about $0.5-1\\ 10^{51}\\ \\mathrm{erg}$ \\citep{frail:01}. The most discussed scenario to explain the GRB phenomenon is made of three steps :\\\\ \\noindent\\textbf{Central engine :} The source of GRBs must be able to release a very large amount of energy in a few seconds. The two most popular candidates are either the merger of compact objects (neutron star binaries or neutron star-- black hole systems \\citep{narayan:92,mochkovitch:93}) or the gravitational collapse of a massive star into a black hole (collapsars/hypernovae \\citep{woosley:93,paczynski:98}). Such events lead to the formation of very similar systems made of a stellar mass black hole surrounded by a thick torus. The collapsar model seems to be favored in the case of long bursts by observational evidences that GRBs are located well inside their host galaxy and often associated to star-forming regions \\citep{paczynski:98,djorgovski:01}. The released energy  is first injected into an optically thick wind, which is accelerated via an unknown mechanism, probably involving MHD processes \\citep{thompson:94,meszaros:97,spruit:01} and becomes eventually relativistic. The existence of such a relativistic wind has been directly inferred from the observations of radio scintillation in GRB 970508 \\citep{frail:97} and is also needed to solve the compactness problem and avoid photon-photon annihilation along the line of sight. Average Lorentz factors larger than 100 are required \\citep{baring:97,lithwick:01}. The next two steps explain how the kinetic energy of this relativistic wind is converted into radiation at large distances from the source, when the wind has become optically thin.\\\\ \\noindent\\textbf{Internal shocks :} the production of gamma-rays is usually associated to the formation of shocks within the wind itself \\citep{rees:94}. Such internal shocks can appear if the initial distribution of the Lorentz factor is highly variable, which is very likely considering the unsteady nature of the envisaged sources \\citep{macfadyen:99}. This model has been studied in details \\citep{kobayashi:97,daigne:98,daigne:00}. The main difficulties which are encountered are a rather low efficiency for the conversion of the wind kinetic energy into gamma-rays (a few percents only) and problems in reproducing with synchrotron emission the slope of the low energy part of the spectrum \\citep{ghisellini:00}. Despite this difficulty, the model can successfully reproduce the main features of the bursts observed by BATSE.\\\\ \\noindent\\textbf{External shock :} the relativistic wind is decelerated later by the external medium. This phase of deceleration is probably the best understood of the three steps and reproduces very well the afterglow properties \\citep{wijers:97}. The dynamics of the wind during the deceleration phase is described by the solution of the relativistic Sedov problem \\citep{blandford:76} and the observed afterglow is due to synchrotron emission produced by relativistic electrons accelerated behind the strong forward shock propagating in the external medium \\citep{sari:98}.\\\\  The work presented in this paper focuses on the prompt emission. The spectrum of this emission as observed by BATSE and Beppo-SAX is non-thermal and is well fitted by the 4-parameter ``GRB-function'' proposed by \\citet{band:93}. This function is made of two smoothly connected power-laws. This non-thermal emission probably originates from the radiation of a population of highly relativistic electrons accelerated behind the shock waves propagating within the wind during the internal shock phase.\\\\  Prior to the internal shock phase, the relativistic wind has to become transparent. At this transition, a thermal emission is produced, that could contribute to the observed prompt emission. Parts of the wind can also become opaque at larger radii if internal shocks create pairs in large number. These opaque regions can produce additional thermal components when they become transparent again \\citep{meszaros:00}. Other thermal contributions can be expected, for example when the jet breaks out at the boundary of the stellar envelope in the collapsar scenario \\citep{ramirez-ruiz:02}. In this paper, we restrict our analysis to the photospheric thermal component. A similar problem has been studied by \\citet{lyutikov:00} in the different context of strongly magnetized winds emitted by rapidly rotating pulsars.\\\\  The paper is organized as follows : in sect.~\\ref{sec:photosphere} we obtain the position of the photosphere of a relativistic wind with a highly variable initial distribution of the Lorentz factor, as expected in the internal shock model. We then compute the corresponding photospheric thermal emission in sect.~\\ref{sec:photosphericemission} and compare it to the non-thermal emission from the internal shocks in sect.~\\ref{sec:Comparison}. The results are discussed in sect.~\\ref{sec:Discussion} and the conclusions are summarized in sect.~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} In the framework of the internal shock model for gamma-ray bursts, we have computed in a detailed way the photospheric emission of an ultra-relativistic wind with a variable initial distribution of the Lorentz factor. We have compared the obtained spectrum and time profile to the non-thermal contribution of the internal shocks. Our main results are the following :\\\\ \\noindent \\textit{(1) The photosphere in the standard fireball model is too hot and luminous.} In the standard fireball model where the initial temperature of the fireball is about 1 MeV, the internal energy is still large when the wind becomes transparent and the photosphere is therefore hot and luminous. The consequence is that the photospheric thermal component in the X-ray/gamma-ray range is in most cases at least as bright as the non-thermal component due to the internal shocks (even if the internal shock efficiency is high). This is in contradiction with the observations of BATSE and Beppo-SAX showing non-thermal spectra.\\\\ \\noindent \\textit{(2) MHD winds are favored.} Results in much better agreement with the observations are obtained when it is assumed that only a small fraction $\\lambda$ of the energy released by the source is initially injected under internal energy form in a fireball. Most of the energy could for instance be initially under magnetic form, a large fraction of the Poynting flux being eventually converted into kinetic energy at large distances. For a typical internal shock efficiency of a few percents, values of $\\lambda \\la 0.01$ are required, which means that not more than 1\\% of the energy is initially deposited in the ejected matter (whose initial temperature is then of about a few hundreds keV).\\\\ \\noindent \\textit{(3) X-ray thermal precursors can be obtained.} A consequence of this strong assumption is that moderately low $\\lambda$ ($\\lambda \\simeq $ a few percents) lead to the presence of thermal X-ray precursors if the distribution of the Lorentz factor is not too variable in the initial phase of wind production. The characteristics of these precursors (spectral range, duration, intensity) are very comparable to the X-ray precursor activity observed in several GRBs by GINGA and WATCH/GRANAT.\\\\ \\noindent \\textit{(4) The optical photospheric emission is very weak.} For very small $\\lambda$ values, the photospheric emission can be shifted to even lower energies. However, we have shown that it also becomes much too weak to explain the prompt optical emission observed by ROTSE in GRB 990123.\\\\  A good test of the results presented in this paper would be the detection of X-ray precursors by an instrument with good spectral capabilities, so that a thermal origin could be firmly established. A determination of the photospheric temperature would put an interesting constraint on the $\\lambda/\\!\\!\\left.f_{\\gamma}\\right.$ ratio and then on the wind acceleration mechanism. Moreover, if the photospheric thermal emission could be clearly detected (for instance in the soft X-ray range), it would provide a direct information about the initial distribution of the Lorentz factor in the wind before the internal shocks start.\\\\ \\vspace*{-3ex}"
    },
    "0207/physics0207007_arXiv.txt": {
        "abstract": "A new vibrational level of the H$_2^+$ molecular ion with binding energy of $1.09\\times 10^{-9}$a.u.$\\approx30$~neV below the first dissociation limit is predicted, using highly accurate numerical nonrelativistic quantum calculations, which go beyond the Born-Oppenheimer approximation. It is the first excited vibrational level $v=1$ of the 2p$\\sigma_u$ electronic state, antisymmetric with respect to the exchange of the two protons, with orbital angular momentum $L=0.$ It manifests itself as a huge p-H scattering length of $a=750\\pm 5$ Bohr radii. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207383_arXiv.txt": {
        "abstract": "The bright X-ray binary X2127+119 in the core of the globular cluster M15 has long been thought to be in an unusual evolutionary state, in which the binary is embedded in a common envelope. Support for this idea comes from X2127+119's absorption lines, which are blue shifted at all orbital phases, indicating the existence of outflows from the system. A common-envelope scenario implies that the absorption lines should exhibit maximum blue shift near mid-eclipse (binary phase 0.0). We have re-analysed INT spectra of X2127+119 obtained in 1986, 1987 and 1988 using the latest orbital ephemeris (substantially different from that used in the original analysis), and find that the orbital phase at which the absorption lines show a maximum blue shift is not 0.0, but rather 0.25 -- 0.3. These results indicate that a common-envelope scenario for X2127+119 may not work. In addition, from spectrograms of the \\heii line, we report the first tentative detection of X2127+119's companion star. ",
        "introduction": "The X-ray binary X2127+119 (AC211) in the core of the globular cluster M15 is optically one of the brightest and most well-studied of the low-mass X-ray binaries (LMXBs), and yet it remains a highly enigmatic system. Its optical and X-ray light-curves, along with its very low $L_{\\rm X}$/$L_{\\rm opt}$, make it a classic accretion disc corona (ADC) source (Fabian, Guilbert \\& Callanan 1987; Naylor et al. 1988), in which the system is seen almost edge-on and the compact object and hot, luminous inner disc are obscured by the accretion disc rim. The X-ray flux we observe comes entirely from photon-scattering by a large corona above the disc, and is only a small fraction of the source's intrinsic X-ray flux.  The X-ray source X2127+119 has recently been discovered by {\\it Chandra} to be two separate sources (White \\& Angelini 2001). The second X-ray binary is 2.7 arcsec from AC211, is 4 magnitudes fainter in the U-band, but its {\\it Chandra} count-rate is 2.5 times higher. The discovery of a second LMXB in the core of M15 explains one of most puzzling aspects of AC211: luminous X-ray bursts showing expansion of the neutron star photosphere have been observed in X2127+119 (Dotani et al. 1990, van Paradijs et al. 1990, Smale 2001), which has been difficult to reconcile with AC211's optical and X-ray light curves which showed it to be an ADC -- to observe X-ray bursts with photospheric expansion the neutron star surface has to be visible, but in ADC sources the neutron star is hidden from view by the accretion disc. This problem goes away if it is the second LMXB, and not AC211, which is the burster. The {\\it Chandra} discovery may solve one mystery but it produces another: the fact that AC211's X-ray luminosity is even fainter than previously believed makes its unusually high optical luminosity even more puzzling, and suggests that the central X-ray source hidden from view must be exceptionally luminous, possibly indicating a very high mass-transfer rate from the companion star.  Determining whether AC211 has an extremely high mass-transfer rate and is in an unusual evolutionary state is important: it is relevant to the understanding of LMXB evolution and to the understanding of stellar interactions within, and the evolution of, globular clusters. Theoretical determinations of the numbers of neutron stars in globular clusters and the efficiency with which they interact with stars in the cluster cores to form binaries, combined with the large number of millisecond radio pulsars (end-products of LMXB evolution) observed in the clusters, imply that we see far fewer LMXBs in globular clusters than we should. This may imply that the evolution of LMXBs within globular clusters is accelerated, and that their lifetimes are very short (less than 10$^8$ years; Hut, Murphy \\& Verbunt 1991), which in turn implies high mass-transfer rates.  Not much is known about either the compact object or the donor star comprising AC211, although its location in the core of a globular cluster enables us to make reasonable assumptions about their nature. No globular cluster X-ray binary is known to contain a black hole, and neutron stars are abundant in globular clusters: theory predicts that the number of neutron stars in a dense globular cluster such as M15 is \\til 4000 (Hut, Murphy \\& Verbunt 1991). The chances therefore that the compact object in AC211 is a neutron star are high. The system has an orbital period of 17.1 hours, indicating that it is unlikely that the companion star is on the main sequence (at that orbital period a main sequence donor star would not fill its Roche lobe). The main sequence turn off point for M15 occurs at \\til 0.8 M$_{\\odot}$ (Fahlman et al. 1985), so we assume that AC211's donor star is an 0.8-M$_{\\odot}$ star which has begun to evolve off the main sequence. ",
        "conclusions": "The evolutionary state of X2127+119/AC211 is of great interest because of theoretical predictions that globular cluster LMXBs evolve much faster than their galactic cousins (Hut, Murphy \\& Verbunt 1991). That AC211 appeared to be embedded in a common envelope seemed to imply that at least one globular cluster LMXB did have an unusual evolutionary history. Our results, however, show that AC211 is unlikely to be a common envelope system because the \\he radial velocities appear to be inconsistent with an L$_2$ outflow. We may also have detected the irradiated face of the elusive companion star in the \\heii line profiles.  The Balmer and \\he lines indicate the presence of outflows from the system, but without better data it is difficult to determine their exact nature. There are several possible scenarios which may explain the nature of these outflows, but each requires rigorous modelling before it can be considered seriously as a mechanism for producing our observed \\he line profiles."
    },
    "0207/astro-ph0207660_arXiv.txt": {
        "abstract": "Infrared OH lines at 1.5 - 1.7 $\\mu$m in the $H$ band were obtained with the NIRSPEC high-resolution spectrograph at the 10m Keck Telescope for a sample of seven metal-poor stars. Detailed analyses have been carried out, based on optical high-resolution data obtained with the FEROS spectrograph at ESO. Stellar parameters  were derived by adopting infrared flux method effective temperatures, trigonometric and/or evolutionary gravities and metallicities from \\ion{Fe}{2} lines. We obtain that the sample stars with metallicities [Fe/H] $<$ $-$2.2 show a mean oxygen abundance [O/Fe] $\\approx$ 0.54, for a solar oxygen abundance of $\\epsilon$(O) = $\\approx$ 8.87, or [O/Fe] $\\approx$ 0.64 if $\\epsilon$(O) = 8.77 is assumed. ",
        "introduction": "Oxygen abundances in metal-poor stars are a key information for understanding the early phases of Galactic chemical evolution. Very few data on oxygen abundances are available for stars with metallicities around [Fe/H] $\\approx$ -3.0, and results obtained from different lines are often discrepant.  The infrared (IR) X$^2 \\Pi$ vibration-rotation transitions of OH lines in the $H$ band were first used in halo stars by Balachandran \\& Carney (1996) for the dwarf  HD 103095, having derived [Fe/H] = $-$1.22 and [O/Fe] = 0.29. Balachandran et al. (2001, 2002) and Mel\\'endez et al. (2001) presented new oxygen abundance determinations in metal-poor stars, from OH lines in the $H$ band.  Their [O/Fe] values tend to show agreement with those derived from [OI] lines in the metallicity range -1.0 $\\simless$ [Fe/H] $\\simless$ -2.5; three among  the four analyzed stars of lower metallicities give higher values.   The oxygen abundance determination using IR OH lines may solve the problem of disagreement between the values obtained from different oxygen abundance indicators. The IR OH lines in the H band of the first overtone have measurable intensities down to [Fe/H] $\\sim$ -3.0 for giants with effective temperatures T$_{\\rm eff}$ $\\sim$ 4500 K,  and [Fe/H] $\\sim$ -3.5 for T$_{\\rm eff}$ $\\sim$ 4300 K, whereas for dwarfs the lines are stronger, and with effective temperatures T$_{\\rm eff}$ $\\sim$ 4500 K, stars of [Fe/H] $\\sim$ -3.5 could be measured. Otherwise only the fundamental transition lines at 3.5 $\\mu$m, in the L band, are measurable. (The difficulty with these latter lines comes, however, from the thermal background and lower sensitivity of the instruments.) The relative intensity of the OH lines in the $L$ and $H$ bands can be seen in Hinkle et al. (1995) for Arcturus, and for the Sun in Livingston \\& Wallace (1991), in which only the $L$-band lines are seen.  IR vibration-rotation OH lines tend to form in LTE, whereas the electronic transition UV lines tend to form in non-LTE (Hinkle \\& Lambert 1975). Non-LTE effects in the formation of OH will affect both sets of lines, according to Asplund \\& Garc\\'{\\i}a-P\\'erez (2001, hereafter AGP01).   In the present work we use IR OH lines in the region 1.5 - 1.7 $\\mu$m in order to derive oxygen-to-iron ratios for a sample of seven metal-poor stars, most of them in the metallicity range -3.0 $<$ [Fe/H] $<$ -2.0.  In Sect. 2 the observations are presented. In Sect. 3 the detailed analyses are described. In Sect. 4 the results from IR OH lines are discussed, and in Sect. 5 conclusions are drawn. ",
        "conclusions": "We obtained high-resolution infrared spectra in the $H$ band in order to derive oxygen abundances from IR OH lines. In order to have a homogeneous set of stellar parameters, we carried out detailed analyses using equivalent widths of iron lines measured on high-resolution spectra from the FEROS spectrograph at ESO.  \\begin{enumerate} \\item {The sample stars with metallicities [Fe/H] $<$ $-$2.2 show a slight increase in [O/Fe], with respect to the results for higher metallicity stars. A mean of [O/Fe] = 0.5 is found, similar to the value derived from the forbidden [OI] lines.}  \\item{ A clear increase in the oxygen abundance derived with increasing equivalent width is found, such that the stronger lines give higher abundances. We adopt an oxygen abundance resulting from a mean of the weaker OH lines. It is also important to note that the effect is more pronounced for the most metal-poor stars.}  \\item{ It is clear that oxygen abundance determinations for the more metal-poor stars derived from very high S/N forbidden [\\ion{O}{1}] and IR OH lines are needed, as well as 3D model atmospheres calculations, in the case of IR OH lines.}  \\end{enumerate}"
    },
    "0207/astro-ph0207311_arXiv.txt": {
        "abstract": "We provide theoretical procedures and practical recipes to simulate non-Gaussian correlated, homogeneous random fields with prescribed marginal distributions and cross-correlation structure, either in a $N$-dimensional Cartesian space or on the celestial sphere. We illustrate our methods using far-infrared maps obtained with the Infrared Space Observatory. However, the methodology presented here can be used in other astrophysical applications that require modeling correlated features in sky maps, for example, the simulation of multifrequency sky maps where backgrounds, sources and noise are correlated and can be modeled by random fields. ",
        "introduction": "Random fields are widely used in astrophysics to model realistic scenarios of physical processes that depend on random components. For example, they are widely used in cosmology to model different types of galactic and extragalactic backgrounds \\citep[e.g.,][]{martinez}. The physical characteristics of the backgrounds are translated into statistical characteristics of the field. The fields are usually required to be ergodic so that information about the model can be extracted from a single realization of the field. In addition, they are assumed to be isotropic (the correlation between two points depends only on the distance that separates them) or homogeneous (the correlation depends on the difference of their position vectors) to reflect the geometry of the cosmological model. Given the complexity of the cosmological models, characteristics of the fields are usually studied via Monte Carlo simulations where model predictions are compared with observations.  To simulate a Gaussian random field we only need to specify the mean and spectrum (or correlation function) but there are non-Gaussian models whose predictions also have to be tested against observations. For example, standard inflationary models predict Gaussian temperature fluctuations of the cosmic microwave background but there are other cosmological models that predict non-Gaussian fluctuations \\citep{avelino,peebles}. Homogeneous non-Gaussian random fields are more difficult to define and simulate, and since they are not uniquely determined by their first two moments, we often have to accept only a partial second order description of the fields.   \\citet{vio01} (hereafter {\\it VAW}) presented numerical methods for the simulation of homogeneous scalar random fields $R(\\tb)$ with prescribed one-dimensional distribution function (marginal distribution) $ F_R(r)$ and correlation structure. These methods can be used, for example, to generate cosmic microwave background maps with a given spectrum and a marginal distribution that allows for asymmetry or kurtosis in the pixel temperatures. In this paper we generalize these methods to the case of pair-wise correlated random fields defined on the same physical space. This will allow us, for example, to simulate backgrounds and source fields that are not independent, as in star-forming regions where the already formed stars are linked to the surrounding gaseous and dusty environment from which they originate.  We consider random fields defined in $N$-dimensional ``parameter spaces'' where the coordinates $\\tb = (t_1,t_2,\\ldots,t_N)$ may correspond to spatial/angular coordinates (spatial random fields), time (time processes), a mix of these two (spatio-temporal random fields), or to even more general situations. We also consider the case of random fields defined on the sphere, that is, random fields that depend only on the direction in the sky. If the multivariate random field is $\\Rb(\\tb) = \\{ R_1(\\tb)$, $R_2(\\tb)$,$\\ldots$, $R_M(\\tb) \\}$, the goal is to generate a field with components $R_i$ having (possibly different) prescribed marginal distributions $F_{R_i}(r_i)$, prescribed correlation functions, and prescribed pair-wise cross-correlations.  In theory, a complete description of a random field requires the definition of all finite-dimensional joint distributions, but, unless the field is Gaussian, this is a terribly difficult task. In practical applications one only considers a second order description of the field that specifies the marginals $F_{R_i}(r_i)$, the means $\\mu_{R_i}={\\rm E}[\\,R_i(\\tb)\\,]$ and the cross-covariance functions \\begin{equation} \\label{eq:covariance} \\xi_{R_i R_j}(\\tb_1,\\tb_2)= {\\rm E} \\left[\\, (R_i(\\tb_1)-\\mu_{R_i}) ~(R_j(\\tb_2) -\\mu_{R_j})\\, \\right], \\end{equation} where ${\\rm E}[\\cdot]$ stands for expected value (ensemble average).  As mentioned before, it is often possible, even necessary, to adopt some simplifying assumptions like isotropy or homogeneity of the field $\\Rb(\\tb)$. In the multidimensional context, isotropy means that the cross-covariance function depends on the length $\\tau=\\Vert \\taub \\Vert$ of the vector $\\taub= \\tb_1-\\tb_2$ but not on its direction: $\\xi_{R_i R_j}(\\tb_1,\\tb_2)= \\xi_{R_i R_j}(\\Vert \\tb_1-\\tb_2 \\Vert)$. In this case the cross-correlation function (normalized covariance function) depends only on $\\tau$ \\begin{equation} \\label{eq:rhos} \\rhob_{\\Rb}(\\tau) = \\left( \\begin{array}{ccc} \\rho_{R_1 R_1}(\\tau) & &\\rho_{R_1 R_M}(\\tau) \\\\ \\vdots & \\ddots & \\vdots \\\\ \\rho_{R_M R_1}(\\tau) & &\\rho_{R_M R_M}(\\tau) \\end{array} \\right), \\end{equation} where \\begin{equation} \\label{eq:rho} \\rho_{R_i R_j}(\\tau) = {\\rm E} \\left[ \\frac{(R_i(\\tb) - \\mu_{R_i}) ~(R_j(\\tb + \\taub) - \\mu_{R_j})}{\\sigma_{R_i} \\sigma_{R_j}} \\right] \\end{equation} and $\\sigma_{R_i}^2$ is the variance corresponding to the marginal $F_{R_i}(r_i)$. By definition, $\\rho_{R_i R_i}(0) = 1$ for $i=1,2,\\ldots,M$.  Although the isotropic case is of great interest in astronomical applications, here we consider multidimensional homogeneous random fields in general, that is, random fields whose covariance function depends on $\\taub$. In this case, the correlation function $\\rhob_{\\Rb}(\\taub)$ is defined as in equation (\\ref{eq:rhos}) but with $\\taub$ instead of $\\tau$.  The rest of the paper is organized as follows. In Section \\ref{sec:notes} we describe a general procedure for generating non-Gaussian random fields by point-wise transformations of Gaussian ones. In Section \\ref{sec:practical} we provide practical recipes for the methods outlined in Section \\ref{sec:notes} and show how they can be used in either $N$-dimensional Cartesian spaces or on the celestial sphere. Section \\ref{sec:examples} presents some examples. We show simulated maps of the background and localized source field components of a far-IR sky map of the Infrared Space Observatory (ISO). In the ISO map, the locations and intensities of the sources are correlated with the background field and only a multidimensional simulation may reproduce such physical characteristic. ",
        "conclusions": "We have described spectral methods for the simulation of homogeneous non-Gaussian multivariate random fields with prescribed cross-correlation and marginal distributions. The basic idea is to apply pointwise transformations to homogeneous Gaussian fields to guarantee the homogeneity of the field. The transformations are chosen to obtain the desired marginal distributions and cross-correlation function. The proposed methodology works for a wide range of cross-correlation functions and marginal distributions of general interest. Among these, some Gaussian mixture models may be useful to model fields of localized sources.  We have illustrated the methods by simulating sky maps with background and source components that are physically linked. The same methods can also be applied to simulate and/or analyse more complex maps, such as those from all-sky surveys of the Planck mission (Mandolesi et al., 1998; Puget et al., 1998). The objective of the Planck Satellite is to observe the Cosmic Microwave Background anisotropies which account, however, for only a very small part of the much more conspicuous backgrounds, either Galactic and Extragalactic. Sky maps will be produced at nine different frequencies but strong correlations are expected among different maps because of the wide frequency distribution of the expected emissions. Many algorithms have been already tested to recover as much information as possible from the observed maps (see e.g. Maino et al., 2001 and references therein). But to extract important cosmological information the methods should provide information on all the contributing components and should be able to model the correlations among them. The methods we have presented here are based on theoretically simple assumptions but do require data, or other available physical information, to select the correct model parameters (cross-correlation function, marginal distributions, Gaussian mixture parameters, etc.). Appropriate modeling of more complex sky maps is the goal of future research."
    },
    "0207/astro-ph0207127_arXiv.txt": {
        "abstract": "We present a simple model of hot gas in galaxy clusters, assuming hydrostatic equilibrium and energy balance between radiative cooling and thermal conduction.  For five clusters, A1795, A1835, A2199, A2390 and RXJ1347.5-1145, the model gives a good description of the observed radial profiles of electron density and temperature, provided we take the thermal conductivity $\\kappa$ to be about 30\\% of the Spitzer conductivity.  Since the required $\\kappa$ is consistent with the recent theoretical estimate of Narayan \\& Medvedev (2001) for a turbulent magnetized plasma, we consider a conduction-based equilibrium model to be viable for these clusters.  We further show that the hot gas is thermally stable because of the presence of conduction.  For five other clusters, A2052, A2597, Hydra A, Ser 159-03 and 3C295, the model requires unphysically large values of $\\kappa$ to fit the data.  These clusters must have some additional source of heat, possibly an active galactic nucleus since all the clusters have strong radio galaxies at their centers.  We suggest that thermal conduction, though not dominant in these clusters, may nevertheless play a significant role by preventing the gas from becoming thermally unstable. ",
        "introduction": "For many years it was thought that radiative losses via X-ray emission in galaxy clusters leads to a substantial inflow of gas, and mass dropout, in the form of a ``cooling flow'' (see \\citealt{fabi94} for a review).  Mass deposition rates were estimated to be as much as several hundred $M_\\odot{\\rm yr^{-1}}$ in some clusters (e.g., \\citealt{pere98}).  However, there was little direct evidence for the mass dropout in any band other than X-rays \\citep{fabi94}.  Recent high resolution X-ray data from {\\it XMM-Newton} and {\\it Chandra} indicate that there is no evidence for the multi-temperature gas that one expects if there is substantial mass dropout (\\citealt{pete01, tamu01, bohr01, fabi01, mole01, mats02, john02}). The new data strongly suggest that mass dropout must be prevented by some additional source of heat that balances radiative losses.  Two possibilities are currently being investigated for the heat source: (i) energy in jets, outflows, and radiation from a central active galactic nucleus (\\citealt{pedl90, tabo93, chur00, chur02, ciot01, brue02}), and (ii) thermal energy from outer regions of the cluster transported to the central cooling gas by conduction (\\citealt{nara01}, hereafter NM01).  Conduction in clusters has been discussed by several authors over the years in various contexts \\citep{binn81, tuck83, bert86, breg88, gaet89, rosn89, davi92, pist96}, but its importance was always considered doubtful.  For conduction to have any significant effect on the cooling gas in a cluster, the effective isotropic coefficient of conduction $\\kappa$ has to be a reasonable fraction of the classical \\citet{spit62} conductivity $\\kappa_{Sp}$.  However, conventional wisdom is that magnetic fields strongly suppress conduction perpendicular to the field.  Therefore, while conduction may be very efficient parallel to the field, the overall isotropic $\\kappa$ is expected to be $\\ll\\kappa_{Sp}$.  This picture has changed with the recent work of NM01 who, following earlier work by \\citet{rech78}, \\citet{chan98}, \\citet{chan99} and \\citet{maly01}, showed that a turbulent MHD medium in which the fluctuation spectrum extends over two or more decades of spatial scales has an effective $\\kappa$ that is a substantial fraction of $\\kappa_{Sp}$.  This is because of chaotic transverse wandering of field lines as a result of strong MHD turbulence \\citep{gold95}, which leads to a large enhancement of cross-field diffusion.  NM01 estimated that the ratio $f=\\kappa/\\kappa_{Sp}$ in a turbulent MHD medium is approximately $0.2$, though, given the uncertainties in their model, the value could probably lie anywhere in the range $\\sim0.1-0.4$.  NM01 showed via a simple order-of-magnitude estimate that thermal conduction with $f\\sim0.2$ is sufficient to balance the radiative losses of the cooling gas in X-ray clusters.  Based on this, they suggested that conduction might be a large part of the explanation for the lack of mass dropout in X-ray clusters.  Their suggestion finds support in the work of \\citet{voig02} and \\citet{fabi02} who show that conduction with $f$ in the range $0.1-1$ can indeed be a dominant heat source in X-ray clusters except perhaps in the very central regions.  The aim of the present paper is to study conduction in galaxy clusters in more detail.  Our goals are two-fold.  First, we consider a very simple model in which the gas is in hydrostatic equilibrium and in thermal balance, with cooling exactly compensated by heat conduction. We investigate whether this basic model can fit the observed shapes of the electron density profile $n_e(r)$ and temperature profile $T(r)$ of clusters for which high-resolution observations are available.  We believe this approach, which is complementary to that of \\citet{voig02}, is a good test of the conduction hypothesis, and a reliable method of estimating what value of $f$ is needed in various clusters to explain the observations.  Second, having found a steady state model for a given cluster, we check whether the hot gas is thermally stable. We briefly discuss the importance and plausibility of other energy sources.  The paper is organized as follows.  In \\S2 we write down the relevant differential equations and boundary conditions for a steady equilibrium model of hot gas in a cluster.  In \\S3 we compare model predictions with data on 10 clusters, and in \\S4 we discuss the thermal stability of the gas in these clusters.  We conclude with a discussion in \\S5.  Throughout the paper, we take $H_0=70$ km sec$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$, rescaling observational results whenever the original papers have used a different cosmology for their analysis. ",
        "conclusions": "NM01 found that thermal conduction with $f\\equiv\\kappa/\\kappa_{Sp}$ of order a few tenths can explain the gross energetics of hot gas in galaxy clusters.  They showed that the heat flux transported to the cooler central regions of the cluster from further out is roughly consistent with what is needed to replace the energy lost through radiative cooling.  \\citet{gruz02} and \\citet{fabi02} confirmed this conclusion.  In this paper, we have carried out a further test of the conduction model by investigating the shapes of the density and temperature profiles, $n_e(r)$ and $T(r)$, in an equilibrium cluster.  We solve for these profiles self-consistently, as compared to previous authors, e.g., \\citet{bert86, breg88, voig02} and \\citet{fabi02}, who either used the observations directly or employed simple analytical expressions for the shapes of $n_e(r)$ and/or $T(r)$.  Among the 10 clusters that we have studied, 5 clusters, namely, A1795, A1835, A2199, A2390 and RXJ1347.5-1145, can be fitted well with a pure conduction model and with values of $f\\sim0.2-0.4$ (Table 2, Figs. 1, 2, 3).  Since the model involves no fitting parameters other than $f$, which is used primarily to set the normalization of $n_e$, and $r_c$, which is not really a fitting parameter but rather is given one of two values (\\S3), the good agreement between the model profiles of $n_e(r)$ and $T(r)$ and the data is very encouraging.  For these five clusters, the best-fit values of $f$ lie within the range expected on theoretical grounds (NM01).  We believe that these are strong arguments in support of the conduction model.  We should note, however, that the theoretical estimate $f\\sim0.1-0.4$ obtained by NM01 is based on a specific (plausible) model of the magnetic field topology, suggested by the work of Goldreich \\& Sridhar (1995, 1997), but the answer might change substantially if the field topology (which is poorly understood in clusters) is very different.  The other 5 clusters, namely A2052, A2597, Hydra A, Ser 159-03 and 3C295, require larger values of $f$, of order unity or greater, and we consider them to be inconsistent with a pure conduction model.  It is interesting that these five clusters all have active nuclei with strong radio activity, so that nuclear activity might well supply the additional heat required by our models.  It is unclear exactly how an AGN would heat the cluster medium.  Turbulent mixing of hot gas in a jet with the surrounding cooler gas is a possibility \\citep{chur02, brue02}, though the effect of magnetic fields on this mixing is poorly understood.  Heating of the gas by cosmic rays \\citep{loew91} or hard X-rays \\citep{ciot01} from the AGN are other possibilities.  In \\S4 we showed by means of a local linear stability analysis that the five clusters for which we have obtained good fits with a pure conduction model are all thermally stable.  This is an interesting result because, without conduction, cooling gas in clusters is generally thermally unstable.  We find it reassuring that a single mechanism, namely conduction, is able both to supply the heat lost via radiative cooling and to control the thermal instability in these clusters.  In the other five clusters, which require an additional source of energy such as a central AGN, the thermal stability of the gas is not assured.  One way of ensuring stability is to have a heat source with a strong non-linear dependence on density, e.g., $\\alpha>1.5$ in equation (\\ref{heating}); this appears somewhat unnatural.  Alternatively, conduction, though not an important energy source in these clusters, might still play a role in controlling the instability.  Based on the above results, we suggest that AGN heating and conduction both are important in clusters.  In some clusters, AGN heating dominates and conduction plays a secondary, though still important, role in helping to stabilize the system.  In other clusters, AGN activity is weak, and conduction supplies the energy for cooling as well as prevents the gas from becoming unstable.  While our work suggests that conduction is effective over much of the volume of a cluster, we note that {\\it Chandra} observations have revealed sharp temperature jumps in some clusters (e.g., \\citealt{mark00}; Vikhlinin, Markevitch \\& Murray 2001ab).  These cold fronts are associated with the interface between the intracluster medium and an infalling galaxy or group.  The observations clearly indicate that conduction is strongly suppressed across the fronts. \\citet{vikh01b} have argued that the magnetic field is stretched out parallel to the cold front by the motion of the infalling galaxy/group and that this explains the very low conductivity.  \\citet{mazz02} present interesting evidence for a possible Kelvin-Helmholtz instability in the cold front in A3667, confirming that the fronts are probably transient features with lifetimes of only a couple of dynamical times.  The presence of these fronts is thus unlikely to affect the strong conduction that we have hypothesized over the bulk of the cluster.  If cold fronts survive for longer than the cooling time of the gas, then the thermally isolated cool gas would experience runaway cooling.  Since such runaway cooling is ruled out by the absence of substantial mass dropout, the lifetimes of cold fronts cannot be longer than the cooling time.  \\bigskip The authors thank Bill Forman, Christine Jones Forman and Larry David for useful discussions, and the referee for helpful comments on the manuscript.  RN thanks the National Science Foundation for support under grant AST 9820686."
    },
    "0207/astro-ph0207257_arXiv.txt": {
        "abstract": "We examine the AAVSO light curve of U Geminorum from 1908 to 2002, with particular focus on the October 1985 outburst. This outburst was longer than any other seen in U Gem by about a factor of 2, and appears to be unique among all dwarf nova outbursts seen in systems with orbital periods longer than 3 hr in that one can measure the decay time scale during the initial slow decay. This rate is $\\sim26\\pm6$  d mag$^{-1}$. Using estimates of the rate of accretion during outburst taken from Froning et al., one can show that $\\sim10^{24}$ g of gas was accreted  onto the white dwarf during the outburst which constrains the surface density in the outer accretion disk to be $\\sim600$ g cm$^{-2}$. The observed time scale for decay is consistent with that expected in U Gem, given its orbital period and disk mass at the time the outburst began. The data are not of sufficient quality to be able to ascertain a deviation from  exponentiality in the decay light curve (as in the SU Ursa Majoris stars' superoutbursts). When coupled with the viscous time inferred from the (short orbital period) SU UMa stars, the U Gem viscous time scale lends support to the standard model in which the decays in dwarf novae can either be viscous or thermal, with the  ratio between them being roughly $h/r$ where $h$ is the vertical pressure scale height in the disk. Indeed, dwarf novae are the only systems for which one can be reasonably certain of the identification of ``viscous'' and ``thermal'' decays. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207582_arXiv.txt": {
        "abstract": "Atomic physics calculations of radiative cooling are used to calculate criteria for the overstability of radiating shocks. Our calculations explain the measurement of shock overstability by Grun et al. and explain why the overstability was not observed in other experiments. The methodology described here can be especially useful in astrophysical situations where the relevant properties leading to an overstability can be measured spectroscopically, but the effective adiabatic index is harder to determine. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207531_arXiv.txt": {
        "abstract": "\\vspace{1pc} The ANTARES deep-sea neutrino telescope will be located at a depth of 2400~m in the Mediterranean Sea. Deployment of the detector will commence this Autumn and is expected to be completed by the end of 2004. With a surface area of the order of 0.1~km$^2$ it will be one of the largest European detectors. The aim of neutrino telescopes is to detect high-energy neutrinos from astrophysical sources whilst also providing information on the early Universe. Successful operation of ANTARES in a deep sea environment constitutes an important milestone towards the ultimate goal of the construction of an underwater neutrino telescope at the cubic-kilometre scale. The sky coverage of astrophysical sources offered by a Mediterranean neutrino telescope is complementary to any similar device at the South Pole. The current status of the project is discussed and the expected performance of the detector is described in the context of the scientific programme of the project which comprises astrophysical studies, dark matter searches and neutrino oscillations. ",
        "introduction": "A promising challenge for exploring the Universe is the detection of high-energy ($\\gtrsim 1$~GeV) neutrinos. Such neutrinos could be produced by powerful cosmic accelerators, for example those in supernova remnants, active galactic nuclei, compact binaries, such as micro-quasars, and those producing gamma-ray bursts. Neutrino astronomy complements high-energy gamma astronomy since the early Universe cannot be probed with high-energy photons due to photon-matter and photon-photon interactions - gamma rays of a few hundred TeV from the Galactic Centre cannot survive their journey to the Earth. The weakly interacting nature of neutrinos, coupled with the fact that they point to their source of origin without deviation makes them unique `probes' with which to investigate regions at distances larger than 50~Mpc.  The production of neutrinos in cosmic accelerators requires hadronic mechanisms to be active in the source. Whilst most of the observed high energy emissions to date can be explained by electromagnetic processes, recent observations of gamma-rays up to 5~TeV from SNR RXJ1713-39 by the Cherenkov telescope CANGAROO \\cite{CANGAROO} suggest that electromagnetic mechanisms are inconsistent with the data but $\\pi^0$ decay models can readily explain the measured energy spectrum. There is, however, some debate on this result \\cite{Pohl}.  The technique employed by neutrino telescopes is dictated by the small neutrino cross section and the large background due to atmospheric muons. Natural Cherenkov radiators, such as water or ice, provide a large active volume at reasonable cost and the indirect detection of neutrinos through muons produced in charged current interactions increase the 'target' region by the muon range. Detectors are constructed at large depths where the atmospheric muon flux is significantly reduced compared to that at the surface. Upward-going muons produced by neutrinos having crossed the Earth, are recognised as the products of neutrino interactions in or close to the instrumented region.  The Cherenkov light emitted by charged particles in deep water or ice is detected using an array of photomultiplier tubes (PMTs) which are housed, together with some associated electronic components, in a high pressure-resistant glass sphere known as an optical module (OM). The muon direction and energy are measured using the arrival times and amplitudes of the PMT pulses. The detector sensitivity increases with energy due to the increase in the $\\nu$-N cross section, the longer muon range and the increase in the amount of emitted Cherenkov light through secondary particles. ",
        "conclusions": "The ANTARES neutrino telescope is foreseen to be fully deployed by the end of 2004. The R\\&D phase of the project, which started in 1996, has finished. During this phase a detailed assessment of the main requirements for an undersea neutrino telescope was made. The phase culminated in 1999 with the successful operation for several months of a demonstrator string which allowed the reconstruction of atmospheric muons. The project has now entered the construction and deployment phase for a 0.1~km$^2$ scale detector. An electro-optical cable has been successfully deployed and a prototype string and a string dedicated to environmental parameter measurements will be deployed during this Autumn.  Simulations tools are providing relevant information concerning the scientific programme of the experiment, namely neutrino astrophysics, dark matter and neutrino oscillations. The predicted ANTARES sensitivity to fluxes from cosmic accelerators are about an order of magnitude better than results presented to date by other neutrino telescopes.  \\begin{figure}[htb] \\vspace{-.5cm} \\includegraphics[width=19pc, height=15pc]{figure9.eps} \\vspace{-1 cm} \\caption{Distribution of $E_{\\mu}/\\cos\\theta$ which would be obtained by ANTARES after 4 years of data taking (data points) assuming $\\Delta m^2 = 3.2 \\cdot 10^{-3}$ eV$^2$ and maximal mixing. Errors are statistical only. The histogram is obtained assuming no oscillations.} \\label{fig10} \\vspace{-.5cm} \\end{figure}"
    },
    "0207/astro-ph0207641_arXiv.txt": {
        "abstract": "High-resolution, 2-D hydrodynamical simulations with a large dynamic range are performed to study the turbulent nature of the interstellar medium (ISM) in galactic disks. The simulations are global, where the self-gravity of the ISM, realistic radiative cooling, and galactic rotation are taken into account. In the analysis undertaken here, feedback processes from stellar energy source are omitted. We find that the velocity field of the disk in a non-linear phase shows a steady power-law energy spectrum over three-orders of magnitude in wave number. This implies that the random velocity field can be modeled as fully-developed, stationary turbulence.  Gravitational and thermal instabilities under the influence of galactic rotation contribute to form the turbulent velocity field. The Toomre effective $Q$ value, in the non-linear phase, ranges over a wide range, and gravitationally stable and unstable regions are distributed patchily in the disk. These results suggest that large-scale galactic rotation coupled with the self-gravity of the gas can be the ultimate energy sources that maintain the turbulence in the local ISM. We find that our models of turbulent rotating disks are consistent with the velocity dispersion of an extended HI disk in the dwarf galaxy, NGC 2915, where there is no prominent active star formation.  Numerical simulations show that the stellar bar in NGC 2915 enhances the velocity dispersion, and it also drives spiral arms as observed in the HI disk. ",
        "introduction": "Observational and theoretical studies have suggested that there are supersonic hydrodynamical or magneto-hydrodynamical turbulent motions in molecular clouds and also in various phases of the ISM \\citep{mccray79,larson81,quiro83,balbus91,gold95,gold97,franco99}. A number of numerical studies showed, however, that the turbulence in the clouds decays as $t^{-\\eta}$, with $ \\eta \\sim 1$ \\citep{mac98, ostriker01}. The dissipation time of the turbulence is of the order of the flow crossing time or smaller, even in the presence of strong magnetic fields \\citep{stone98}.  These numerical experiments suggest that energy input is necessary to maintain the turbulent motion in the molecular clouds.  Some energy sources originating in stellar activity have been proposed: stellar winds from young stars \\citep{norman80}, photoionization \\citep{mckee89}, and supernova explosions \\citep{MO}.  However, these processes can not be the main energy sources for keeping the turbulent motion in molecular clouds that do not harbor stellar activities \\citep{will98}.  In this case, external shocks caused by distant stellar energy release may produce vortices in the clouds \\citep{korn00}.  In differentially rotating disks with a magnetic field, the magneto-rotational instability \\citep{balbus91} is a probable cause for the turbulent motion. \\citet{sell99} have studied the MHD turbulence in extended HI disks, applied their model to the extended HI disk of NGC 1058, and have showed that the observed uniform velocity dispersion may be modeled by the MHD-driven turbulence.  In terms of the origin of the non-self-gravitating, pure hydrodynamical turbulence, on the other hand, Balbus, Hawley and Stone (1996) have claimed, on the basis of theoretical and numerical studies, that Keplerian disks are linearly and non-linearly stable. They rejected self-sustained hydrodynamical turbulence with outward angular momentum transport, because the turbulence cannot gain energy from the differential rotation. However, \\citet{kato97} suggested a steady hydrodynamical turbulence in Keplerian disks, which are caused by the pressure-strain tensor. Richard \\& Zahn (1999) have re-analiezed laboratory experiments of the Couette-Taylor flow, and concluded that turbulence may be sustained by differential rotation when $d\\Omega/dR < 0$. \\cite{godon99} showed that purely hydrodynamical perturbations can develop initially into either sheared disturbances or coherent vortices (see also Papaloizou \\& Pringle 1985). However, the perturbations decay and do not evolve into a self-sustained turbulence.  Another important mechanism that generates turbulence in a rotating disk is self-gravitational, local or global instability of the gas \\citep{gold65a,gold65b}. Although there are some simulations of self-gravitating, turbulent gas in a local shearing box (e.g. V$\\acute{\\rm{a}}$zquez-Semadeni et al. 1995; Gammie 2001), the local approximation would not be adequate to study the nature of self-gravity dominated turbulence \\citep{balbus99}. The global hydrodynamical simulations given by  \\citet{LKA2} have revealed spiral unstable modes in the differentially rotating disk, but the development of self-sustained turbulence was not observed.  In the this paper, we study the gravity-driven turbulence in galactic disks using two-dimensional, global hydrodynamical simulations.  We solve the basic hydrodynamical equations and the Poisson equation numerically, taking into account realistic radiative cooling.  We use the same numerical technique discussed by Wada \\& Norman (1999,2001) (hereafter WN99, WN01), in which kpc-scale dynamics of the multi-phase ISM can be followed with a sub-pc resolution. WN99 and WN01 suggested that a gravitationally unstable and thermally unstable disk can evolve into a globally quasi-stable disk with quasi-stationary turbulent velocity field. Here we perform simulations with a higher spatial resolution, and make a detailed analysis of the turbulent structure.  In \\S 2, we briefly summarize the numerical method and models, and in \\S 3, we show that turbulent energy spectra are achieved over a wide dynamic range (from kpc to pc) in the direct numerical experiments.  In order to study the effects of pure turbulence in the disk, we here ignore the energy feedback due to supernovae and stellar winds from massive stars \\citep{NF96}. In \\S 4, we discuss what is required for realistic models of the ISM, and we apply our model to the HI disk of the dwarf galaxy, NGC 2915. Extended HI disks are particularly relevant objects to compare with our models, because they are not significantly affected by the star formation \\citep{dickey90,meurer96}. We summarize our results in \\S 5. ",
        "conclusions": "\\subsection{Towards more realistic modeling of the ISM} Modeling the ISM on a galactic scale is a challenging problem in numerical astrophysics. It is required that realistic simulations include many elementary processes, such as various radiative cooling/heating processes, thermal conduction, interaction between the magnetic field and the ISM, self-gravity of the gas, energy feedback from supernovae, etc. Moreover, the simulations should be ultimately three-dimensional and global, that is the whole galactic disk should be solved. In this sense, our model and any other past numerical simulations of the ISM are still highly idealized (see a review by V$\\acute{\\rm{a}}$zquez-Semadeni 2002). Our model is global, and realistic radiative cooling and self-gravity of the gas are taken into account, but the magnetic field, which is important for the evolution of the interstellar turbulence (see references in \\S 1), is ignored. Therefore we should be careful to apply the present results to the real ISM. On the other hand, most MHD simulations of the ISM adopt the local shearing box approximation. There are many local 2-D models, e.g. self-gravitating ISM with radiative cooling (Vazquez-Semadeni et al. 1996), and  self-gravitating ISM with isothermal or adiabatic equation of state \\citep{kim01,gammie01}. Some models in 3-D consider radiative cooling and supernova feedback \\citep{korpi99}, but most of them are local, non-selfgravitating, and an isothermal equation of state (EOS) is assumed \\citep{mac99}. There was some global, 3-D MHD simulations with polytropic EOS calculated for the galactic central region, but they ignored self-gravity of the gas (e.g. Machida, Hayashi, \\& Matsumoto 2000). Using self-gravitating, global hydrodynamical simulations with radiative cooling, on the other hand, Wada (2001) showed that the ISM in the galactic central region has the complicated filamentary, multi-phase structure as seen in the 2-D simulations presented here. The results in the present paper and the past attempts are complementary, and they should be improved  with more consistent numerical modeling of the ISM in the future.  Finally we would like to comment on plausibility of the 2-D approximation for the ISM. The real ISM in a galactic disk should behave as three-dimensional turbulence below the scale height of the cold gas ($\\sim 100$ pc). Since the turbulence is driven by local gravitational instability, we expect that the ISM in 3-D behaves like 2-D turbulence on larger scales, therefore the inertial range of $d{\\rm ln} E(k)/d{\\rm ln} k \\sim -1$ would also appear for $k \\lesssim 40$ in 3-D ($k \\lesssim 200$ in 2-D). For $k \\gtrsim 40$, the energy spectrum would be Kolmogorov-like, i.e. $E(k) \\propto k^{-5/3}$, or if it is shock-dominated, $E(k) \\propto k^{-2}$ would be expected. In a real ISM on a smaller scale, the main energy sources would be magneto-hydrodynamical instability and/or energy feedback from stars as well as the energy decay from larger scales. Transition from 2-D to 3-D turbulence and its effect on the energy spectra are now important problems that can be addressed in studies of the gas dynamics in galactic disks, This would be clarified utilizing three-dimensional, global simulations of the ISM with magnetic field on a galactic scale.  One might wonder if the spatial resolution in our simulation is meaninglessly fine (i.e. $\\sim 0.5$ pc), if we apply the model to the real ISM whose scale height is $\\sim 100$ pc. However, as shown in Fig. 4, the transition from the inertia range to the `dissipative' part is affected by the numerical resolution. The maximum wave length of the inertia range is about 10 times larger than the grid size. Therefore even for 2-D simulations, the spatial resolution should be much finer than the assumed scale heigh of the disk.  \\subsection{Origin of Velocity Dispersion in the HI Disk of NGC 2915} The HI disk in the dwarf galaxy, NGC 2915, extends to over five times the Holmberg radius, and there is no active star formation observed outside the central region. Therefore, this galaxy is relevant to our study of the dynamics of the turbulent gas disk without the influence of the stellar energy feedback. Meurer, Mackie, \\& Carignan (1994) and Meurer et al. (1996) observed NGC 2915 ($D = 5.3$ Mpc, $M_B = -15.9$ mag, $R_{\\rm Ho}=2.93$ kpc) using the ATCA (Australia Telescope Compact Array) with a linear resolution is 640 pc  and the AAT (Anglo-Astrallia Telescope).  The total mass of the HI disk is $9.6\\times 10^8 M_\\sun$ and $R_{\\rm HI} = 14.9$ kpc. The integrated star-formation rate is 0.05 $M_\\sun$ yr$^{-1}$, and most of the star-formation is near the center of the galaxy.  The HI intensity map shows that there are spiral arms extending well beyond the optical extent.  The isovelocity contours show that the disk is rotating with a small amount of non-circular motion.  The line-of-sight velocity dispersion of the HI is $\\sim 20-40$ km s$^{-1}$ inside the optical radius, and it is $\\sim 8$ km s$^{-1}$ in the extended HI disk. The central question is what causes this large velocity dispersion in the extended HI disk?  We apply the same numerical method in \\S 2 to model the HI disk in NGC 2915 to obtain the radial distribution of the velocity dispersion in a non-star forming disk. We scale the potential model (eq.(5)) and the gas disk to represent the HI disk of NGC 2915. Two rotation curves are assumed for an axisymmetric component of the external potential: 1) the potential derived from the HI rotation curve (Model A in Fig. 16 of Meurer et al. 1996), where the core radius $a$ in eq.(5) is 4 kpc, and 2) models with a smaller core radius ($a = 2$ kpc).  The latter model was chosen, because the HI rotation curves derived from tracing peak intensities of position-velocity maps do not necessarily represent {\\it true} mass distribution of the galaxies, especially for the central part \\citep{sofue99}.  Sofue et al. (1999) claimed that the true central rotation curves tend to be steeper than those implied from the peak intensities in PV maps.  NGC 2915 has a central optical bar.  The observed spiral patterns could be resonance-driven structures by this central bar. The velocity dispersion in the disk is generally enhanced by the non-axisymmetric potential and its resonances. The pattern speed of the bar can be directly measured \\citep{tre84}, and it is $\\Omega_p = 8.0\\pm 2.4$ km s$^{-1}$ kpc$^{-1}$ \\citep{bureau99}. In order to investigate the effects of the central bar, we also performed runs with a non-axisymmetric potential (bar-potential) with the same pattern speed as observed (Bureau et al. 1999), and where the length of the bar is taken to be equal to the core radius $a$.  The non-axisymmetric part of the potential is assumed to be in the form $ \\Phi_1 (R, \\phi, t) = \\varepsilon (R) \\Phi_0 \\cos 2(\\phi -\\Omega_{\\rm p} t) \\ , $ where $\\varepsilon (R) $ is given as $ \\varepsilon(R)= \\varepsilon _0 {aR^2}/{(R^2+a^2)^{3/2}}.  $ The parameter $\\varepsilon_0$ represents the strength of the bar \\citep{wad01c}. The gas is initially distributed in an axisymmetric, 15 kpc radius disk with a exponential-like surface density profile to resemble the observed HI distribution (Fig. 15 in Meurer et al.  1996). We have run models with parameters as $6 \\le \\Omega_p \\le 11 $ km s$^{-1}$ kpc$^{-1}$, $0.05 \\le \\epsilon_0 \\le 0.15$, and $a = 2$ and 4 kpc.  The velocity dispersions of an axisymmetric model $A$ ($a=4$ kpc) at $t=5.4$ Gyr are plotted as a function of the radius in Fig. 9a. The velocity field is sampled at every 20 grid points (i.e. 292 pc) for $x$ or $y$-directions, and they are averaged in the same size as the observed beam size to calculate the dispersion. We find that the dispersion is about a factor 4 smaller than the observed value (20-40 km s$^{-1}$) at $R < 2$ kpc, a factor 2 smaller than the observed 8 km s$^{-1}$ at $R > 2$ kpc.  The velocity dispersions of a bar model $B$ ($a= 2$ kpc, $\\varepsilon_0 = 0.1, \\Omega_p = 8$ km s$^{-1}$ kpc$^{-1}$) is plotted in Fig. 9b.  The central velocity dispersion in this model is $\\sim 20-30$ km s$^{-1}$, and the it is about 3--4 km s$^{-1}$ in the outer disk.  The central velocity dispersion is comparable with the observed value, but the dispersion in the outer disk is still factor 2 smaller than observations.  Fig. 10 shows $E(k)$ for the two models $A$ and $B$. Both spectra show double power-law shapes, $E(k) \\propto k^{-1}$ and $k^{-1.5}$ in $ 10 \\lesssim k \\lesssim 100$ for models $A$ and $B$, respectively. The slope of the axisymmetric model is comparable to that in Fig. 5. The steeper slope in the bar model would be due to stronger shocks caused by the bar. The spectra also shows that the bar model has a factor of 2-10 larger energy than the axisymmetric model in the most scales.  In Fig. 11 and 12, we plot the density and PV maps of the bar model $B$.  Two major spiral arms, which are formed from many spirals and clumps, are evident as well as an inhomogeneous compact disk.  The PV diagram shows a large velocity dispersion at $R < 3$ kpc, which corresponds to the clumpy nuclear disk.   The Toomre's $Q$ parameter estimated by the observations is 5--9 \\citep{meurer96}. However, this does {\\it not}  mean that there are no local gravitationally unstable regions in the HI disk. The beam size in the observations is much larger than the typical size of the unstable regions. Therefore, we can observe the global stability of the HI disk, but probably the local instability is observed as the velocity dispersion in the beam size.  After exploring the parameter space ($\\epsilon_0,\\, \\Omega_p$, and $a$), we conclude that a set of parameters, i.e. $\\epsilon_0 \\sim 0.1, \\Omega_p \\sim 8 $ km s$^{-1}$ kpc$^{-1}$, and $a\\sim 2$ kpc, reasonably reproduces the observed radial distribution of the velocity dispersion, and the HI morphology. However, the observed value of the velocity dispersion in the extended disk is about factor two larger than that in the best-fit model $B$. This might be due to the limitation of our model (\\S 4.1), or because of the following reasons. Fig. 10 suggests that stronger bars can enhance the velocity dispersion in the disk by about factor of two. If the observed velocity dispersion is a result of {\\it hidden} gravitational instability in the HI disk, the dispersion could be larger in a more massive disk. That is, the observed HI mass does not necessarily represent the total mass of the gas in the galaxy.  There might be a component of high density, compact molecular clouds in the extended HI disk, which could be observed by the next-generation radio interferometer, ALMA."
    },
    "0207/astro-ph0207194_arXiv.txt": {
        "abstract": "The galactic kinematics of Mira variables derived from radial velocities, Hipparcos proper motions and an infrared period-luminosity relation are reviewed. Local Miras in the 145-200day period range show a large asymmetric drift and a high net outward motion in the Galaxy. Interpretations of this phenomenon are considered and (following Feast and Whitelock 2000) it is suggested that they are outlying members of the bulge-bar population and indicate that this bar extends beyond the solar circle. ",
        "introduction": "The galactic kinematics of Mira variables have for a long while been of importance in helping us understand both the nature and evolution of Miras as well as the structure of our own Galaxy. This paper is primarily concerned with the second point - what do we learn from Miras about galactic structure? The paper also concentrates on local kinematics, leaving out a detailed discussion of the kinematics of the galactic bulge, much work on which has of course been done here in Japan. It has been possible to take a fresh look at the local kinematics of Miras using Hipparcos astrometry, extensive new infrared photometry and published radial velocities. This was done in a series of three papers (Whitelock, Marang and Feast (2000), paper I: Whitelock and Feast (2000), paper II: Feast and Whitelock (2000a), paper III). The present paper summarizes some of the relevant results from these papers and extends the discussion of the kinematics. ",
        "conclusions": ""
    },
    "0207/astro-ph0207477_arXiv.txt": {
        "abstract": "We present some results from a systematic survey for disks around spectroscopically identified young brown dwarfs and very low mass stars. We find that $\\approx$75\\% of our sample show intrinsic IR excesses, indicative of circum(sub)stellar disks.  The observed excesses are well-correlated with H$\\alpha$ emission, consistent with a common disk accretion origin.  Because the excesses are modest, conventional analyses using only IR colors would have missed most of the sources with disks.  In the same star-forming regions, we find that disks around brown dwarfs and T~Tauri stars are contemporaneous; assuming coevality, this demonstrates that substellar disks are at least as long-lived as stellar disks.  Altogether, the frequency and properties of circumstellar disks are similar from the stellar regime down to the substellar and planetary-mass regime. This offers compelling evidence of a common origin for most stars and brown dwarfs. ",
        "introduction": "While the number of known brown dwarfs is growing rapidly, the origin of these objects is an unanswered question.  One insight into the formation mechanism(s) for brown dwarfs is whether young substellar objects possess circumstellar disks.  There is abundant observational evidence and theoretical expectation for accretion disks around young solar-type stars.  Thus, the presence of disks around young brown dwarfs would be naturally accommodated in ``star-like'' formation scenarios.  On the other hand, scenarios involving dynamical interactions (e.g., collisions and/or ejections) are likely to be hostile to circumstellar disks.  Evidence for circumstellar disks around individual young (few~Myr) brown dwarfs has recently been found, from H$\\alpha$ emission, near-IR excesses, and mid-IR detections.  However, it is difficult to determine the {\\em frequency} of disks around brown dwarfs from studies to date due to a combination of small number statistics, sample selection inhomogeneity, and, most importantly, choice of wavelength.  A priori, brown dwarf disks are expected to be harder to detect than disks around stars because of lower contrast.  Substellar objects are less luminous and have shallower gravitational potentials; hence, the inner regions of their disks are likely to be cooler and thus could have negligible excesses in the $JHK$ (1.1$-$2.4~\\micron) bands, which have been used by most previous studies. ",
        "conclusions": ""
    },
    "0207/astro-ph0207463_arXiv.txt": {
        "abstract": "{This paper discusses  evidence for and properties of disks associated to brown dwarfs in the star-forming region \\rhooph.  We selected nine objects from the ISOCAM survey of Bontemps et al.~\\cite{Bonea01} that have detections in the two mid-infrared bands (6.7 and 14.3 \\um), relatively low extinction and low luminosity. We present  low-resolution near-infrared spectra in the J, H and K bands, and determine for each source spectral type, extinction, effective temperature and luminosity by comparing the spectra to those of field dwarfs and to the most recent model stellar atmospheres. The results indicate that eight objects have spectral types M6--M7.5, effective temperature of 2600--2700 K, one has a later spectral type (M8.5) and lower temperature (about 2400 K). The derived extinctions range between \\AV$\\sim$2 and 8 mag. The location of the objects on the HR diagram, in spite of the uncertainties of the evolutionary tracks for young objects of substellar mass, indicates that all the objects are very young and have masses below about 0.08 \\Msun. The coolest object in our sample  has mass in the range 8-12 \\MJ\\ (0.008--0.012 \\Msun). In all cases, the mid-infrared excess is consistent with the predictions of models of disks irradiated by the central object, showing that circumstellar disks are commonly associated to young brown dwarfs and planetary-mass objects. Finally, we discuss possible variations of the disk geometry among different objects, as well as the possibility of using these data to discriminate between various formation scenarios. ",
        "introduction": "A large number  of objects with sub-stellar mass  are now known, with masses ranging  from the hydrogen burning limit that divides stars from brown dwarfs (BDs; \\Mstar$\\simless 0.075$ \\Msun)  to values comparable to the mass of giant planets and below the deuterium burning limit ($\\simless$ 0.013 \\Msun). Their discovery in regions of star formation has provoked  an intense debate on the formation mechanism of such objects. Do they form, as solar mass stars do, from the collapse of a molecular core (Shu et al.~\\cite{Shu87})? Are they stellar embryos, whose further growth is prevented by dynamical ejections from small stellar systems (Reipurth \\& Clarke~\\cite{RC01}; Bate et al.~\\cite{Bate02})? Or are they ``planets\", i.e., objects that form in gravitationally unstable regions of circumstellar disks (Papaloizou \\& Terquem~\\cite{PT01};  Lin et al.~\\cite{Linea98})? Is there a single formation process for all substellar objects? What is the lowest mass for the gravitational collapse mechanism?  A crucial contribution to this debate is expected from studies  of the circumstellar disks (if any) associated with sub-stellar objects, since different theories make very different predictions.   Disks are a necessary step in any formation mechanism that involves accretion from a parental core. If BDs  form from core collapse, they should be associated to disks similar in properties to those found around low mass \\pms\\ stars (T Tauri stars; TTS). A prediction of the stellar embryo theory is that the disks should be truncated by the ejection mechanism, so that  they should be small and short-lived. In the ``planetary\" hypothesis,  any circumstellar disk should be even less substantial.  In some young BDs, emission in excess of that due to the photosphere has been detected in the near (Oasa et al.~\\cite{OTS99}; Muench et al.~\\cite{MLAL01}; Wilking et al.~\\cite{WGM99}) and  mid-infrared (Persi et al.~\\cite{Pea00}; Comer\\'on et al.~\\cite{Comea98}, \\cite{CNK00}; Bontemps et al.~\\cite{Bonea01}), and has been interpreted, by analogy with TTS,  as evidence for circumstellar disks. In an earlier study (Natta \\& Testi \\cite{NT01}; Paper I), we discussed the properties of three objects in Chamaeleon I for which we could find in the literature ground-based spectroscopy and photometry as well as ISOCAM measurements at 6.7 and 14.3 \\um\\ (Comer\\'on et al.~\\cite{CNK00}; Persi et al.~\\cite{Pea00}). One of these objects is a bona-fide BD,  while the two others are close to the threshold between stars and BDs. We found that the excess emission was clearly detectable only in the mid-infrared, because the stellar photosphere overwhelms the disk emission in  the three near-infrared bands. The observed SEDs are well described by disk models similar to those of TTS,  assuming that the heating is due to irradiation from the central star.  This first result provides strong support for the idea that BDs form like stars, from the contraction of a molecular core. Hence, we decided to extend our study of disk properties to a larger number of substellar mass objects in regions of star formation, possibly down to objects of few Jupiter masses. With this in mind, we  selected a small but well defined sample of nine objects in the \\rhooph\\ region, that were detected by ISOCAM at both 6.7 and 14.3 \\um\\ (Bontemps et al.~\\cite{Bonea01}). We obtained near-infrared spectra for all of them (see \\S 2), which we used to derive the basic parameters of the central objects, namely effective temperature, luminosity and mass (\\S 3). Because of the adopted selection criteria, all of these objects have   excess emission in the mid-IR. We model the expected disk emission for each object and show the results in \\S 4. We discuss the implications of our findings in \\S 5 and present conclusions in \\S 6.   \\section {Observations and data reduction}  \\subsection {Selection criteria}  We chose nine BD candidates from the sample of Class~II objects detected at both 6.7 and 14.3 \\um\\ by Bontemps et al.~(\\cite{Bonea01}). We selected all  objects with visual extinction less than $\\sim$8.5~mag and luminosity less than $\\sim$0.04~\\Lsun\\ according to Bontemps et al.~(\\cite{Bonea01}). The first criterion ensures the possibility of obtaining high signal to noise specta across the entire near infrared range. The low luminosity was required to increase the chance of selecting objects in the range of masses we are interested in.  The location of the selected objects in the ISOCAM color-magnitude diagram is shown in Figure~\\ref{fflcol} (filled circles). \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fflcol.eps}} \\caption{ISOCAM color-magnitude plot (adapted from Bontemps et al.~\\cite{Bonea01}). The symbols are asteriscs for  Class I sources, open circles for Class II sources, stars for Class III sources. Filled circles show the objects in our sample. The dotted line indicates the ISOCAM completeness limit. } \\label{fflcol} \\end{figure} We note that all the objects are close to or below the completeness limit of the ISOCAM survey, as expected for such low luminosities. In colors, our sample span the whole range covered by the Class~II objects (essentially classical TTS).  Some of the selected sources were  known from previous studies to be very low-mass objects. In Table~1 we give the ISO source number, the J2000 coordinates, other designations and references to previous studies. Finding charts are provided in Appendix~\\ref{images}. We will comment on the comparison between the literature source parameters and those derived in this paper in Section~\\ref{sdiscspar}. The results on one of the sources in our sample (\\#033) have been already presented in Testi et al.~(\\cite{Tea02a}, hereafter Paper~II); they have been re-analyzed and reported again here for an easier comparison with the rest of the sample.   \\subsection {Near-infrared spectroscopy}  Near-infrared spectra for the objects in our sample were acquired in the period July 4--9, 2001 at the Telescopio Nazionale Galileo (TNG), using the multi-mode Near-Infrared Camera Spectrograph (NICS; Baffa et al.~\\cite{Bea01}). The Amici device (Oliva~\\cite{O00}), a prism based, high-throughput optical element unique to NICS, was used as disperser, coupled with a 0.5\\arcsec\\ wide, 4.2\\arcmin\\ long slit; the resulting effective resolution is approximately $\\Delta\\lambda/\\lambda\\sim$100, approximately constant across the entire spectral range (0.85--2.45 \\um). An identical instrumental configuration was used for the observations of field dwarfs of known spectral type (Testi et al.~\\cite{Tea01}; \\cite{Tea02b}). Data reduction and calibration was performed as described in Testi et al.~(\\cite{Tea01};~\\cite{Tea02a}).  \\subsection {Broad-band photometry}  On August 1 and 3 2001, we obtained moderately deep Gunn-i integrations using DFOSC and the Danish 1.54~m telescope at the ESO La~Silla observatory. Following standard bias, flat fielding and sky subtraction, typically 6 individual 15~min dithered frames were coadded to produce the final images. Photometric calibration was achieved by observing a set of stars from the Landolt~(\\cite{L92}) catalogue, for which i-band AB magnitudes were computed using the transformations given in Fukugita et al.~(\\cite{Fea96}). Given the uncertainties in the transformations and the non perfect observing conditions, the uncertainties in the  photometry are rather large.  For all sources near-IR J,H,K$_s$ photometry is available from the 2MASS second incremental data release. Additional L' and R-band photometry were taken from Comer\\'on et al.~(\\cite{Comea98}).    \\section {Spectral classification and stellar properties}   \\begin{table} \\begin{flushleft} \\caption{ Sample objects and i-band photometry} \\vskip 0.1cm \\begin{tabular}{ccccc} \\hline\\hline (1)   & (2)    & (3)& (4)  & (5)      \\\\ Object&\\multicolumn{2}{c} {Coordinates}&  i-band  & Other \\\\ (ISO\\#)& \\multicolumn{2}{c}{(J2000.0)}  &  (AB mag)&   Names \\\\ \\hline &            &          &                 &         \\\\ 023  & 16 26 18.8 & -24 26 09& 20.34 $\\pm$0.15 & SKS1-10 \\\\ 030  & 16 26 21.4 & -24 25 59& 16.50 $\\pm$0.10 & SKS1-13\\\\ &            &          &                 &GY5 \\\\ 032  & 16 26 21.7 & -24 44 43& 16.26 $\\pm$0.10 & --     \\\\ 033  & 16 26 22.2 & -24 24 05& 21.73 $\\pm$0.20 & SKS3-13\\\\ &            &          &                 & GY11\\\\ 102  & 16 27 06.5 & -24 41 50& 15.75 $\\pm$0.10 & GY204\\\\ 160  & 16 27 37.4 & -24 17 58& --              & --    \\\\ 164  & 16 27 38.6 & -24 38 39& 18.18 $\\pm$0.10 & SKS1-49\\\\ &            &          &                 &   GY310\\\\ 176  & 16 27 46.3 & -24 31 41& --              & GY350\\\\ \\smallskip 193  & 16 28 12.2 & -24 11 37& --              & --       \\\\ \\hline \\hline \\label{logobs} \\end{tabular} \\end{flushleft} References for Column 5. SKS: Strom et al.~(\\cite{SKS95}); GY: Greene \\& Young~(\\cite{GY92}) \\end{table}  \\begin{table} \\begin{flushleft} \\caption{ Derived Object Properties} \\vskip 0.1cm \\begin{tabular}{clcccc} \\hline\\hline (1)& (2) & (3)& (4)& (5)& (6)\\\\ Object&  ST& \\Teff  & L$_\\star$& \\AV & M$_\\star$ \\\\ (ISO\\#)&    & (K)&     (L$_\\odot$)& (mag)& (M$_J$)   \\\\ \\hline &    &      &     &    &      \\\\ 023  &  M7& 2650 & 0.04& 8.0& 30$-$50 \\\\ 030  &  M6& 2700 & 0.07& 3.0& 40$-$80 \\\\ 032  &  M7.5& 2600 & 0.06& 2.0& 30$-$50\\\\ 033  &  M8.5& 2400 & 0.008& 7.0& 8$-$12\\\\ 102  &  M6& 2700 & 0.08& 3.0& 40$-$80\\\\ 160  &  M6& 2700 & 0.04& 6.0& 30$-$60\\\\ 164  &  M6& 2700 & 0.09& 6.0& 40$-$80\\\\ 176  &  M6& 2650 & 0.07& 7.0& 30$-$70\\\\ \\smallskip 193  &  M6& 2650 & 0.1& 7.5& 40$-$80\\\\ \\hline \\hline \\end{tabular} \\end{flushleft} \\end{table} \\label{stars}   The observed near-infrared spectra obtained at the TNG, normalized to the mean flux in the interval 1.1 -- 1.75~\\um, are shown in Fig.~\\ref{ffield} and ~\\ref{fatm}. We  derive for each object effective temperature and luminosity in the following manner. We first obtain the extinction and spectral type by comparing the source spectra to those of field dwarfs. We then use the derived extinction value to obtain the effective temperature through the comparison with reddened model atmospheres. The luminosity is computed from the dereddened J-band magnitude using the appropriate bolometric correction derived from the model atmosphere.  The first step is illustrated in Figure~\\ref{ffield}, where we compare the observed spectra with a set of reddened field dwarfs, also obtained at the TNG with the same instrumentation (Testi et al.~\\cite{Tea01}, \\cite{Tea02b}). We adopt the extinction law appropriate for \\rhooph\\ (R=4.2; Cardelli et al.~\\cite{CCM89}).  The overall shape of the spectrum from 1 to 2.4\\um\\ depends strongly on the spectral type of the object and extinction along the line of sight. There is, however, a degree of degeneracy, so that a cooler, less reddened object looks similar to a hotter, more reddened one. Therefore, we have also considered other features, such as the shape of the H band, the drop due to water absorption at the red edge of the J band and the intensity of some of the features visible in the J band.  Although some of  these characteristics depend somewhat on the gravity, so that we cannot expect a perfect match between  the young BDs and the field dwarfs, the fits are very good for most objects. Of the nine targets, three have extinction \\AV$\\le$ 3 mag, and six $\\ge$6.0 mag; eight  out of nine objects have a spectral type M6--M7, with $\\pm$one subclass uncertainty, while \\#033 has a later spectral type (M8.5). We give the results in Table~2, Columns 2 and 5. Note that even if the extinction is determined in the wavelength range 0.8-2.4 \\um, for convenience  we express it in terms of \\AV , the extinction in the visual, (\\AV = \\AJ/0.313; Cardelli et al.~\\cite{CCM89}).  In Fig.~\\ref{fatm}, the same TNG spectra are compared to low gravity, $\\log g = 3.5$, model stellar atmospheres (Allard et al.~\\cite{Aea00}; \\cite{Aea01}), smoothed to the appropriate resolution and reddened using the value of A$_V$ derived above. We obtain from this comparison the best value of the effective temperature \\Tstar, as well as a check on the adopted value of \\AV. Our estimates of \\Tstar\\ have an uncertainty of typically $\\pm 100$K; we assign values of 2600--2700 K to all the objects of spectral type M6--M7.5, while  \\#033 is definitely cooler (\\Tstar $\\sim$ 2400). The vales of \\Tstar\\ are given in Table ~2, Column 3. The robustness of these results and the uncertainties in the derivation of A$_V$ and \\Teff\\ are described in more detail in Appendix~\\ref{a33}.   Since all the objects in our sample  have a mid-infrared excess, we have considered the possibility that excess emission is present also in the near-infrared, and is affecting our determination of the stellar parameters. We have therefore subtracted the maximum contribution expected from an irradiated  disk (flared, seen face-on; see \\S 4) from  the observed spectra. We found no significant change in the derived stellar parameters.   The luminosity of the objects is shown in Table~2, Column 4. It has been computed from the dereddened J flux, and the ratio of the J to the total flux given by the appropriate stellar atmosphere model.  These bolometric corrections are virtually identical to those of Wilking et al.~(\\cite {WGM99}) and Leggett et al.~(\\cite {Lea02}). For all stars, the adopted distance is D=150 pc. The uncertainties on \\Lstar\\ are difficult to determine  accurately. We estimate that they probably range from 20\\% to 30\\%, mostly due to uncertainties on \\AV. The bolometric correction for the J band changes very little with the atmospheric parameters, but an additional uncertainty (the same for all stars) may come from the uncertainty in the assumed distance.   Finally, we have performed a last check on the reliability of our estimated parameters using  our i-band photometry. For each star, we computed from   model atmospheres and extinctions  theoretical values of the magnitude in the i, J, H, K and L bands by convolving the flux distribution with the appropriate filter responses. The results are shown in Fig.\\ref{mags}, where we compare model predictions and observations for the 6 stars for which i-band photometry was  obtained. \\footnote[1]{ Note that in Paper~II, Fig.4 shows in the inset broad-band fluxes of \\#033 (GY~11) dereddened by \\AV =7.5 mag, rather than 7.0 as quoted.} The agreement of the i-band observed and predicted magnitudes is generally rather good, given the extreme sensitivity  of the model predictions to the exact shape of the i-band filter, with the possible exception of \\#032, which would need \\AV=3 mag, rather than the 2 mag determined from the comparison with field dwarfs. The corresponding change in luminosity would be of 35\\%.   Fig.~\\ref{HR} shows the location of the nine \\rhooph\\ objects in the HR diagram. In the three panels, we overlay them to three different sets of evolutionary tracks, computed by D'Antona \\& Mazzitelli~(\\cite{DM97}), Chabrier et al.~(\\cite{Chea00}) and Burrows et al.~(\\cite{Bea97}), respectively.  The derived masses (Table~2, Column 6) depend on the adopted tracks, hence we report the corresponding range of values. All objects appear to be very young, with ages lower than 1 Myr and probably of the order of a few $10^5$ yr. It is well known that at such ages evolutionary tracks are not very reliable (Baraffe et al.~\\cite{Bar02}), and that the parameters derived from the location on the HR diagram are only indicative. However, in spite of the uncertainties in both tracks and observations,   we   estimate that our sample contains one  very low mass object (\\#033), with a mass of only $\\sim$8--12~\\MJ\\ (Paper~II), and a  group of objects with masses in the BD range, of which about half (\\#023, \\#032, \\#160 and \\#176) are very likely  BDs.   The clustering of eight out of nine of our objects in a narrow region of the HR diagram is a result of our selection criteria  and can be understood as follows. TTS in \\rhooph\\ have typical ages of 1 Myr, with very few stars as old as 3 Myr (Palla \\& Stahler~\\cite{PS00}). The lack of older BDs in the sample is easy to understand, since the limited sensitivity of the ISOCAM survey (especially at 14.3 \\um) strongly biases towards the highest luminosity, hence the youngest sources, and we  expect to find in our sample only   BDs younger than the average TTS. Older, more massive objects  could, in principle,  fall in our sample. In practice, we found that this was not the case, given \\rhooph\\ typical age. We have applied the procedure adopted by Bontemps et al.~(\\cite{Bonea01}) to a ``theoretical\" star with mass 0.2 \\Msun\\ and age of 2 Myr, using model-predicted J, H, K magnitudes (Baraffe et al.~\\cite{Bea98}) and \\AV $\\leq$9 mag; such a star would have a computed luminosity higher than our selected upper limit (\\Lstar$\\simless$ 0.04 \\Lsun), and would therefore not be included in our sample.  Younger stars would be even more luminous.     \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fig_field.ps}} \\caption{ Observed TNG/Amici spectra of the sample objects. In each panel, we show (solid  line) the spectrum of one object compared with the reddened spectra of field M-dwarfs ( dotted lines) of different spectral types, from Testi et al.~(\\cite{Tea02b}). All spectra are normalized  to the mean flux in the 1.1--1.75 \\um\\ range and shifted with constant offsets for clarity. The field dwarf spectra have been reddened by the value of \\AV\\ shown in each panel. } \\label{ffield} \\end{figure*}  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fig_atm.ps}} \\caption{Same as Fig.~\\ref{ffield}, but in this case the red dotted spectra are reddened theoretical atmospheric models (Allard et al.~\\cite{Aea00}), with \\Tstar\\ as labelled and log g=3.5. } \\label{fatm} \\end{figure*}  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fig_mags.ps}} \\caption{ Comparison between dereddened observed  magnitudes (i, J,H,K, L'), shown by squares, and the prediction of model atmospheres with parameters as in Table 2 (dashed lines) and of the same model atmospheres with additional disk emission (flared, face-on; solid lines) for the six stars for which i-band magnitudes are available. Disk models are described in \\S 4. } \\label{mags} \\end{figure*}   \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fhrd.eps}} \\caption{ HR diagram for three sets of evolutionary tracks: D'Antona \\& Mazzitelli~(\\cite{DM97}) in the left panel, Chabrier et al.~(\\cite{Chea00}) and Baraffe et al.~(\\cite{Bea98}; for the 0.2 \\Msun) in the mid panel, Burrows et al.~(\\cite{Bea97}) in the right panel. Solid lines refer to objects of different mass, as labelled: hydrogen burning stars in black (red), deuterium burning BDs in medium grey (green), objects below the deuterium burning limit in light grey (cyan). Isochrones  are shown as dotted lines, and labelled with the appropriate age. On each panel, the location of the nine observed objects is shown by dots with error bars. } \\label{HR} \\end{figure*}  \\section {Disk models}  All nine  objects  have  mid-infrared fluxes measured with ISOCAM in at least two bands (centered at 6.7 and 14.3 \\um; Bontemps et al.~\\cite{Bonea01}). In three cases, there are additional ISOCAM observations in three narrower bands, centered at 3.6, 4.5 and 6.0 \\um; Comer\\'on et al.~\\cite{Comea98}). The ISOCAM points are shown for each object in Fig.~\\ref{seds}, together with our calibrated and de-reddened TNG spectra.  For each system we  compute the SED predicted by disk models, assuming that the disk is heated by the radiation of the central object. We ignore in this paper any possible viscous heating within the disk (see \\S 6 for a brief discussion). We follow as in Paper~I and~II the method outlined by Chiang \\& Goldreich~(\\cite{CG97}; CG97), with some improvements and modifications  (Natta et al.~\\cite{Nea01}; Chiang et al.~\\cite{Cea01}). CG97 consider a disk in hydrostatic equilibrium in the vertical direction (flared), and describe at each radius the vertical temperature structure of the disk in terms of two components: the disk surface, i.e., the external layer of the disk which is optically thin to the stellar radiation, and the disk midplane. These models allow a quick and reasonably accurate description of the expected SED, more than adequate for the purposes of this paper.  The disk is a scaled-down version of TTS typical disks. It extends inward to the stellar radius, and outward to \\Rd =$1\\times 10^{15}$ cm (67 AU). The total mass is \\Md $\\sim$0.03 \\Mstar, and the surface density varies as R$^{-1}$. The dust in the disk midplane has opacity $\\kappa=0.01 (\\lambda/1.3{\\rm mm})^{-1}$ cm$^2$ g$^{-1}$ (Beckwith et al.~\\cite{BSCG90}). For the dust on the disk surface, we take the mixture of carbonaceous materials and silicates that provides a good fit to the SEDs of several \\pms\\ stars (Natta et al.~\\cite{Nea01}), i.e., a MRN distribution of graphite and astronomical silicates with $dn/da \\propto a^{-3.5}$, $a_{min}$=100 \\AA, $a_{max}$=1 \\um, 30\\% of cosmic C and all Si into grains.  The results of the model calculations are shown in Fig.~\\ref{seds}. The stellar parameters (\\Tstar, \\Lstar, \\Mstar) are taken from Table~2. As pointed out in Paper I, most of the disk parameters are irrelevant for the calculation of the mid-infrared disk emission, or appear in combinations, and cannot be determined individually (see also Chiang et al.~\\cite{Cea01}). As long as the disk midplane remains optically thick to mid-infrared radiation, the only parameters that affect the SED  in the near and mid-infrared are the geometrical shape of the disk (i.e., the flaring angle), the inclination to the line of sight and, to some degree, the disk inner radius $R_i$. There is also some dependence of the shape of the SED on the surface dust model; however, since the luminosity intercepted and re-radiated by the optically thin surface layers is fixed, variations due to (reasonable) changes of the grain properties  are well within the uncertainty of the existing observations.  The upper solid curves in Fig.~\\ref{seds} show the SEDs of flared disks with $R_i$=\\Rstar, seen face-on. They all have strong silicate emission at 10 \\um\\ and a rather flat spectral slope between the two ISO bands at 6.7 and 14.3 \\um,  of order $\\alpha \\sim 0.6-0.8$ ($\\nu F_\\nu \\propto \\nu^\\alpha$). If, rather than extending all the way to the stellar surface the disk is truncated further out, as predicted by magnetospheric accretion models in TTS, at each radius the surface of a flared disk intercepts and reprocesses a larger fraction of the stellar radiation. The disk emission increases correspondingly  at all wavelengths but in the near-infrared, where one is sensitive to the lack of the hottest disk dust.  A  model with $R_i\\sim$ 3\\Rstar\\ is shown (dashed line) for \\#033, where, as discussed in Paper II, the inner hole may account for the large observed mid-infrared excess.  Large variations of the predicted SED occur if the disk shape changes. On each panel, we  show the predictions of geometrically thin, ``flat\" disks (lower solid lines), i.e., disks where the grains are not well mixed with the gas, but have collapsed  onto the disk midplane. Also for these models, we have adopted the CG97 formalism, which remains adequate in all the cases where the disk heating is dominated by the stellar irradiation. If the surface contribution to the SED were negligible, one would recover for these disks the well known temperature profile $T\\propto R^{-3/4}$ and the power-law slope of the SED $\\nu F_\\nu \\propto \\nu^{4/3}$ (Adams and Shu \\cite{AS86}). Our calculations show that also in flat disks the surface contributes to the mid-infrared flux, as shown by the presence in the SED of the silicate feature in emission; however, the midplane emission is larger than the surface contribution at all wavelengths but  in the region $\\sim$ 8--12 \\um, where the silicate feature dominates, so that  the spectral slope between the two ISO points  is  always very close to 4/3. At all wavelengths larger than $\\sim$ 2.5 \\um, the emission of a flat disk is significantly lower than that of a flared one.  Finally, we show on three Panels of Fig.~\\ref{seds} the predictions of tilted flared disks, seen by the observer with inclinations of 69$^o$ (\\#102), 80$^o$ (\\#164) and 86$^o$ (\\#193) respectively (0$^o$ for face-on disks).  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fig_seds.ps}}  \\caption{ Disk and photosphere predicted SEDs. In each panel, the red dots with error bars show the ISOCAM observed fluxes (Comer\\'on et al.~\\cite{Comea98}; Bontemps et al.~\\cite{Bonea01}). The black solid line is the dereddened and calibrated TNG/Amici spectrum. The green jagged line shows the SED of the  photosphere. The combined SED of the photosphere plus disk is shown by blue lines; in each panel, the two solid curves refer to face-on flared  (upper curve) and flat disks (lower curve), with \\Rin=\\Rstar. For \\#033, the dot-dashed curve shows the SED of a face-on, flared disk with \\Rin=3\\Rstar. Finally, we show on three panels the SEDs of tilted flared disks (dotted lines), seen by the observer with inclinations of 69$^o$ (\\#102), 80$^o$ (\\#164) and 86$^o$ (\\#193) respectively (0$^o$ for face-on disks). } \\label{seds} \\end{figure*}    The comparison of the ISO observations to the model predictions shows  that irradiated disk models can account  for the observed mid-infrared excess. More precisely, and in spite of the large uncertainties of the ISO data, inspection of Fig.~\\ref{seds} shows that there are five stars out of nine (\\#030, \\#032, \\#102, \\#160, \\#176) that are extremely well fit by flat disk models. Two objects (\\#023 and \\#033) seem to require flared, face-on disks, while two others have a lower mid-infrared excess, consistent with disks seen  rather edge-on. However, given the large error bars and the model uncertainties, most objects with flat disks are also consistent with flared disk models with large inclination, as shown for the case of \\#102.    \\section {Discussion}   \\subsection {Photospheric Parameters} \\label{sdiscspar}  The agreement between the object spectra and those of field dwarfs is good beyond our expectations.  The largest differences are of 20\\% at most, generally at the peak of the H band, with no systematic difference between objects with large or low extinction, nor between stronger and weaker sources. Even the relatively narrow features that appear in the spectra around 1.1 \\um\\ are often well matched in the two sets of spectra. This indicates that, at the resolution of our observations, one should not expect strong gravity effects. We have checked that this is indeed the case by comparing model atmosphere spectra smoothed to the observed resolution for stars of different gravity (Allard et al.~\\cite{Aea01}), ranging from 3.5 to 6.0. All the models with gravity in the interval 3.5--5.0 are practically identical, at our spectral resolution and in this temperature range.  The comparison of our spectra with model atmosphere predictions is somewhat less satisfactory, especially in the  H band, where the shape of the feature  peaked at about 1.7 \\um\\ (resulting from water absorption features at shorter and longer wavelengths) is narrower in the models than observed, and around 1.3 \\um, where the models tend to predict more emission than is observed. Note, however, that this is not always the case (see, for example, \\#032 and \\#164). Still, the agreement is in general rather good, with differences that never reach more 30\\%, again with no dependence on the extinction nor on the observed signal.  The comparison of our determinations of the photospheric parameters of individual objects with previous spectroscopic determinations in the literature shows that in some cases there is good agreement, while in others there are discrepancies that are not easily understood. For example, Wilking et al.~(\\cite{WGM99}) assign similar spectral types to \\#023, \\#030, but a significantly later one (M8.5) to \\#164, based on K band R$\\sim$300 spectroscopy. For the same object, Luhman \\& Rieke ~(\\cite{LR99}) estimate a spectral type M7, based on intermediate resolution K band spectroscopy, similar to our classification M6. The same authors, on the other hand, attribute to \\#030 a somewhat earlier spectral type (M5-M6). The case of \\#033 (GY~11) has been discussed in detail in Paper~II. A likely reason for differences in the spectral classification is that our scheme is based on the overall spectral shape, while the others rely on fitting individual spectral features, which in the infrared show large scatter for late M objects (e.g., Luhman \\& Rieke~\\cite{LR98}).  On the more general issue of the effective temperature scale of  young BDs, we attribute temperatures in the range 2600-2700 K to our group of objects with spectral types M6--M7.5. Our only object with later spectral type (M8.5) has \\Tstar =2400$\\pm$100 K. In a preliminary analysis of our sample field dwarfs (Testi et al.~\\cite{Tea02b}), we derive a similar effective temperature-spectral type correspondence. This is not significantly different from the scale used by Wilking et al.~(\\cite{WGM99}) in their study of candidate BDs in \\rhooph. It is, however, at odds with some recent results, that tend to attribute to young BDs of similar spectral types temperatures higher than our values (Lucas et al.~ \\cite{Lucea01}; Lodieu et al. ~\\cite{Lodea02}). Further work, on larger samples of BDs in young star forming regions is clearly required.   \\subsection{ The disk hypothesis}  The comparison between  models and observations, discussed in the previous section, proves that the mid-infrared excess associated to many young BDs can be accounted for by the emission of circumstellar disks heated by the radiation of the central object.   Few disk properties are  constrained by the existing observations, and we do not want to overinterpret our results, given the large uncertainties of the observed fluxes, and the simplicity of the adopted models. However, in our limited sample of nine stars we find disks of different flavours, and, in particular, an indication that many BDs  may have flat disks. If we consider also the three objects in Cha I studied in Paper I, we have three objects with clear evidence of flared disks, and nine where flat disks seem more appropriate, although we cannot  rule out almost edge-on flared disks for some of them (see also Apai et al.~\\cite{apai02}). This is  potentially an interesting result, since it seems natural to associate flat disks with dust sedimentation toward the midplane. In our selection of ISO sources, we have an obvious strong bias against objects with flat disks, since we required that the sources  were detected by ISO in both bands. So, the fact that our objects with the lowest 6.7 \\um\\ fluxes (\\1\\ and \\#033) have flared disks is not surprising. However, there is   no bias against selecting flared disk objects of higher luminosity, and we find   only one (\\#023). The possibility of dust settling in these very young low-mass objects is intriguing. However, it needs to be confirmed by high-quality photometric observations at longer wavelengths, before entering into further speculations.  The ejected embryos hypothesis does not exclude that BDs may have a small, and therefore short-lived, circumstellar disk. Estimates by Bate et al.~(\\cite{Bate02}) give  disk radii of about 20 AU or less. The existing infrared data do not allow us to rule out such possibility, since the SED of a model with \\Rd=20 AU will differ from the SED of a disk with \\Rd=75 AU only at wavelengths $\\simgreat$ 40 \\um. The mass of the disk is not predicted by the Bate et al.~(\\cite{Bate02}) calculations, nor constrained by the existing observations, since the only constraint we can set is that the disk has to be optically thick in the mid-infrared. This, however, only requires a disk mass of 10$^{-5}$--10$^{-6}$ \\Msun (or \\Md/Mstar$\\sim 10^{-4}$), which is still consistent with a typical disk (having \\Md/Mstar$\\sim$ 0.03, \\Rd=75 AU), truncated at \\Rd=20 AU. Until far-infrared and millimeter data  become available, the only way to validate these models is to determine the fraction of disks in  unbiased samples of BDs of known age.    Finally, one should remember that our analysis relies on the assumption that the ISO sources coincide with the objects we identify in the near-infrared. In some cases, this is likely to be true (see Appendix~\\ref{images} and the discussion of \\#033 in Paper~II). In other cases, it is impossible  to check the validity of this assumption, given the large ISO beam and the presence of other red objects in the near-infrared images. However, the good agreement between the observations and the model predictions, which depend essentially only on the stellar properties we derive from the spectroscopy, is encouraging. Further tests of the association of the observed mid-infrared excess with the identified stars could be obtained by accurate images in the L and M bands, where we predict that the disk emission should be dominant (see Fig.~\\ref{seds}).      \\section {Conclusions}  We have discussed in this paper  a sample of nine very low-mass objects in the \\rhooph\\ star forming region that have evidence for circumstellar warm dust. We selected from the ISOCAM sample of Bontemps et al.~(\\cite{Bonea01}) those objects that have mid-infrared detections in both the 6.7 and the 14.3 \\um\\ bands, relatively low extinction and low luminosity. We determined first if these BD candidates were indeed bona-fide BDs, and then we checked if the observed infrared excess was consistent with the predictions of disk models, similar in properties to those associated to T Tauri stars.  Our strategy was very successfull. The low-resolution near-infrared spectra obtained at the TNG allowed us to determine  for each object spectral type and extinction, by comparison with field dwarfs observed with the same instrumental set-up, as well as  effective temperature and luminosity, by comparison with model atmosphere predictions. The comparison with various sets of evolutionary tracks on the HR diagram shows that all the nine sources are very young, low-mass objects. In particular, one (\\#033 or GY 11, already discussed in Paper II) has a mass of 8-12 \\MJ, while the others have masses in the BD mass range; four of them are very likely bona-fide BDs.  In all objects, the mid-infrared excess  is consistent with the predictions of disks irradiated by the central object. We find no evidence of strong accretion occurring in these systems, based on the fact the observed near-infrared fluxes are dominated by the emission of the photospheres, and there is very little contribution (if any) from hot dust. However, it is not clear to which degree the near-infrared excess  in  very low-luminosity objects is a sensitive indicator of accretion (see, for an example of an actively accreting object with no near-infrared excess, Fern\\'andez \\& Comer\\'on \\cite{FC01}), and this issue should be explored more quantitatively in the future.  The existing data indicate that the disks must be optically thick at mid-infrared wavelengths; in some cases they must be flared (i.e., gas-rich with well-mixed dust and gas), while in others it is possible that they are geometrically flat, i.e., that dust has settled to the disk midplane. However, data at longer wavelengths are necessary to further investigate this point, and we do not want to put too much weight on this rather weak evidence.  In the same sobering vein, we want to point out that our results do not  discriminate yet between  different formation mechanisms, namely between the possibility that BDs form from the gravitational collapse of individual, very low-mass cores, and the ejected embryo theory. We fit the observed mid-infrared excess with a scaled-down version of disks around the more massive TTS.  This, however,  just implies   that ``normal\" disks can account for the existing observations, since few parameters are actually constrained. As already pointed out in Paper I and II, only observations at long wavelengths can measure the disk radius and mass, since the lower limits that we can derive from the conditions that the disk is optically thick in the mid-infrared are hardly significant.  Having stressed all the limitations of our results, let us now point out that this is the first sample of very low mass objects in a star forming regions where evidence for circumstellar disks has been found and investigated in detail. Our accurate near-infrared spectroscopy, which allows us to estimate a reliable value of the mass of the objects, proves  that disks exist around  low mass objects, well into the range of brown dwarfs. In one case, \\#033, our data provide strong indications that an object with mass close to or below the deuterium burning limit also has a circumstellar disk.  In addition to providing the beginning of a census of disk properties around BDs, our models  indicate that the excess due to the cold disks irradiated by a central BD can only be detected by deep photometry in the L and M bands. We expect that  major progress in our understanding of BD formation will be obtained by combining near-infrared low resolution spectroscopy with photometry in J,H,K,L,M of unbiased (i.e., not a-priori selected because they have a mid-infrared excess, as here) samples of BD candidates in star forming regions of different age.     \\begin{acknowledgements} We thank Carsten Dominik and Michael Meyer for useful discussions. It is a pleasure to acknowledge the TNG and ESO staff for their excellent support during observations. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This work was partly supported by  ASI grant  ARS 1/R/27/00 to the Osservatorio di Arcetri. \\end{acknowledgements}   \\appendix ",
        "conclusions": ""
    },
    "0207/astro-ph0207180_arXiv.txt": {
        "abstract": "We examine power spectra from the Abell/ACO rich cluster survey and the 2dF Galaxy Redshift Survey (2dFGRS) for observational evidence of features produced by the baryons. A non-negligible baryon fraction produces relatively sharp oscillatory features at specific wavenumbers in the matter power spectrum. However, the mere existence of baryons will also produce a global suppression of the power spectrum. We look for both of these features using the false discovery rate (FDR) statistic. We show that the window effects on the Abell/ACO power spectrum are minimal, which has allowed for the discovery of discrete oscillatory features in the power spectrum. On the other hand, there are no statistically significant oscillatory features in the 2dFGRS power spectrum, which is expected from the survey's broad window function. After accounting for window effects we apply a scale-independent bias to the 2dFGRS power spectrum, $P_{Abell}(k) = b^2P_{2dF}(k)$ and $b = 3.2$. We find that the overall shapes of the Abell/ACO and the biased 2dFGRS power spectra are entirely consistent over the range $0.02 \\le k \\le 0.15h$Mpc$^{-1}$. We examine the range of $\\Omega_{matter}$ and baryon fraction, for which these surveys could detect significant suppression in power. The reported baryon fractions for both the Abell/ACO and 2dFGRS surveys are high enough to cause a detectable suppression in power (after accounting for errors, windows and $k$-space sampling). Using the same technique, we also examine, given the best fit baryon density obtained from BBN, whether it is possible to detect additional suppression due to dark matter-baryon interaction. We find that the limit on dark matter cross section/mass derived from these surveys are the same as those ruled out in a recent study by Chen, Hannestad and Scherrer. ",
        "introduction": "During its first $\\simeq100,000$ years, the Universe was filled with a fully ionized plasma with a tight coupling between the photons and baryons via Thomson scattering. A direct consequence of this coupling is the acoustic oscillation of both the primordial temperature and density fluctuations (within the horizon) caused by the trade--off between gravitational collapse and photon pressure.  The relics of these acoustic oscillations are predicted to be visible as alternating peaks and valleys in the CMB temperature power spectrum. The relative amplitudes and locations of the features provide powerful constraints on the cosmological parameters ({\\it e.g.} $\\Omega_{total}, \\Omega_{b}$, etc). The BOOMERANG, MAXIMA, and DASI experiments announced the first high confidence detection of these acoustic oscillations in the temperature power spectrum of the Cosmic Microwave Background (CMB) radiation (Miller et al. 1999; Melchiorri et al. 2000; Balbi et al. 2000; Lee et al. 2001; Netterfield et al. 2002; Halverson et al. 2002; de Bernardis et al. 2002; Miller et al. 2002a).  Like the CMB experiments mentioned above, we have also seen the recent emergence of large and/or deep extragalactic datasets to directly measure the 3-dimensional luminous matter power spectrum, $P(k)$. Specifically, the PSCz galaxy survey (Saunders et al. 2000), the Abell/ACO cluster survey (Miller et al. 2002b), and the 2dF Galaxy Redshift Survey (Colless et al. 2001) have all released newly measured power spectra (Hamilton, Tegmark, and Padmanabhan, 2000; Hamilton and Tegmark 2002; Miller \\& Batuski  2001; Miller, Nichol \\& Batuski 2001a,b; Percival et al. 2001; Tegmark, Hamilton \\& Xu 2002). These surveys are beginning to allow us to search for oscillatory features in the matter power spectrum (frozen into the matter distribution from acoustic oscillations on the largest-scales). Using some of these large (in volume and/or number) surveys, Miller et al. (2001a) found statistically significant features (dips in power) at $k \\sim 0.04$ and $0.1h$Mpc$^{-1}$ in the $P(k)$. While much less obvious, features (at similar wavenumbers) were also qualitatively observed in the recently determined 2dFGRS $P(k)$ (Percival et al. 2001) and the 2dF QSO power spectrum (Hoyle et al. 2002). The existence of such oscillatory features in the matter power spectrum permits one only to say whether they are consistent with baryonic features, since features could also be due to other reasons (e.g. features in the primordial $P(k)$).  Besides the features resulting from acoustic oscillations, the presence of a non-negligible baryon fraction in the Universe has another observational consequence on the power spectrum: the overall power below the sound horizon is suppressed over a large range in $k$ (see Eisenstein and Hu 1998--hereafter EH98). The suppression occurs in both the cold dark matter and baryonic transfer functions. As discussed in EH98, the main effect of the baryons is the suppression of the growth rates between the equality and drag epochs. Therefore, the effect of baryons in the matter power spectrum can be detected via this suppression over many $k$-modes, as opposed to looking for the sharp features which result from the acoustic oscillations.  However, there are some ``degeneracies'' in the $P(k)$ suppression. In addition to the baryonic damping discussed above, dark matter (hereafter--DM) scattering off the baryons will also suppress the power spectrum on scales smaller than the sound horizon (Chen, Hannestad, and Scherrer 2002), and hot dark matter will also have the effect of suppressing power on small scales. Eisenstein and Hu (1999) note that the suppression of the transfer function due to  baryons will dominate over neutrinos.  Others have begun to address the effects of hot dark matter (see Wang, Tegmark, and Zaldariagga 2002 and Elgaroy et al. 2002), and so we focus here on suppression effects related to the baryons: baryon suppression and DM-baryon scattering.  In Section 2 we compare the 2dFGRS power spectrum to that of the Abell/ACO cluster spectrum, taking into account effects from their respective window functions. By controlling the False Discovery Rate (Miller et al. 2001a,c), we look for any statistically significant oscillatory features in the 2dFGRS power spectrum. In Section 3, we then examine whether either of these surveys could have detected the baryons through suppression. In Section 4, we examine the effect of DM--baryonic scattering interactions on the $P(k)$. We then summarize and discuss our findings in Sections 5 and 6. ",
        "conclusions": "We have examined two large surveys for the effects of baryons on the luminous matter power spectrum. We have shown that the window effects on the measured $P(k)$ are minimal ($<5$\\%) for the Abell/ACO survey, while they are significant in the 2dFGRS. The 2dFGRS window not only smoothes out any features, but it also moves those features to higher $k$.  We have shown that the features detected in the Abell/ACO $P(k)$ by Miller, Nichol, and Batuski (2001), are entirely erased by the 2dFGRS window. There are features in the 2dFGRS power spectrum visible to the eye, but they are not significant after a thorough statistical analysis.  After convolving the Abell/ACO power spectrum with the 2dFGRS window, we find that the shapes of the two power spectra are entirely consistent from $k = 0.02h$Mpc$^{-1}$ to $k = 0.15h$Mpc$^{-1}$. Neither survey shows evidence for a turnover in power toward a scale-invariant spectrum.  Since the presence of baryons suppresses the power spectrum on scales smaller than the sound horizon, we examined whether either of these surveys could have detected this suppression. We utilized a multiple hypothesis testing approach (by controlling the False Discovery Rate) as opposed to an omnibus test (e.g. a $\\chi^2$ test). This allowed us to demand that the detection of the overall suppression, and not simply the baryonic features. Since neither the 2dFGRS nor the Abell/ACO have accurate power measurements on the largest scales (where there is no baryonic suppression), we consider model power spectra which have the same attributes (windows, sampling, errors, etc) of the two surveys. We find that the Abell/ACO survey and the 2dFGRS survey are complementary in that the former can probe to smaller $\\Omega_m$ due to its volume and window, while the latter can probe to smaller $\\Omega_b/\\Omega_m$ due to its better sampling and smaller errors.  Recently, Chen et al. (2002) had shown that scattering interactions between dark matter and the baryons will also cause significant suppression in the matter power spectrum. We fix $\\Omega_m$ and $\\Omega_b$ and examine a range of dark matter masses and DM-baryon cross sections to see whether either of these surveys could detect this suppression. We find that such interaction is detectable if the interaction strength is above or close to the current limit put by the CMB and large scale structure data. However, we note that if the interaction strength between the dark matter and baryon is as strong as the Spergel-Steinhardt dark matter self-interaction as speculated in Wandelt et al. (2000), it should already be detectable."
    },
    "0207/astro-ph0207149_arXiv.txt": {
        "abstract": "Evidence of TeV emission from GRB970417a has been previously reported using data from the Milagrito detector~\\cite{atkins00b}. Constraints on the TeV fluence and the energy spectrum are now derived using additional data from a scaler system that recorded the rate of signals from the Milagrito photomultipliers. This analysis shows that if emission from GRB970417a has been observed, it must contain photons with energies above 650 GeV.  Some consequences of this observation are discussed. ",
        "introduction": "Some of the most important contributions to our understanding of gamma-ray bursts (GRBs) have come from observations of afterglows over a wide spectral range. Comparisons between these observations and predictions of GRB afterglow properties both as a function of time and of wavelength have provided stringent tests of GRB models~\\cite{paradijs00}. However, far less is known about the multiwavelength spectrum during the prompt phase of GRBs because of its very short duration.  Almost all GRBs have been detected in the energy range between 20 keV and 1 MeV~\\cite{fishman95}. A few have been observed above 100 MeV by EGRET~\\cite{schneid92,sommer94,hurley94,schneid95} indicating that at least some GRB spectra extend up to hundreds of MeV. However, the upper extent of GRB energy spectra is unknown. There may be a second (higher energy) component of emission, similar to that seen in several TeV sources~\\cite{dermer99,pilla98}.  The Milagro gamma-ray observatory, which began full operations in January 2000, is a wide field-of-view instrument that operates with a duty cycle near 100 \\%. It is particularly well suited to extending observations of the prompt phase of GRBs up to TeV energies. A prototype of Milagro, Milagrito~\\cite{atkins00a}, found evidence for TeV emission from one of the 54 gamma-ray bursts observed by BATSE that were within the Milagrito field of view with a probability of $1.5 \\times 10^{-3}$ of being a background fluctuation~\\cite{atkins00b}.  An excess of events was observed from the direction of GRB 970417a during the time BATSE observed emission. In this paper, we use additional data from Milagrito to examine the fluence and spectrum of the possible high-energy emission of GRB970417a.  Milagrito detected secondary particles reaching the ground produced by the interaction of TeV gamma rays in the atmosphere. The Milagrito detector consisted of a single layer of 228 photomultiplier tubes (PMT) placed on a 2.8 m $\\times$ 2.8 m grid approximately 1 meter below the surface of a large covered pond of water~\\cite{atkins00a}. An air shower was registered when relativistic charged particles, radiating Cherenkov light in the water, caused $\\ge$ 100 PMTs to detect light within a 200-ns time interval.  The direction of the gamma ray initiating the shower was reconstructed from the relative timing of the PMT signals.  It is very difficult to obtain information on the energy of the individual events contributing to the TeV gamma-ray excess observed with Milagrito. While the observed number of shower particles is related to the energy of the primary gamma-ray, it also depends on the height in the atmosphere of the first interaction and on the distance of the shower core from the pond. The area of Milagrito was small relative to the lateral extent of a typical shower, so it is usually not possible to determine the location of the core.  Consequently, the number of PMTs hit in the pond is only weakly related to the energy of the gamma-ray primary. Also, the trigger required 100 of the 228 PMTs to register a signal, so there is very little dynamic range over which to identify a variation due to different source spectra.  Instead, information about the distribution of the gamma-ray energies can be obtained using additional information about the summed count rates of the individual PMTs in the pond. Very-high-energy (VHE) gamma rays of too low an energy to trigger the detector can give rise to an increase in these rates. Measurement of the rates therefore allows constraints to be placed on the spectrum of the putative TeV gamma-ray flux observed from GRB970417a. ",
        "conclusions": "As shown above, if the excess triggered events observed by Milagrito were indeed associated with GRB970417a, then they must be due to photons of at least 650 GeV, the highest energy photons detected from a GRB. Figure~\\ref{fig:nufnu} shows that the VHE fluence must be at least an order of magnitude greater than the sub-MeV fluence measured by BATSE. In the following, assuming that high-energy emission from GRB970417a was detected by Milagrito, this section explores some of the consequences of the observation.  The Milagrito observations at VHE energies are not consistent with a straightforward extrapolation of the spectrum measured by BATSE, so they would be the first evidence for the existence of a second (higher energy) emission component in a gamma-ray burst. We note that while the implied fluence at VHE energies is considerably larger than that measured by BATSE, the actual peak of the spectral energy distribution may lie below the energy range measured by BATSE. Thus, while it appears that the power in the high energy component is greater than that in the low energy component, this need not necessarily be the case since a low energy peak was not measured. However, this would be a very unusual GRB spectrum, as peak energies are typically between 100 keV and 1 MeV~\\cite{preece}.  The existence of a higher energy component of emission is predicted by many emission models. For models in which the sub-MeV emission is due to non-thermal emission from a population of relativistic electrons, a fraction of these photons can be upscattered to the GeV-TeV range via inverse Compton scattering~\\cite{meszaros93, chiang99, pilla98}. The emergence of this component is expected to be prompt and coincident with the synchrotron radiation seen in the keV-MeV range.  Protons may also be accelerated to very high energies: GeV-TeV gamma rays can be produced by synchrotron radiation from these protons, or from decaying pions produced by high energy protons interacting with photons~\\cite{boettcher98}.  The possible VHE detection of GRB 970417a has significant implications on the distance scale and energy production in the GRB.  VHE gamma rays interact with lower-energy photons to produce electron-positron pairs~\\cite{gould66,jelley66,primack99,salamon98}.  Thus gamma rays from distant sources may suffer significant attenuation on the intergalactic background radiation fields as they traverse the Universe en route to Earth. There is an energy dependent horizon beyond which gamma rays cannot be detected.  The magnitude of this effect is a function of gamma-ray energy, the density and spectrum of the background radiation fields, and the distance to the source.  The observation of photons with energies of at least 650 GeV coupled with an estimate of the opacity for VHE photons would constrain the redshift of GRB 970417a.  The results of Primack et al (1999) imply that the opacity of the Universe to 650 GeV photons is one at a redshift of $\\sim$0.1.  While it is possible that GRB 970417a may lie beyond this, it would then require an enormous source flux of VHE photons to produce the excess observed by Milagrito.  If the redshift of GRB970417a is less than ~0.1, then with the exception of GRB980425 it is much closer than all of the GRBs with measured redshifts. However only long, bright bursts have been localized sufficiently well to allow the measurement of afterglows, so the sample of GRBs with measured redshifts may not be representative of the entire class of GRBs. While GRB970147a was a long burst, it was also dim and would not have been localized by past or current detectors.  There have been several attempts to extend information on the GRB redshift distribution to bursts without measured redshifts. Several authors have investigated the gamma-ray properties of bursts with known redshifts (and thus luminosities) to find observable gamma-ray parameters that may be indicators of the luminosity. The spectral evolution of pulse structures appears to be anticorrelated with peak luminosity (i.e. bursts with a long lag between low-energy and high-energy detection by BATSE are dimmer)~\\cite{norris00} and bursts with a greater degree of variability appear to be more distant~\\cite{fenimore01}.  Although GRB970417a was too dim at sub-MeV energies to allow the calculation of the variability parameter to be used as a luminosity indicator, it was included in a recent study of the lag-luminosity relation~\\cite{norris02}. However, because GRB970417a was both dim and of relatively short duration, the large uncertainty in the measured lag precludes obtaining a reliable estimate of its luminosity~\\cite{norris02a}.  Other distance indicators are obtained from the global properties of GRBs due to deviations from Euclidean geometry for sources at cosmological distances.  For example, Schmidt (2001) shows that GRBs with harder MeV spectra are more distant than those with soft spectra, and in fact the spectrum of GRB970417a, as measured by BATSE, is soft.  In addition to suffering attenuation via pair production in intergalactic space, VHE photons may also be absorbed in the source itself. The very rapid variability (small emission region) and high luminosity of GRBs implies a very large photon density. If the source is non-relativistic, the optical depth of high-energy photons is so large that the photons could not emerge. If the emission region is moving relativistically, then the pair production optical depth is decreased because the photon energies and densities in the rest frame of the emission region are lower than they appear to the observer~\\cite{baring97,lithwick01}. Using the method of (Lithwick and Sari), we can use the requirement that the burst be optically thin to 650 GeV photons to place a lower limit on the Lorentz factor ($\\gamma$) of the expansion. Assuming a variability timescale of 1 second, a broken power law spectral fit to the BATSE data and a redshift of 0.1 we find a lower limit to the Lorentz factor of 95 for GRB970417a. Therefore, unusually large Lorentz factors are not required for GRB970417a, primarily because the implied luminosity of such a nearby burst is low resulting in a relatively low photon density.  The future of these studies holds great interest. Milagro, a more sensitive detector than Milagrito, is now operating.  SWIFT will launch in 2003 and is expected to detect and localize several hundred bursts per year.  A large fraction of the SWIFT detections will have measured distances.  Milagro will observe these nearby GRBs detected by SWIFT to determine the fraction of GRBs with TeV emission as well as the flux of that TeV emission."
    },
    "0207/astro-ph0207655_arXiv.txt": {
        "abstract": "The Next Generation Sky Survey (NGSS) is a proposed NASA MIDEX mission to map the entire sky in four infrared bandpasses -- 3.5, 4.7, 12, and 23 $\\mu$m. The seven-month mission will use a 50-cm telescope and four-channel imager to survey the sky from a circular orbit above the Earth. Expected sensitivities will be half a million times that of COBE/DIRBE at 3.5 and 4.7 $\\mu$m and a thousand times that of IRAS at 12 and 23 $\\mu$m. NGSS will be particularly sensitive to brown dwarfs cooler than those presently known. Deep absorption in the methane fundamental band at 3.3 $\\mu$m and a predicted 5-$\\mu$m overluminosity will produce uniquely red 3.5-to-4.7 $\\mu$m colors for such objects. For a limiting volume of 25 pc, NGSS will completely inventory the Solar Neighborhood for brown dwarfs as cool as Gl 229B. At 10 pc, the census will be complete to 500 K. Assuming a field mass function with $\\alpha = 1$, there could be one or more brown dwarfs warmer than 150 K lying closer to the Sun than Proxima Centauri and detectable primarily at NGSS wavelengths. NGSS will enable estimates of the brown dwarf mass and luminosity functions to very cool temperatures and will provide both astrometric references and science targets for NGST. ",
        "introduction": "Large-area surveys such as the Two Micron All Sky Survey (2MASS; see other contribution by Kirkpatrick) and the Sloan Digital Sky Survey (SDSS; see contribution by Covey) have revolutionized the field of brown dwarf science. However, these surveys are able to probe only to temperatures near 1000 K at a distance of 10 pc, which leaves the realm of cooler brown dwarfs almost entirely unexplored. Fortunately, the Next Generation Sky Survey (NGSS) has been proposed in a large part to characterize cooler brown dwarfs in the Solar Neighborhood. In \\S2 the NGSS project will be briefly described. Its role in the future of brown dwarf studies is highlighted in \\S3, and the current status of the mission is given in \\S4. ",
        "conclusions": ""
    },
    "0207/gr-qc0207048_arXiv.txt": {
        "abstract": "We summarize recent results concerning the evolution of second order perturbations in flat dust irrotational FLRW models with $\\Lambda\\ne 0$. We show that asymptotically these perturbations tend to constants in time, in agreement with the cosmic no-hair conjecture. We solve numerically the second order scalar perturbation equation, and very briefly discuss its all time behaviour and some possible implications for the structure formation. ",
        "introduction": "First order perturbations of Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) models fail to account for a number important relativistic properties that arise when non-linearities are taken into account. An important example of this is the so called mode coupling. Even if one considers purely scalar perturbations, a comoving scale $k$ at second order is sourced by any other scale $k^\\prime$. In addition, the scalar, vector and tensor perturbations which are decoupled to the first order, become coupled once nonlinear perturbations are taken into account. Thus for example taking into account such couplings at second order results in the initial pure scalar metric perturbations to necessarily generate second order vector and tensor perturbation modes (see e.g.\\cite{MMB}). In view of this, it is important to develop higher order perturbation schemes which go beyond the first order and can thus account for the physical effects which are not taken into account up to the linear order. This would also improve the level of accuracy of results as well allowing the study of stability of the results obtained using first order perturbations.  Among concrete motivations for the use of a second order perturbative scheme is the need for more accurate results from the new generation of gravitational wave detectors as well as the increased computational precision required in order to analyse the cosmic microwave background anisotropies.\\cite{Pyne-Carroll,Mollerach-Matarrese} Both these problems are likely, in future, to require analyses that go beyond the linear order. In addition, the non-linear analysis in the context of cosmological modeling is also relevant on scales much smaller than the cosmological horizon where it might substantially modify the first order results\\cite{Campos-Tomimura} and therefore, as we shall also see below, become important for the study of the structure formation. Furthermore, as far as observations are concerned, second order effects can be important on small scales in order to account for non-linear effects such as the Rees-Sciama\\cite{Rees-Sciama} and gravitational lensing effects (see e.g.\\cite{Schneider-etal}). For a list of references concerning these questions we refer the reader to\\cite{Mollerach-Matarrese}.  Second order perturbations have so far not been widely studied in general relativity. In the cosmological context they seem to have been first employed by Tomita\\cite{Tomita} to study the evolution of scalar perturbations in the Einstein--de-Sitter model using a synchronous gauge. This study was repeated by Mataresse et. al.\\cite{Matarrese-silent}, who obtained similar results using a comoving approach and Russ et al.\\cite{Russ-etal} who have included second order terms resulting from a coupling between growing and decaying scalar perturbation modes. There have also been recent works\\cite{Mukhanov-etal96} showing that the back reaction of the second order perturbations (to be precise, perturbations quadratic in first order terms) are important in early universe scenarios. Furthermore, it has been shown that in inflationary models the magnitude of the second order perturbations can be comparable to those in the first order and therefore a non-linear perturbative analysis may be crucial.\\cite{Bassett2}  The domain of applicability of the theory of second order perturbations that we shall describe is that of small perturbations about an homogeneous and isotropic background. We shall employ it as an approximate framework in order to study the effect of non-linearities in structure formation scenarios as well as the study of the stability of cosmological results obtained using first order perturbative schemes with respect to nonlinear perturbations.  In this paper we shall summarise recent results concerning the exact asymptotic solutions of the second order perturbation equations. An immediate consequence of these results was to prove the non-linear asymptotic stability of the de-Sitter attractor\\cite{BMT2001}, and thus generalise the cosmic no-hair conjecture.\\cite{Gibbons-Hawking} In general, however, the perturbation equations cannot be solved analytically. We shall show preliminary results concerning the numerical integration of the scalar perturbation equation for all times. Finally, we discuss possible implications on the structure formation. We use units in which $8\\pi G=c=1$. Greek indices take values $1,2,3$. ",
        "conclusions": "We have summarised our very recent results\\cite{BMT2001,Mena-thesis} on the evolution of second order perturbations in flat irrotational dust FLRW models. We have shown that second order perturbations tend asymptotically to constants in agreement with the cosmic censorship conjecture.  In order to obtain the behaviour of such perturbations at intermediate times, we have also integrated numerically the scalar perturbation equation and found different transient behaviours over intermediate time scales, depending upon the choice of the initial conditions. This can have interesting consequences for nonlinear phenomena that evolve in time, such as structure formation."
    },
    "0207/astro-ph0207525_arXiv.txt": {
        "abstract": "\\noindent We consider the dimming of photons from high redshift type 1a supernovae through mixing with a pseudoscalar axion field in the intergalactic medium.  We model the electron density using a log-normal probability distribution and assume frozen in magnetic fields.  Assuming the magnetic fields are produced in the early universe we are unable to obtain sufficient dimming in order to explain the apparent acceleration without violating the bounds on the frequency dependence of the dimming.  We also show that any axion mixing leading to a reduction in optical luminosities would also lead to a significant reduction in the polarisation of UV light from intermediate redshift objects which may be detected in the future. ",
        "introduction": "At the time of writing, the present understanding of the energy content of the universe appears to be converging towards a standard model.  Observations of the cosmic microwave background radiation seem to suggest that the total amount of energy in the universe is such that there is zero spatial curvature on the largest observable scales \\cite{maxboom}.  At the same time, observations of the distribution of matter at smaller scales give us information about galaxy clustering and the time of matter-radiation equality, and tell us that the fraction of this total energy which is in the form of matter is only about $30\\%$ \\cite{2df}. This is much larger than the fraction of the total energy which can be in the form of baryonic matter \\cite{pagel}, and together with observations of the rotation curves of galaxies and velocity dispersions of galaxies within clusters, has lead to the concept of dark matter.  However, there still remains the question of the remaining unaccounted for $\\sim 70\\%$ of the total energy density.  Although the data set is effectively still quite small, observations of type 1a supernovae seem to have cast some light on this missing energy \\cite{sn,sndim}.  The apparent dimness of these supernovae at high redshift suggests that there is a non-zero cosmological constant which makes up the missing energy or at least something which acts like one.  If we are to interpret this cosmological constant as some energy field then we find that its energy density $\\rho\\sim (10^{-3} eV)^{4}$ is far away from any of the energy scales (i.e. the Planck scale or the scale of supersymmetry breaking) one would naively expect such a field to have if it were non-zero.  It is also not clear how such a small energy scale could be stable to quantum corrections.  Recently it has been suggested that the dimness of high redshift supernovae is not due to the accelerated expansion of the universe, but rather due to mixing between the photons coming from these objects and a pseudo-scalar axion field \\cite{Csaki:2001yk}.  The Lagrangian density of the photon-axion system is given by \\ben \\mathcal{L}=-\\frac{1}{2}(\\partial^\\mu a\\partial_\\mu a+m^2_{a}a^2) +\\frac{a}{M_a}F_{\\mu\\nu}\\widetilde F^{\\mu\\nu}-\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu} \\een where $F_{\\mu\\nu}$ is the electromagnetic kinetic term and $\\widetilde F_{\\mu\\nu}$ its dual, $a$ the axion field, $m_a$ the axion mass and $M_a$ is its coupling to the photon field.  The coupling of the photon and axion fields in this way means that a photon has a finite probability of mixing with its opposite polarisation and with the axion in the presence of an external magnetic field \\cite{raffelt}, such as the intergalactic magnetic field.   If the magnetic field is distributed randomly along the path of the light this will lead to a gradual reduction in luminosity until only 2/3 of the initial energy of the flux remains in the form of photons.  If the axion mass and coupling scales are chosen to be $m_a\\sim 10^{-16}$eV and $M_a\\sim 4\\times 10^{11}$GeV respectively \\cite{Csaki:2001yk}, the fractional amount of dimming will not vary much over the optical range where the supernovae are observed.  This might enable one to consider models where the extra energy component of the universe is not a cosmological constant but a fluid with a different equation of state, e.g. a string network $\\cite{bucher}$.  The first paper suggesting such a mixing broke the line of sight up into cells within each of which the magnetic field was oriented in a random direction \\cite{Csaki:2001yk}.  Later papers included the effect of a redshift varying magnetic field and the presence of intergalactic electrons \\cite{Deffayet:2001pc}\\cite{Mortsell:2002dd} (see also \\cite{Erlich:2001iq}).  If the electron density is assumed to be approximately uniformly distributed in the intergalactic medium with the value $\\tilde{n}_{e}=1.8\\cdot 10^{-7} \\rm cm^{-3}$ the oscillations will strongly depend on the energy of the photons leading to possible conflict with supernovae spectral observations \\cite{sndim}. However, as pointed out in \\cite{Csaki:2001jk}, the intergalactic medium is not uniform and many regions will have a electron density much below the mean density.  This, according to \\cite{Csaki:2001jk}, will render the energy dependence unobservable at present.  In this paper we study the effects of the variation of the free electron density in the plasma on the photon-axion mixing. Due to the large electrical conductivity of the intergalactic medium the magnetic field can be considered frozen into the plasma. The magnitude of the magnetic field will therefore depend on the electron density as $B\\propto n_{e}^{2/3}$. In regions with a low electron density the mixing probability will be suppressed because in these regions the magnetic field is reduced. Therefore, when both electron density and magnetic fields are allowed to fluctuate, the total mixing probability for a path along the line of sight is sensitive to regions where fluctuations have produced an electron density high above the mean density. It is therefore important to take fluctuations in both the electron density and the magnetic field into account when evaluating the total mixing probability.  Throughout this paper we have assumed that the intergalactic magnetic fields were created in the early universe, i.e. at much higher redshifts than any of the objects we will be considering.  There are some models where the intergalactic magnetic fields are produced at much later epochs.  For example it is possible that super massive black holes may be responsible for the amplification of magnetic fields in the intergalactic medium at late times \\cite{colgate}.  Different results would be obtained if one were to consider such models. ",
        "conclusions": "In this paper we have considered the dimming of light from high redshift supernovae due to mixing between photons and a pseudoscalar axion field in the intervening intergalactic plasma.  We have tried to simulate this process in a more detailed way than previous studies by taking a realistic density probability distribution for the electron density.  We have also included the effect of a fluctuating magnetic field due to it being frozen into the background plasma.  The presence of such an axion does not completely solve the problem of the missing $\\sim 70\\%$ of the energy of the universe.  Rather we aimed to show that such a mixing might make it possible to explain the missing dark energy of the universe via an alternative source of stress energy other than a cosmological constant.  As a candidate for this energy we considered the case of a string network with equation of state $\\omega_S=-1/3$ and calculated the combination of axion-photon coupling and magnetic field that would lead to enough dimming.  Assuming a primordial origin for the intergalactic magnetic fields we found no values of these parameters which could explain the dimming without violating the bound on the frequency dependence of this dimming already calculated by the observational supernovae groups.  However, the colour bound that we observed was only a factor of about 50\\% larger than that observed, not many orders of magnitude, and there is a small inherent error expected in our analysis due to deviations from our assumed behaviour of the variance parameter $\\sigma(z)$ in the density probability function (\\ref{Pdelta}) which will become greater at high redshifts.  We found that for $\\eta=1.5\\times 10^{-26}{\\rm cm}^{-1}$ we obtained nearly enough dimming and only marginally violated he colour bound.  We therefore have adopted this as our optimal value of $\\eta$.  One of the main conclusions of this work is therefore that a photon-axion coupling of the sort discussed in \\cite{Csaki:2001yk} leads to very particular predictions as to the redshift dependence of dimming at different frequencies.  If one considers an alternative origin for the intergalactic magnetic fields other than their primordial production it will be possible to change the bounds obtained in this paper since the correlation between magnetic field strength and electron density used in our calculations will be modified.  We also calculated the expected reduction in the polarisation of light from high redshift objects due to photon-axion mixing.  After a brief literature search we were unable to find any observations of high redshift AGN with large enough polarisations to be at odds with a dimming of $\\eta=1.5\\times 10^{-26}{\\rm cm}^{-1}$.  However, we find the frequency dependence of this effect to be so large that we believe such a dimming would lead to no observed sources at redshift $z\\sim 1.5$ with UV (100 nm) polarisation as high as $10\\%$.  Any distribution of magnetic fields which could lead to significant dimming of the optical flux coming from a distant object should also lead to a large reduction in any polarisation of the UV photons coming from that object.  This prediction is independent of the detailed history of cosmic magnetic fields.  It should be possible for this issue to be further investigated by making space based polarimetric observations."
    },
    "0207/astro-ph0207239_arXiv.txt": {
        "abstract": "{ A statistical analysis of a large data set of H$\\alpha$ flares comprising almost \\mbox{100\\,000} single events that occurred during the period January 1975 to December 1999 is presented. We analyzed the flares evolution steps, i.e. duration, rise times, decay times and event asymmetries. Moreover, these parameters characterizing the temporal behavior of flares, as well as the spatial distribution on the solar disk, i.e. \\mbox{N-S} and \\mbox{E-W} asymmetries, are analyzed in terms of their dependency on the solar cycle. The main results are: 1)~The duration, rise and decay times increase with increasing importance class. The increase is more pronounced for the decay times than for the rise times. The same relation is valid with regard to the brightness classes but in a weaker manner. 2)~The event asymmetry indices, which characterize the proportion of the decay to the rise time of an event, are predominantly positive ($\\approx 90\\%$). For about 50\\% of the events the decay time is even more than 4 times as long as the rise time. 3)~The event asymmetries increase with the importance class. 4)~The flare duration and decay times vary in phase with the solar cycle; the rise times do not. 5)~The event asymmetries do not reveal a distinct correlation with the solar cycle. However, they drop during times of solar minima, which can be explained by the shorter decay times found during minimum activity. 6)~There exists a significant \\mbox{N-S} asymmetry over longer periods, and the dominance of one hemisphere over the other can persist for more than one cycle. 7)~For certain cycles there may be evidence that the N-S asymmetry evolves with the solar cycle, but in general this is not the case. 8)~There exists a slight but significant \\mbox{E-W} asymmetry with a prolonged eastern excess. ",
        "introduction": "With the discovery of the first solar flare observed independently by R.C. Carrington and R. Hodgson on September 1, 1859, knowledge about these energetic phenomena on the Sun has steadily increased. The statistical investigations of the characteristics of solar H$\\alpha$ flares essentially started in the 1930s, when a worldwide surveillance of the Sun based on Hale's spectrohelioscope was established (Cliver 1995). Since then many papers have been published analyzing different statistical aspects of solar flares. However, in most of the papers only single aspects are investigated and/or the data set comprises a quite limited period. Moreover, there exists no recent paper extensively studying the statistical properties of solar H$\\alpha$ flares.  In the present paper we make use of the substantial data collection of solar H$\\alpha$ flares in the Solar Geophysical Data (SGD). We selected the period January~1975 (since then the H$\\alpha$ flares are listed with the same content and format) to December~1999. With this selection we have a homogeneous data set comprising almost \\mbox{100\\,000} single flare events, which provides a significant statistical basis. Furthermore, as the selected period entirely covers two solar cycles, 21 and 22, and the rising phase of solar cycle~23 until the end of year 1999, the data set also enables us to analyze dependency of the flare characteristics on the solar cycle.  The paper is structured as follows. In Sect.~2 a characterization of the data set is given. Sect.~3 describes the applied methods. In Sect.~4 the results are presented and discussed, comprising a statistical analysis of the temporal flare parameters combining the data of the overall period, such as flare duration, rise times, decay times (Sect.~4.1) and event asymmetries (Sect.~4.2). Also, we analyze the above-mentioned temporal and spatial flare characteristics, namely the N-S and E-W asymmetry, and its dependency on the solar cycle (Sect.~4.3). Sect.~5 contains a summary of the main results and the conclusions. ",
        "conclusions": "\\label{discussion}  \\subsection{Temporal parameters}  The main results of the analysis regarding temporal flare parameters, as duration, rise times, decay times and event asymmetries are summarized in the following. 1)~On average, the duration, rise and decay times of flares increase with increasing importance class. The increase is more pronounced for the decay times (factor 4--5 between subflares and flares of importance~$>$1) than for the rise times (factor 2--3). The same relation holds for the brightness classes but in a weaker manner. 2)~The event asymmetries, which characterize the proportion of the decay to the rise time of a flare, are predominantly positive ($ \\approx 90\\%$). For more than 50\\% of all flares the decay phase is even more than 4 times as long as the rising phase. 3)~On average, the event asymmetries increase with the importance class. 4)~The duration changes in phase with the solar cycle, i.e. on average the flare duration is longer during periods of maximum activity than during solar minima. Since the rise times do not reveal a distinct correlation with the solar cycle but the decay times do, the variations of the duration in accordance with the solar cycle are due to the variations of the decay times. 5)~The event asymmetries do not reveal a distinct correlation with the solar cycle. However, they decrease during solar minima.  The facts that on the average the flare duration increases with the importance class and the rise times are shorter than the decay times are reported in a number of previous papers (see the references cited in Table~\\ref{previous results}). However, the new outcome of the present analysis is, on the one hand, that the increase of the duration with the importance class in particular results from the increase of the decay times, which is significantly more pronounced than the increase of the rise times. On the other hand, by the concept of the event asymmetry, we were able to give a quantitative description of the asymmetry in the flare development, which revealed that on average the asymmetry also increases with the importance class. Both results suggest that, with respect to the temporal behavior, the cooling phase of the H$\\alpha$ flare is more strongly affected by the flare size than the phase of heating-up the chromospheric plasma at the flare site.  Furthermore, a significant change in the duration and decay times with the solar cycle was found. On average, during solar maximum the decay times are larger than during solar minimum, with $t_{\\rm decay}^{\\rm max} \\approx 1.5 \\cdot t_{\\rm decay}^{\\rm min}$. The rise times do not reveal a significant variation in accordance with the cycle. The combination of both facts can also account for the drop in the event asymmetries found during solar minima. These results suggest that the change in the flare duration is mainly caused by the variations of the decay times, giving further evidence that the flare cooling phase is more sensitive to changes in the physical conditions of the chromospheric plasma than the rising phase.   \\subsection{Spatial distributions}  The main outcomes of the present analysis regarding the spatial distribution of flares over the different hemispheres are: 1)~There exists a significant \\mbox{N-S} asymmetry over longer periods, and the dominance of one hemisphere over the other can persist for more than one cycle. 2)~For certain cycles there may be evidence that the N-S asymmetry evolves in coincidence with the solar cycle, but in general this is not the case. 3)~There exists a slight but significant \\mbox{E-W} asymmetry with a prolonged eastern excess during solar cycles~21 and~22. Combining the results obtained by previous authors and the present paper, possibly an eastern excess of solar flares existed during solar cycles~17 to~22.  The existence of a \\mbox{N-S} asymmetry is generally accepted but still not definitely interpreted. One possible explanation of the N-S asymmetry of solar activity phenomena is that a time difference in the development of solar activity on the northern and the southern hemisphere exists (e.g., Tritakis et al. 1997). However, in this case the N-S asymmetry is expected to evolve in coincidence with the solar cycle, which is rather ambiguous. Another explanation utilizes the concept of ``superactive regions\", which are large, complex, active regions containing sunspots (Bai 1987, 1988). Such superactive regions produce the majority of solar flare events and appear preferentially in certain areas of the Sun, so-called ``active zones\". As shown by Bai et al. (1988) in the past, active zones were present on the Sun, which persisted for several solar cycles. In that frame, the \\mbox{N-S} asymmetry can be attributed to the existence of active zones in the northern and southern hemispheres, which can persist over long periods.  Also the fact that flares are not uniformly spread in heliographic longitude over the solar disk (note that this non-uniformity has a different meaning than the E-W asymmetry), pointed out by Heras et al. (1990) and Li et al. (1998), can be understood in the framework of the active zones concept. The longitudes, at which the active zones are located, are more flare-productive than other longitude ranges. The active zones induce a pronounced and non-random flare activity, which is superimposed onto an episodic and random flare activity coming from the other (less active) regions of the solar disk (Heras et al. 1990).  However, the idea of active zones cannot account for an \\mbox{E-W} asymmetry persisting over time scales larger than those of the solar rotation.  The first report of an E-W asymmetry in solar activity was given by Maunder (1907), who found that the total spot area and total number of spot groups were larger in the eastern than in the western hemisphere. A possible interpretation of this effect was given by Minnaert (1946). A forward tilt of the vertical sunspot axis causes a ``physical\" foreshortening of the spots, which acts more strongly on spots on the western than on those on the eastern hemisphere. However, it cannot be seen how such an effect could also account for an eastern excess in the flare occurrence rate. In recent papers (Mavromichalaki et al. 1994, Tritakis et al. 1997) it has also been found that the emission of the eastern hemisphere of the corona systematically predominates. However, the coronal \\mbox{E-W} asymmetry as well as the \\mbox{E-W} asymmetry of solar flares are still unexplained and controversial issues."
    },
    "0207/astro-ph0207596_arXiv.txt": {
        "abstract": "We review recent observational and theoretical results concerning the presence of actinide nuclei on the surfaces of old halo stars and their use as an age determinant. We present model calculations which show that the observed universality of abundances for $56 < Z < 75$ elements in these stars does not necessarily imply a unique astrophysical site for the $r$-process. Neither does it imply a universality of abundances of nuclei outside of this range. In particular, we show that a variety of astrophysical $r$-process models can be constructed which reproduce the same observed universal $r$-process curve for $56 < Z < 75$ nuclei, yet have vastly different abundances for $Z \\ge 75$ and possibly $Z < 56$ as well. This introduces an uncertainty into the  use of the Th/Eu chronometer as a means to estimate the ages of the metal deficient stars. We do find, however, that the U/Th ratio is a robust chronometer. This is because the initial production ratio of U to Th is almost independent of the astrophysical nucleosynthesis environment. The largest remaining uncertainties in the U/Th initial production ratio are due to  the input nuclear physics models. ",
        "introduction": "Rapid neutron-capture (the $r$-process) is responsible for producing about half of the elements heavier than iron.  It is believed to occur in an explosive stellar environment in which the neutron capture time scale is much shorter than typical beta-decay lifetimes near the line of stable nuclei. The nuclear reaction flow can then proceed through  extremely neutron-rich unstable nuclei. The fact that the heavy radioactive actinide nuclei, U and Th, are generated in the $r$-process is of particular interest. These nuclides have half lives [$t_{1/2}(^{238}$U)$= 4.47 \\times 10^9~y$, $t_{1/2}(^{232}$Th)$= 1.40 \\times 10^{10}~y$] which are comparable to the cosmic age. These chronometers have taken on renewed recent attention as their absorption lines  have been identified \\citep[e.g.][]{sneden,cayrel, honda03a, honda03b} in metal deficient stars.  The inferred abundances of Th and/or U can be used to estimate stellar ages. Metal deficient stars are believed to be the oldest stars in the Galaxy, and their surface abundances have probably not changed (except for radioactive decay) since these stars were formed. Moreover, their age can be regarded as the Galactic age and a lower limit to the cosmic age. As a chronometer, this method is particularly appealing since it avoids the usual Galactic chemical evolution model dependence \\citep{meyerschramm} associated with Solar-System radio cosmochronometry. The Th or U on the surface of an old low-metallicity halo star were probably generated in a single nucleosynthesis event. Hence, the surface abundance of radioactive element $Y_{r}$ (r=Th or U) is given to a very good approximation by \\begin{equation} Y_r (\\Delta T)=Y_r(0)exp(-\\Delta T/\\tau_r) \\end{equation} where $\\Delta T$ is time since nucleosynthesis and $\\tau_r$ is the mean life of the r-element, {\\it i.e.} $\\tau_r=t_{1/2}/ln 2$. For each radioactive element, one can solve Eq. (1) to find $\\Delta T$. It is best to utilize abundance ratios relative to an element with a nearby absorption feature (e.g. Eu). Therefore, $\\Delta T$ is usually given by \\begin{equation} \\Delta T = 46.7 (\\log ({\\rm Th/Eu})_{0}-\\log ({\\rm Th/Eu})_{\\rm T})~~, \\end{equation} \\begin{equation} \\Delta T= 21.8 (\\log({\\rm U/Th})_{0} - \\log({\\rm U/Th})_{\\rm T})~~, \\end{equation} in units of Gyr, where the index $0$ denotes the initial production ratio, while the index $T$ refers to the presently observed value. (In these equations we denote $Y_{Th}=Th$, $Y_{Eu}=Eu$, and $Y_U=U$.) The only uncertainties are, therefore, those in the determination of the present stellar abundances themselves, and those due to uncertainties in model estimates of the initial production ratios. [That is, as long as the produced actinide nuclei have not passed through stellar CNO burning where they might be destroyed by photo-induced fission \\citep{mmd89}.]  In view of the significance of this independent measure of the Galactic age, it becomes important to scrutinize and quantify these remaining uncertainties as much as possible.  In this paper we are primarily concerned with the astrophysical uncertainties in the initial production abundances. We show that there is indeed considerable uncertainty in using only a single Th or U radiochronometer, even when the universally observed Solar-System $r$-process abundances are well reproduced for lower-mass nuclei.  We also establish that the Th/U chronometer is quite robust and somewhat independent of astrophysical model uncertainties.  For some time \\citep[cf.][]{truran,mathews90} r-process elements in metal deficient stars have been interpreted as evidence for a universal $r$-process abundance distribution in operation in the early Galaxy. In particular, more recent observations \\citep{sneden,Sned98,sneden2,johnson} all show similar abundance distributions for $Z > 56$ elements. This feature is often referred to as the ``universality'' of the $r$-process. Hence, it is generally believed that, at least for $Z > 56$ elements, the astrophysical site and associated yields of r-process nucleosynthesis are unique. Based upon this assumption, \\citet{sneden} estimated the ages of several stars using the ratio of Th/Eu at the time of formation of the Solar System as the initial production ratio [even though Solar-System material has experienced multiple supernovae before its formation and has experienced Th decay]. The analyses of these stars all indicated similar present ages of about 14 Gyr $\\pm$ 4.  There are, however, some reasons to question the assumption of a universal $r$-process abundance curve.  The material out of which these metal poor stars were formed is likely to have experienced only one or two supernovae before incorporation into the star. Depending upon which particular progenitor supernova was in operation, there might be substantially different abundance distribution curves for these stars, compared to the ensemble average represented in Solar-System material \\citep[cf.][]{Ishi99}. Moreover, quite recently, \\citet{cayrel} have reported the observation of peculiar r-process elements in the metal deficient star CS31082-001. This is the first star for which a uranium line was also detected. This star has the strange feature, however, that the ratio of Th to Eu is greater than that of Solar material. This would imply (on the bases of Th/Eu) that this star is younger than the Sun. This seems unlikely in view of its low metallicity, [Fe/H] $\\sim -2.9$. Furthermore, the Th/Eu age is in contradiction with the U/Th age for this star which is 12 $\\pm$ 3 Gyr. Further evidence of Th/Eu uncertainty is apparent in the recent data of \\citet{honda02} and \\citet{honda03a,honda03b}. They have reported on r-process element abundance distributions in two other metal-deficient stars. These stars also show different abundance distribution patterns for $Z \\ge 75$.  Even the stars studied in Sneden et al.(1996) and Sneden et al.(2000) may show deviations from a universal $r$-process distribution for lighter Z $<$ 56 nuclei. For example Sneden et al.(2000) noted divergence from the solar r-process for a few elements in CS22892-052. \\citet{cowan} also noted some divergence in the star BD +17 3248, but those data are somewhat uncertain. This, together with meteoritic evidence, has been taken as an indication \\citep{qian} that two different $r$-process environments could be in operation. All in all, the observational data  seem to indicate that the universality of r-process elements may be broken for $Z \\ge 75$ elements and possibly Z $\\leq$ 56 as well.  In this paper, we wish to clarify the astrophysical model dependence for these stellar r-process chronometers. Several previous investigations \\citep{gor,wanajo02,sha,q02,otsuki03} of the uncertainties in actinide chronometers can be found in the literature. The paper by \\citet{gor} and \\citet{sha} were primarily directed toward understanding the nuclear uncertainties. For example, \\citet{gor} considered 32 different models for $r$-process actinide production based upon different nuclear physics input for fission barriers, nuclear masses, beta-decay rates, etc. In \\citet{sha}, it was found that the Pb abundance was useful to constrain nuclear models. These calculations, however, were made in the context of schematic ``canonical event'' models which were constrained to reproduce Solar-System $r$-process abundances. Hence, universality for all elements up to Z=82 was imposed. In another related work, \\citet{wanajo02} have attempted to clarify the age of CS31082-001 in the context of a specific ``neutrino-driven wind'' model, but with only electron fraction and the outer boundary temperature as free parameters. \\citet{q02} even shows on the basis of the astrophysical observations that the production of the elements with $A>130$ is not robust. All of these works have demonstrated some uncertainties in the Th/Eu chronometer, while better results are obtained using the U/ Th chronometer.  The present work differs from the previous works in two important ways: 1) Unlike \\citet{gor} and \\citet{sha}, our focus is on the astrophysical rather than nuclear uncertainties; 2) Rather than to focus on a particular wind model as in \\citet{wanajo02} and \\citet{otsuki03}, we consider wide range of plausible astrophysical models and parameters all of which reproduce the universality in the abundances for 56 $<$ Z $<$ 75 noted in observations. We also explicitly considerneutrino interaction effects.  From the observations and our theoretical calculations, we conclude that the astrophysical site for r-process nucleosynthesis is probably not unique, even though there is an observed universality of abundances. That is, a variety of astrophysical models can be constructed which reproduce the same apparent universal abundance curve for elements with $56 < Z < 75$, but which give large variations in abundances for elements with $Z \\ge 75$ and/or $Z \\le 56$. Hence, it is dangerous to use the Th/Eu chronometer to estimate the age of metal deficient stars. Moreover, the production ratio of the Th/Eu chronometer is strongly dependent upon  the nucleosynthesis environment. On the other hand, we demonstrate that the U/Th production ratio is almost independent of the nucleosynthesis environment, confirming that this ratio is robust as a  chronometer.  In what follows, we first review in more detail the recent observational results for metal deficient halo stars in section 2. Results of theoretical calculations are shown in section 3. In section 4, we will summarize the viability of the nuclear cosmochronometers. ",
        "conclusions": "Our theoretical calculations indicate that the coincidence of the observed abundance distribution for $56<Z<75$ elements with the Solar r-process abundances does not necessarily  mean  that all $r$-process events occur in the same universal environment. Moreover, different abundance distributions for $Z \\ge 75 $ and $Z \\le 56$ elements are produced even when the universal $56 < Z < 75 $ abundances are reproduced. Hence, it is not possible to predict the initial production ratio of Th to Eu in each metal deficient star without somehow independently identifying each production environment. We therefore conclude that there is substantial uncertainty when using the Th/Eu chronometer to estimate the ages of metal deficient stars.  We note that the reason for the uncertainty is that in realistic (e.g. SN wind) r-process models, one tends to run out of neutrons as actinide elements are being formed. It might be possible, however, that an r-process environment could occur in which there were abundant neutrons present as the actinide nuclei are synthesized. In this case, the r-process could proceed to the fission termination and recycling. If most stars experienced this, then the uncertainty in the Th/Eu chronometer would be reduced and the CS31082-001 star would be an interesting exception. Nevertheless, since it is difficult to produce fission recycling in realistic (e.g. wind) models, we caution against this assumption. Another possibility to reduce the Th/Eu uncertainty is that a combination of light-element (Z $\\le$ 56) abundances and Pb-region abundances \\citep{sha} might help to identify the environment. We also note that variations of the light-element abundances are quite natural in neutrino-driven wind models even when a universality of heavier elements is imposed. Hence, one may not need to propose two different $r$-process sites \\citep{qian} to account for the variations of light elements relative to the universal heavy nuclei. Although two sites are probably still necessary to explain the meteoritic evidence. To some extent, however, the light-element abundances may simply reflect the star-to-star variations in a single neutrino-driven wind scenario for the $r$-process.  On the other hand, we find that the U/Th chronometer is a robust means to estimate the cosmic age. From our theoretical calculations, we have shown that the U/Th initial production ratio is almost independent of the nucleosynthesis environment as suggested by \\citet{gor} and \\citet{wanajo02}. Indeed, the present work vindicates the conclusions of those papers for wider range of astrophysical environments. Hence, using this chronometer, one can estimate the age of metal deficient stars with little correction for the astrophysical uncertainties. Unfortunately, some uncertainty due to the nuclear physics still remains \\citep[e.g.][]{sha}. Further quantitative studies of the uncertainties due to various nuclear physics input such as the mass formula, beta-decay rates, alpha-decay rates, $\\beta$-delayed fission and neutron emission are strongly desired. Efforts along this line are currently underway.  Regarding the universal abundance curve, we have shown that there is a dependence of the $56<Z<70$ abundances on the dynamical time scale. We also note that there is a dependence of these abundances on the detailed nuclear physics models. The hope is that with sufficiently precise observed abundances one might be able to both identify the true dynamical timescale and the correct nuclear physics models for the r-process. More accurate observational data are likely to be available in the future due to the development of large aperture telescopes and high resolution spectroscopy. Clearly, more observations of $r$-process abundances in metal poor stars are desired.   This work has been supported in part by Grants-in-Aid for Scientific Research (10640236, 10044103, 11127220, 12047233) of the Ministry of Education, Science, Sports, and Culture of Japan. Work at the University of Notre Dame was supported by the U.S. Department of Energy under Nuclear Theory Grant DE-FG02-95-ER40934. One of the authors (KO) wishes to acknowledge support from the Japanese Society for the Promotion of Science. KO and TK also thank the Institute for Nuclear Theory at the University of Washington for their hospitality and the Department of Energy for their partial support during the completion of this work."
    },
    "0207/astro-ph0207243_arXiv.txt": {
        "abstract": "The ESO Distant Cluster Survey (EDisCS) is a photometric and spectroscopic study of the galaxy cluster population at two epochs, $z \\simeq 0.5$ and $z\\simeq 0.8$, drawn from the Las Campanas Distant Cluster Survey (LCDCS).  We report results from the initial candidate confirmation stage of the program and use these results to probe the properties of the LCDCS. Of the 30 candidates targeted, we find statistically significant overdensities of red galaxies near 28.  Of the ten additional candidates serendipitously observed within the fields of the targeted 30, we detect red galaxy overdensities near six.  We test the robustness of the published LCDCS estimated redshifts to misidentification of the brighest cluster galaxy (BCG) in the survey data, and measure the spatial alignment of the published cluster coordinates, the peak red galaxy overdensity, and the brightest cluster galaxy. We conclude that for LCDCS clusters out to $z \\sim 0.8$ , 1) the LCDCS coordinates agree with the centroid of the red galaxy overdensity to within 25\\arcsec ($\\sim 150 h^{-1}$ kpc) for 34 out of 37 candidates with $3\\sigma$ galaxy overdensities, 2) BCGs are typically coincident with the centroid of the red galaxy population to within a projected separation of 200$h^{-1}$ kpc (32 out of 34 confirmed candidates), 3) the red galaxy population is strongly concentrated, and 4) the misidentification of the BCG in the LCDCS causes a redshift error $>$ 0.1 in 15-20\\% of the LCDCS candidates.  These findings together help explain the success of the surface brightness fluctuations detection method. ",
        "introduction": "Observations of the distant galaxy cluster population are being driven by a new generation of catalogs that provide statistically significant samples of hundreds to thousands of candidate clusters between redshifts 0.5 and 1.  These catalogs can either serve as large statistical samples without appealing to any further observations, for example to measure the cluster-cluster correlation function \\citep{gon2002}, or as input for follow-up studies of selected subsamples.  The European Southern Observatory Distant Cluster Survey (EDisCS) is a detailed follow-up study of 20 clusters, 10 at $z \\simeq 0.5$ and 10 at $z \\simeq 0.8$, drawn from the 1073 candidate clusters cataloged by the Las Campanas Distant Cluster Survey \\citep[LCDCS;][]{thesis,gon2001}.  This paper describes results from the preliminary effort to confirm the set of cluster candidates that will be the focus of the more extensive observations of the EDisCS.  The value of a cluster catalog is greatly enhanced for any application if the catalog includes measurements of the redshift and mass of each candidate cluster.  Due to the size of recent catalogs, it is impractical to obtain {\\it spectroscopic} redshifts or masses for a significant fraction of the catalog. Most catalogs now provide an estimate of these properties drawn solely from the survey data \\citep[see][]{post96,gon2001}.  Superior survey data, for example deeper images or multiple colors, should improve the reliability of the estimated parameters, but they decrease the observing efficiency. The optimum balance between the fidelity of the cluster catalog and observational efficiency is not evident, and will depend on the scientific aims. \\citet{gon2001} provided a catalog from what is arguably the most observationally efficient method (10 nights at a 1m telescope produced a catalog of $\\sim$ 1000 cluster candidates out to $z \\sim 1$ over an area covering 130 sq. degrees), but which might in turn provide the least robust estimates of the cluster redshifts and masses, and which, even more importantly for some potential uses of the catalogs, may include a larger fraction of false detections. Using observations in multiple filters that are $\\sim$ 3 magnitudes deeper than original survey data, we examine whether the false positive rate quoted originally for the LCDCS is valid and whether the LCDCS cluster coordinates and estimated redshifts are confirmed using deep, multicolor data.  The LCDCS generated a catalog of concentrations of photons on the sky rather than galaxies \\citep{dal1995,zar1997,gon2001}. Significant fluctuations in the background sky are classified into various categories, including high redshift clusters.  The cluster redshift is estimated using the magnitude of the brightest galaxy near the surface brightness fluctuation, which is presumed to be the brightest cluster galaxy (BCG).  The redshift-magnitude relationship for BCGs is calibrated using spectroscopy of a sample of $\\sim$ 20 clusters.  The cluster mass is estimated using the peak brightness of the convolved surface brightness map, calibrated using a sample of $\\sim 10$ clusters with X-ray temperature and velocity dispersion measurements. The uncertainties in each of these estimators, and of the false positive and negative rates, are discussed by \\citet{gon2001} in their presentation of the LCDCS.  Our examination of these issues here utilizes multifilter images of 30 targeted fields, which contain 40 LCDCS candidate clusters, obtained with ESO's Very Large Telescope (VLT) as part of the initial stage of the EDisCS.  The format of this paper is as follows. In \\S 2 we present the data utilized in this analysis and details of the reduction procedure.  We then examine the two-color photometry in \\S 3 to confirm or reject the cluster candidates observed by EDisCS and test whether the fractional contamination is consistent with that given by \\citet{gon2001}. In \\S 4 we test the robustness of the estimated redshifts quoted for the LCDCS, which are based upon the magnitude of the brightest cluster galaxy, with particular emphasis on the potential problem of BCG misidentification.  In this section we also quantify the offsets between the LCDCS coordinates, the locations of the brightest cluster galaxies, and the peak of the projected galaxy distribution. Next, we briefly comment upon the LCDCS mass estimates in \\S 5 and compare the LCDCS surface brightness with other observable quantities.  Finally, a summary of the results and brief discussion of forthcoming work are presented in \\S 5.  For all physical distances in this paper we assume a flat, $\\Lambda$CDM cosmology with $\\Omega_0$=0.3. ",
        "conclusions": "In this paper we present results from the initial phase of the ESO Distant Cluster Survey (EDisCS), using VLT imaging to better characterize the Las Campanas Distant Cluster Survey catalog. We first use smooth density maps of the color-selected galaxy distribution to confirm cluster candidates. We find that 93\\% (28/30) of the EDisCS targets are coincident with statistically significant overdensities of red galaxies, as are 60\\% (6/10) of serendipitously imaged LCDCS candidates. The latter number is consistent with the contamination rate published by \\citet{gon2001} for a randomly-selected subsample of the LCDCS with the same redshift distribution. In addition to confirming a set of promising LCDCS candidates for further study, we also use the photometry to identify the brightest cluster galaxies and use these identifications to test the robustness of the estimated redshifts published in the LCDCS catalog.  We find that misidentification leads to redshift errors $\\Delta z>0.1$ in 6/34 cases (18\\%), which is also consistent with predictions from \\citet{gon2001}.  In addition, we find that the surface brightness detection technique appears to be slightly more sensitive to the overdensity of unresolved cluster galaxies than it is to diffuse emission from the extended halos of brightest cluster galaxies, indicating that the redshifts are most likely to be overestimated for dynamically unrelaxed systems in which the BCG and peak galaxy overdensity are not aligned.  Of more general interest are our findings that 1) the distribution of red cluster galaxies is generally regular and highly centrally concentrated out to $z \\sim 0.8$, and 2) that the BCGs are also found near the concentrations of red galaxies, suggesting that the cores of clusters out to $z \\sim 0.8$ are typically dynamically relaxed. The use of photometric and spectroscopic redshifts obtained as part of EDisCS will help establish whether these conclusions hold once cluster members are identified. The above results verify that the surface brightness fluctuation technique proposed by \\citet{dal1995} and employed by \\citet{gon2001} is an effective method of identifying distant clusters, and demonstrate that the utility of the LCDCS catalog redshifts is not seriously compromised by misidentification of the BCG. Upcoming spectroscopy for the EDisCS will improve upon this analysis by directly testing the robustness of the redshift estimates, as well as the LCDCS predictions for the cluster velocity dispersions. These data define the sample of clusters that will comprise the EDisCS."
    },
    "0207/astro-ph0207419_arXiv.txt": {
        "abstract": "We show that  there are simple one dimensional  problems for which the MHD   code,  ZEUS,  generates   significant  errors,   whereas  upwind conservative schemes perform very well on these problems. ",
        "introduction": "ZEUS  is a  freely  available MHD  code  that is  widely  used by  the Astrophysical  community.   Although Stone  \\&  Norman (1992a,b)  give results for  the Sod  problem (Sod 1978)  and its MHD  equivalent, the Brio and  Wu problem (Brio \\& Wu  1988), ZEUS does not  appear to have been  tested  on  a wide  range  of  Riemann  problems such  as  those described in e.g.  Dai \\& Woodward (1994), Ryu \\& Jones (1995), Falle, Komissarov \\& Joarder (1998)and  Balsara (1998).  Since  ZEUS  is  neither  upwind  for all  characteristic  fields  nor conservative, we  might expect it  to perform significantly  less well than upwind conservative codes (e.g.  Brio \\& Wu 1988; Dai \\& Woodward 1994; Ryu  \\& Jones 1995;  Falle, Komissarov \\& Joarder  1998; Balsara 1998; Powell  et al. 1999).  As we  shall see, this is  indeed true in the sense  that there are  a number of  simple problems for  which the ZEUS solution contains significant errors that are absent in solutions calculated with an upwind conservative scheme. ",
        "conclusions": "It is  evident from these  results that, ZEUS  can be made  just about acceptable for pure gas dynamics if the linear artificial viscosity is multiplied by the smallest local Courant number since the shock errors are small in this case.  However, it is not satisfactory for adiabatic MHD, at least in its present  form.  The shock errors do not occur for an isothermal equation of state, but, since the rarefaction shocks do, ZEUS is also not reliable for  isothermal MHD. It is possible that the rarefaction  shocks in  MHD  waves  can be  removed  without using  an excessive  linear   artificial  viscosity   by  the  addition   of  an appropriate linear artificial resistivity. The shock errors might also be  reduced by  advecting the  total energy  rather than  the internal energy.   However,  even with  such  improvements,  the  low order  of accuracy  makes ZEUS very  inefficient compared  with a  modern upwind scheme.  This should not be taken to mean that conservative upwind codes are in any sense  perfect.  For  example, it is  necessary to  introduce some extra dissipation in  the Riemann solver to remove  the serious errors discussed  by Quirk  (1994)  and some  desirable  properties, such  as strict conservation, may have to be sacrificed in order to satisfy the constraint $\\nabla  \\cdot {\\bf  B} = 0$  in multidimensional  MHD (see e.g. Powell et al. 1999; Balsara 2001).  These results  obviously have implications for the  reliability of the numerous calculations in the  literature that have used ZEUS. Although these effects are  likely to be present in  many cases, the associated errors are not necessarily so  serious as to completely invalidate the calculations. Whether  or not they make any  qualitative difference in any  particular  case  can  only  be  decided  either  by  a  thorough examination  of the results  to see  whether any  of these  errors are present, or by repeating the  calculations using a modern code.  These calculations were performed with the version of Zeus2d available from the NCSA  website, but, since all versions of  ZEUS appear to use the same algorithms,  the results should not depend  on the particular version.   It is also  worth pointing  out that  although we  used the scheme  described by  Falle,  Komissarov and  Joarder (1998),  similar results would probably have been obtained with any modern upwind code.  The  author would  like  to thank  both  the editor  and an  anonymous referee for a number of helpful comments on the original version."
    },
    "0207/astro-ph0207133_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 }} {The photometric data collected during 2001 season OGLE-III planetary/low luminosity object transit campaign were reanalyzed with the new transit search technique -- the BLS method by Kovacs, Zucker and Mazeh. In addition to all presented in our original paper transits, additional 13 objects with transiting low-luminosity companions were discovered. We present here a supplement to our original catalog -- the photometric data, light curves and finding charts of all 13 new objects.  The model fits to the transit light curves indicate that a few new objects may be Jupiter-sized (${R<1.6~R_{\\rm Jup}}$). OGLE-TR-56 is a particularly interesting case. Its transit has only 13~mmag depth, short duration and a period of 1.21190 days. Model fit indicates that the companion may be Saturn-sized if the passage were central.  Spectroscopic follow-up observations are encouraged for final classification of the transiting objects as planets, brown dwarfs or late M-type dwarf stars.  We also provide the most recent ephemerides of other most promising planetary transits: OGLE-TR-10 and OGLE-TR-40 based on observations collected in June 2002.  All photometric data are available to the astronomical community from the OGLE Internet archive.} ",
        "introduction": "In the paper by Udalski \\etal (2002), the OGLE photometric survey presented results of the pilot campaign aiming at discovery of small depth transits caused by planets or low-luminosity small objects like brown dwarfs or late type M-dwarfs. The campaign turned out to be very successful -- 46 objects with transiting objects causing small (${<0.08}$~mag) box-shaped drop of brightness were detected. Simple transit model assuming completely dark transiting object and limb darkening of the primary star was fitted to the light curves of transits, indicating small sizes of transiting companions. In several cases the results indicated Jupiter-sized objects (${R<1.6~R_{\\rm Jup}}$).  Unfortunately, the  photometry alone cannot unambiguously distinguish between Jupiter size planets and other low-luminosity objects: brown dwarfs and late type M dwarfs. All of them have radii of the order of 0.1--0.2~\\RS\\ (${1{-}2~R_{\\rm Jup}}$). Thus, a spectroscopic follow-up and a measurement of the radial velocity amplitude of the stars is needed to determine the masses of transiting companions and final classification.  Large number of discovered objects with transiting companions makes the transit method of planetary search very attractive as the efficiency of search can be much higher than that of standard spectroscopic searches. Moreover, non-planetary companions like brown dwarfs or faint late M-type dwarf are also very interesting astronomically and their parameters are poorly known. It should be noted that photometry of transits combined with precise spectroscopy is the only method of unambiguous determination of all basic parameters like dimensions and masses.  Recently, Kovacs, Zucker and Mazeh (2002) proposed a new method of analysis of large photometric datasets for transit detection: the Box-fitting Least Squares (BLS) algorithm. Encouraged by promising simulations on artificial datasets presented by Kovacs \\etal (2002) we ran the method on photometric data of all objects presented in Udalski \\etal (2002). All of them were easily found by the BLS method. Encouraged even more, we decided to run the algorithm on photometric data of all ${\\approx 52~000}$ Galactic disk stars selected for transit search. Results were very impressive -- 13 additional objects with transiting low-luminosity companions were detected by the BLS method, increasing the total number of transit objects detected in the 2001 OGLE-III campaign data to 59.  This paper is a supplement to the original paper presenting results of the 2001 OGLE-III planetary/low-luminosity object transit campaign (Udalski \\etal 2002). We show here photometry, light curves and finding charts as well as results of preliminary analysis of transit light curves for all 13 new objects detected with the BLS method. Photometry of new OGLE-III stars  with transits is available to the astronomical community from the OGLE Internet archive. ",
        "conclusions": "\\vspace*{12pt} The BLS method of transit detection (Kovacs \\etal 2002) proved to be a very effective and efficient tool for analysis of huge photometric datasets for periodic transits. Application of the method to the real photometric data collected during the 2001 OGLE-III transit campaign allowed to increase significantly the sample of low-luminosity transit objects, which is now totaling at 59. It is worth noting that the BLS method detects transit with smaller depth -- the smallest depth in the case of the OGLE-TR-56 transit is only 13 mmag -- and also transits in more noisy data (fainter objects).  For all new transit objects we performed an analysis of transit light curve similar as in our original paper (Udalski \\etal 2002). We estimated sizes of transiting objects by modeling their light curves assuming a completely dark companion and using formulae provided by Sackett (1999). These were numerically integrated at the appropriate phase points to produce model light curves. Details can be found in Udalski \\etal (2002).  The smallest size of the star and its companion is obtained when the transit is central, \\ie for the inclination ${i=90\\arcd}$. Non-central passage (smaller $i$) requires a larger size of both, the star and its companion. Table~2 lists the minimum size of the transiting companion and the corresponding size of the star (assuming ${M_s=1~\\MS}$) to provide information on the expected dimensions of the transiting objects. It should be remembered that semi-major axis of the orbit scales as $ M^{1/3}\\!$, and therefore the same scaling is appropriate for the sizes of stars and companions listed in Table~2. In the close-up windows in the Appendix we show model light curves for central transit (${i=90\\arcd}$) drawn with a continuous bold line.  \\MakeTable{lcc}{12.5cm}{Dimensions of stars and companions for central passage ($M_s=1~\\MS$)} { \\hline \\noalign{\\vskip3pt} Name       & $R_s$   & $R_c$ \\\\ &[\\RS]&[\\RS]\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} OGLE-TR-47 &    1.95 &    0.23 \\\\ OGLE-TR-48 &    1.21 &    0.16 \\\\ OGLE-TR-49 &    0.97 &    0.15 \\\\ OGLE-TR-50 &    1.16 &    0.22 \\\\ OGLE-TR-51 &    0.85 &    0.13 \\\\ OGLE-TR-52 &    1.27 &    0.20 \\\\ OGLE-TR-53 &    0.76 &    0.12 \\\\ OGLE-TR-54 &    1.38 &    0.19 \\\\ OGLE-TR-55 &    0.93 &    0.11 \\\\ OGLE-TR-56 &    0.80 &    0.08 \\\\ OGLE-TR-57 &    1.85 &    0.26 \\\\ OGLE-TR-58 &    1.75 &    0.26 \\\\ OGLE-TR-59 &    1.15 &    0.16 \\\\ \\hline }  Table~2 indicates that in several cases the new transiting companions are of Jupiter size (${<1.6~R_{\\rm Jup}}$). They may be planets, or brown dwarfs or low-luminosity M-type stars. The spectroscopic follow-up will provide necessary clarification.  The most intriguing object within this group is OGLE-TR-56. Table~2 indicates that the size of the transiting companion is smaller than the Saturn size if its transit is central. However, even for inclination as low as ${i=83\\arcd}$, its radius is still equal to one Jupiter radius. This object is certainly very small. The color of the primary star indicates a star of smaller size than the Sun consistent with Table~2. If confirmed spectroscopically to be a planet, OGLE-TR-56 would have the shortest orbital period (1.21190~days) and would be one of the smallest extrasolar planets.  Alternatively the transiting object in OGLE-TR-56 might be a brown dwarf. Recently, Santos \\etal (2002) reported possible detection of a brown dwarf around HD~41004 with a period of only 1.3 days. Also, possible detection of a brown dwarf in the short period Hyades binary system RHy~403 was reported by Reid and Mahoney (2000).  OGLE-TR-56 could be another example indicating that the so called ``brown dwarf desert'' -- lack of brown dwarfs with short periods on small orbits, is not that empty as generally thought.  The least interesting alternative would be a blend of a regular Galactic bulge eclipsing star with a bright disk star which could mimic the low depth transit. However, this case is not very likely -- the shape of the transit light curve, namely short time of descending and ascending branches of the transit compared to long phase of totality, makes this explanation less probable.  It is worth adding here that two other most promising planetary transit objects: OGLE-TR-40 and OGLE-TR-10, discovered in the original paper (Udalski \\etal 2002), were observed for expected transits on a few nights in June 2002. The main goal of these observations was the improvement of ephemerides that due to short time baseline in 2001 could be inaccurate after longer period of time. For both objects a few transits were covered and the new best ephemerides, based now on more than one hundred orbital cycles, are: \\vspace*{-7pt} $${\\rm HJD}_{\\rm min} = 2452060.03667 + 3.43079\\cdot E \\eqno({\\rm OGLE-TR-40})$$ \\vspace*{-15pt} $${\\rm HJD}_{\\rm min} = 2452070.21900 + 3.10140\\cdot E \\eqno({\\rm OGLE-TR-10})$$ The photometric data of new OGLE-III transit objects  presented in this paper are available in the electronic form from the OGLE archive: \\vskip3pt \\centerline{\\it http://www.astrouw.edu.pl/\\~{}ogle} \\vskip3pt \\centerline{\\it ftp://ftp.astrouw.edu.pl/ogle/ogle3/transits/new$\\_$2001$\\_$transits}  \\vskip3pt \\noindent or its US mirror  \\centerline{\\it http://bulge.princeton.edu/\\~{}ogle} \\vskip3pt \\centerline{\\it ftp://bulge.princeton.edu/ogle/ogle3/transits/new$\\_$2001$\\_$transits}  \\Acknow{We would like to thank Prof.\\ B.\\ Paczy{\\'n}ski for many interesting discussions and comments. The paper was partly supported by the  Polish KBN grant 2P03D01418 to M.\\ Kubiak. Partial support to the OGLE project was provided with the NSF grant AST-9820314 and NASA grant NAG5-12212 to B.~Paczy\\'nski.}"
    },
    "0207/astro-ph0207305_arXiv.txt": {
        "abstract": "We investigate the relationship between accretion rates and the spectral energy distributions (SEDs) of BL Lac objects, using a sample of objects for which published information on the host galaxies, emission-line luminosities, and peak frequencies and luminosities of their SEDs are available. The sample is composed of 43 BL Lac objects which have a relatively continuous distribution of peak frequencies. Under the assumption that the observed emission lines are photoionized by the central accretion disk, we use the line luminosities to estimate the accretion luminosities and hence accretion rates.  We find that low frequency-peaked BL Lac objects (LBLs) span a wide range of accretion rates, whereas high frequency-peaked BL Lac objects (HBLs) cover a more restricted range of lower values.  There appears to be a continuous distribution of accretion rates between the two subclasses of BL Lac objects. We find that the peak frequency of the SED, $\\pknu$, correlates with the accretion rate, approximately with the form $\\pknu\\propto \\Lambda^{-3}$ in HBLs and $\\pknu \\propto \\Lambda^{-0.25}$ in LBLs, where $\\Lambda \\equiv \\lline/c^2$. The peak luminosity of the SED is also correlated with $\\Lambda$.  These results suggest that the accretion rate influences the shape of the SED in BL Lac objects.  They also support models which couple the jet and the accretion disk.  We present a physical scenario to account for the empirical trends. ",
        "introduction": "The spectral energy distributions (SEDs) of BL Lac objects can be largely characterized by two broad peaks. The low and high energy peaks are commonly attributed to synchrotron emission and inverse Compton scattering, respectively.  Depending on the relative strengths of the two peaks, it is now customary to divide BL Lac objects into low frequency-peaked (LBL) and high frequency-peaked (HBL) sources (Giommi \\& Padovani 1994). Typically, HBLs are less variable, less polarized, and contain less dominant radio cores than LBLs (e.g., Laurent-Muehleisen et al. 1993; Perlman \\& Stocke 1993; Jannuzi, Smith \\& Elston 1994). However, the physical origin of the differences between the two classes of BL Lac objects remains a matter of debate.  According to the unification scheme for BL Lac objects, LBLs and HBLs are observed at different viewing angles (see review by Urry \\& Padovani 1995). On the other hand, Sambruna, Maraschi, \\& Urry (1996) have shown that the typical multiwaveband SED of an HBL cannot be produced from the spectrum of an LBL simply by changing the viewing angle alone. They suggested that other physical parameters, such as the magnetic field, may be different in the two classes.  New X-ray and radio surveys, such as those by Perlman et al. (1998), Laurent-Muehleisen et al. (1998), and Caccianiga et al. (1999), have shown that the distribution of peak frequencies in the SEDs of BL Lac objects is more uniform than previously thought.  What is the main physical parameter (or set of parameters) that control the SEDs of BL Lac objects?  Recent observations of the host galaxies of BL Lac objects provide new clues to understanding the physical nature of these systems.  Urry et al. (2000) and Scarpa et al. (2000) systematically studied the morphologies of BL Lac objects using the {\\it Hubble Space Telescope (HST)}\\ and found that there is no significant difference between the host galaxies of HBLs and LBLs. Most of the hosts are normal, giant elliptical galaxies, with no obvious evidence for ongoing mergers or strong interactions with other galaxies. Urry et al. (2000) suggest that, if the mass of the central black hole (BH) in BL Lac objects correlates with the bulge luminosity of the host galaxy, as is the case in nearby, inactive galaxies (e.g., Magorrian et al. 1998; Kormendy \\& Gebhardt 2001), the Eddington ratios\\footnote{The Eddington ratio refers to the ratio of the bolometric accretion luminosity to the Eddington luminosity, $L_{\\rm Edd} \\approx 1.3\\times 10^{38} \\mbh$ erg~s$^{-1}$, where $\\mbh$ is the BH mass in units of $M_{\\odot}$.} in  BL Lac objects must span a large range. A number of authors have suggested that the radio power in radio-loud quasars is controlled by the BH mass (e.g., Laor 2000; Lacy et al. 2001; McLure \\& Dunlop 2001; but see Ho 2002).  If the host galaxies of BL Lac objects obey the BH mass/bulge luminosity relation, then the peak frequencies are evidently not controlled by the BH mass (Urry et al. 2000).  This conclusion has been strengthened by the work of Falomo, Kotilainen \\& Treves (2002), Wu, Liu, \\& Zhang (2002), and Barth, Ho, \\& Sargent (2002a, b), who have obtained more robust BH mass estimates using the relation between BH mass and the stellar velocity dispersion of the bulge (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; Tremaine et al.  2002).  If BH mass is not the main parameter, what about the accretion rate? The importance of the role of accretion rates in BL Lac objects was originally recognized by Rees et al. (1982), who suggested that optically thin tori may power their central engines.  A number of authors have recently revisited this issue in the context of BL Lac objects (Ghisellini \\& Celotti 2001; \\centerline{\\includegraphics[angle=-90,width=8.0cm]{peak.ps}} \\figcaption{\\footnotesize Distribution of peak frequencies for our sample of BL Lac objects. Note that there is no bimodality in the distribution. \\label{fig1}} \\vskip 0.265cm \\noindent Maraschi 2001; B\\\"ottcher \\& Dermer 2002; Cavaliere \\& D'Elia 2002; Cao 2002) and more generally radio-loud quasars (Lacy et al. 2001). Using a large sample of objects which covers a wide range of nuclear activity, from nearly inactive systems to classical Seyfert 1 nuclei and quasars, Ho (2002) shows that the conventional ``radio-loudness'' parameter is strongly inversely correlated with the mass accretion rate.  This paper presents empirical evidence that the SEDs of BL Lac objects depend on the accretion rate. ",
        "conclusions": "Using a sample of BL Lac objects with information on their host galaxies and emission-line luminosities, we show that LBLs have significantly higher accretion rates and Eddington ratios than HBLs.  Both classes are highly sub-Eddington systems that may be accreting via an ADAF.  We find that the peak luminosity of the SED correlates significantly with the accretion rate, lending strong support to the idea that the jet and accretion disk are closely coupled.  Furthermore, we present evidence that the peak frequency of the SED also correlates with the accretion rate, although the functional dependence between the two parameters appears to be quite different between HBLs and LBLs. We argue that these empirical trends can be qualitatively explained by invoking different radiative efficiencies in the jets of the two classes of objects.  HBLs, which have lower accretion rates, evidently manage to convert more of their jet kinetic power into radiation.  By contrast, in LBLs, which have higher accretion rates, a greater fraction of the jet power remains in kinetic form.  However, the underlying causal connection between the radiative efficiency of the jet and the accretion rate, and the manner in which energy is channeled into relativistic electrons (Ghisellini, Celotti \\& Costamante 2002), remain unclear.  It would be of interest to perform spectrophotometric monitoring of nearby BL Lac objects in order to establish more conclusively the origin of their broad emission lines."
    },
    "0207/hep-th0207168_arXiv.txt": {
        "abstract": "Brane Gas Cosmology (BGC) is an approach to unifying string theory and cosmology in which matter is described by a gas of strings and branes in a dilaton gravity background. The Universe is assumed to start out with all spatial dimensions compact and small. It has previously been shown  that in this context, in the approximation of neglecting inhomogeneities and anisotropies, there is a dynamical mechanism which allows only three spatial dimensions to become large. However, previous studies do not lead to any conclusions concerning the isotropy or anisotropy of these three large spatial dimensions. Here, we generalize the equations of BGC to the anisotropic case, and find that isotropization is a natural consequence of the dynamics. ",
        "introduction": "The Standard Big Bang (SBB) cosmology has become an extremely successful model that has been well tested by experiment. However, the model is incomplete.  The underlying theory of classical general relativity and the description of matter as an ideal gas breaks down at the high temperatures of the early Universe, and the solutions of the theory in fact have an initial singularity. Moreover, SBB does not address many important cosmological questions such as the observed homogeneity, spatial flatness, and the origin of structure in the Universe.  Cosmological inflation (see e.g. \\cite{Linde:nc,Liddle:cg} for textbook treatises and \\cite{Watson:2000hb,Brandenberger:1999sw} for shorter reviews) builds on SBB cosmology providing a solution to some of these issues, but it (at least in the context of scalar field-driven inflation) suffers from the same initial singularity problem \\cite{Borde:xh} and other conceptual problems \\cite{Brandenberger:1999sw}, which indicate that inflation cannot be the complete story of early Universe cosmology.  In recent years, many models motivated by string theory and M-theory have emerged as possible solutions to the outstanding problems of early Universe and inflationary cosmology (see e.g. \\cite{Easson:2000mj,Brandenberger:2001ph} for recent but incomplete reviews). Beginning with the work on Pre-Big-Bang cosmology \\cite{Veneziano:1991ek,Gasperini:1992em} it was realized that a dynamical dilaton should play an important role in the very early Universe. More recently, models have become prominent in which  our Universe consists of a 3-brane embedded in a higher dimensional bulk space, with the standard model constrained to live on the brane \\cite{Arkani-Hamed:1998rs,Antoniadis:1998ig,Randall:1999vf,Randall:1999ee,Khoury:2001wf,Kallosh:2001ai,Steinhardt:2001vw}. Although these models can resolve a number of issues, such as the hierarchy problem, they introduce several other difficulties in the process. For example, large extra dimensions should be explained by classical general relativity, and it has been shown this results in problems stabilizing the brane \\cite{Carroll:2001ih}. More importantly, in most of these models the six ``extra'' spatial dimensions are taken to be compactified, {\\em a priori}, with no explanation provided for how this could come about dynamically.  Although it appears to be an important concern for the naturalness of these models, this issue is rarely discussed in the literature.  An alternative approach to string/M-theory cosmology is the string gas or BV scenario. This model began with works \\cite{Brandenberger:1989aj,Tseytlin:1992xk} in which the effects of string gases on the cosmological evolution of the low energy effective string theory background geometry including the dilaton were explored.  The most important result to emerge from these works is a dynamical mechanism, tied to the existence of string winding modes, which yields a nonsingular cosmology and may explain why at most three spatial dimensions can become large if the initial state is chosen to correspond to a Universe which is small in all spatial directions.  This work has been generalized to include the cosmological effects of p-brane gases and leads to the current model of Brane Gas Cosmology (BGC) \\cite{Alexander:2000xv}. In BGC, the Universe starts analogous to the SBB picture, i.e. hot, dense, and with all fundamental degrees of freedom in thermal equilibrium. The Universe is assumed to be toroidal in all nine spatial dimensions and filled with a p-brane gas. The assumption of toroidal geometry of the background space leads to the existence of string winding modes, since the background space  admits cycles on which  branes of the relevant dimensionalities (in particular one branes) can wrap. This wrapping is associated with a winding energy which - in the context of dilaton gravity - acts as a confining potential for the scale factor preventing further expansion of the spatial dimensions.  Also associated with the brane are oscillatory modes described by scalar fields and momentum modes which correspond to the center of mass motion of the brane. The momentum modes are related by T-duality to the winding modes and this duality results in the non-singular behavior of the model. In order for dimensions to decompactify, p-brane winding modes must annihilate with anti-winding modes and it is argued that this only occurs in a maximum of $2p+1$ dimensions \\cite{Alexander:2000xv}. Since strings $(p=1)$ are the lowest dimensional objects that admit winding modes, since they are the lightest of all winding modes and hence fall out of equilibrium later than other winding modes, it follows that the number of large space-time dimensions can be at most (3+1).  In the context of the background equations of dilaton gravity, the winding modes yield a confining potential for the scale factor which also gives rise to a period of cosmological loitering (expansion rate near zero) for the three large spatial dimensions. This is due to the time needed for winding modes to annihilate and produce closed strings or loops \\cite{Brandenberger:2001kj}. This is of great interest since loitering can explain the horizon and relic problems of standard cosmology without resorting to inflation \\footnote{However, to obtain a solution of the flatness and entropy problems, a phase of inflation following the decoupling of the three large spatial dimensions may be required}. It was also shown in \\cite{Brandenberger:2001kj} that by considering loop production at late times BGC naturally evolves into the SBB, with a $3+1$ dimensional, radiation dominated Universe.  There are important issues that remain to be addressed within BGC.  The fact that toroidal geometry was assumed for the background space was used for the existence and topological stability of winding modes. However, K3 or Calabi-Yau manifolds are more realistic choices for backgrounds within string theory and they do not admit 1-cycles (necessary to have topologically stable winding modes). Promising results have recently appeared which indicate that the conclusions of BGC extend to a much wider class of spatial background, including backgrounds which are K3 or Calabi-Yau manifolds \\cite{Easson:2001re,Easther:2002mi}.  Perhaps the main issue to be addressed in BGC is that of spatial inhomogeneities.  That is, we would naturally expect fluxes and p-brane sources to lead to the possibility of catastrophic instabilities of spatial fluctuation modes. Although we do not address this issue here, we plan to study the role of inhomogeneities in followup work.  Other important issues for BGC are the questions of stabilization of the six small extra dimensions and isotropization of the three dimensions that grow large.  Although these topics may seem to be unrelated, they both can be addressed by generalizing the BGC scenario to the anisotropic case.  This paper will concentrate on the isotropization of the three large dimensions, but the generalization of BGC to the anisotropic case achieved in this paper will be valuable to address the issue of stabilization in later work. ",
        "conclusions": "We have generalized the equations of BGC to the anisotropic case including the effects of winding state annihilation and loop production. We address the issue of isotropization of the three large dimensions quantitatively by introducing the anisotropy parameter ${\\cal A}$.  Our analysis indicates that for an arbitrary amount of initial anisotropy, the anisotropy will reach a maximum early in the evolution and then approach zero at later times. Thus, we have explained how isotropy arises as a natural consequence of BGC."
    },
    "0207/hep-ph0207295_arXiv.txt": {
        "abstract": "This article investigates the generation of non-Gaussianity during inflation. In the context of multi-field inflation, we detail a mechanism that can create significant primordial non-Gaussianities in the adiabatic mode while preserving the scale invariance of the power spectrum. This mechanism is based on the generation of non-Gaussian isocurvature fluctuations which are then transfered to the adiabatic modes through a bend in the classical inflaton trajectory. Natural realizations involve quartic self-interaction terms for which a full computation can be performed. The expected statistical properties of the resulting metric fluctuations are shown to be the superposition of a Gaussian and a non-Gaussian contribution of the same variance. The relative weight of these two contributions is related to the total bending in field space. We explicit the non-Gaussian probability distribution function which appears to be described by a single new parameter. Only two new parameters therefore suffice in describing the non-Gaussianity. ",
        "introduction": "The large scale structures of the universe are usually considered to arise from vacuum quantum fluctuations that are amplified during a stage of accelerated expansion. In its simplest version, inflation predicts the existence of an adiabatic initial fluctuation with Gaussian statistics and an almost scale-invariant spectrum~\\cite{inflation}. And it is clear that as long as the evolution is linear from the radiation era, non-Gaussianity can arise only from an ``initial'' non-Gaussianity generated during inflation.  Simple calculations, that we reproduce in the second section of this paper, however show that within single field inflationary framework it is not possible to produce primordial non-Gaussian fluctuations if the slow-roll conditions are preserved throughout the inflationary period during which the seeds of the large-scale structures are generated. Non-Gaussianity can be generated only if the inflation starts from a non-vacuum initial state~\\cite{martin} or if there exist sharp features in the shape of the potential~\\cite{staro}, but it in the latter case it clearly shows up in the density fluctuation power spectrum. It has been noticed however that the situation is somewhat changed when more than one light scalar field are present during inflation. In this case it is to be noticed first that one generically produces a mixture of adiabatic and isocurvature fluctuations~\\cite{linde,polarski,garcia,ms} that can be uncorrelated or correlated~\\cite{langlois,gordon,hwang}. The observational consequences of the existence of these two types of modes have started to be considered~\\cite{moodley1,lr,moodley2} but they clearly depend on which type of matter each of the field decays to. Multi-field inflation also opens the door to the generation of non-Gaussianity simply because the non-linear couplings can be much stronger in the isocurvature direction than in the adiabatic direction~\\cite{la,yamamoto,lm,salopek}. For instance in models in which a Peccei-Quinn symmetry is broken during inflation, it was pointed out by Allen {\\em et al.}~\\cite{allen} that a fourth order derivative term in the effective theory leads to non-Gaussianity in the axion density. In the \"seed\" models including $\\chi^2$~\\cite{chi2}, axion~\\cite{allen,axion}, Goldstone bosons~\\cite{bucher} and topological defects~\\cite{topdef} models, there is a test scalar field that is a pure perturbation and that does not contribute to the background energy density; the energy being quadratic in the field, it induces non-Gaussianity in the perturbations. In such models however the non-Gaussian features are present in the isocurvature modes only and will be observationally relevant only if those modes survive the reheating phase. The phenomenological situation is somewhat different however if a transfer of the modes is possible, that is when the fields are coupled~\\cite{bartolo2,bartolo1}.  The aim of this paper is therefore to explore whether in the context of multi field inflation it is possible to generate non-Gaussian features while preserving the adiabatic slow-roll type power spectrum and what would be its observational signature. We are obviously motivated by the development of Cosmic Microwave Background (CMB) and large scale structures observations that offer an opportunity reconstruct the properties of the primordial metric perturbations. Primordial fluctuations are more directly probed by CMB observations but then the number of modes that can be measured is still small. Up to now, no non-Gaussian signature has been detected in either the 4 year COBE data~\\cite{sandvik,rocha} or the 1 year MAXIMA data~\\cite{santos}. In large-scale structure surveys the number of independent modes one can observe is large but the difficulty is that the non-linear gravitational dynamics~\\cite{revue} generates non-Gaussian couplings that can shadow the primordial ones~\\cite{nous}.  One should then rely on a good understanding of the impact of the gravitational dynamics on the observations. Cosmic shear surveys might offer one of the most serious opportunity to explore such effects in the coming years~\\cite{fb,detection}.  We start (Section~\\ref{0}) by a general overview of the generation of non-Gaussianity in single and multi-field inflation. It will lead us to define a mechanism that can produce such non-Gaussianity. In Section~\\ref{I}, we then consider the evolution of perturbations in two-field inflation and summarize how isocurvature perturbation can be transferred to the adiabatic component when the trajectory in the field space is curved. The isocurvature mode develops non-Gaussianity due to self-interaction; we study in Section~\\ref{II} the evolution of such a self-interacting field in an expanding universe. After having posed the problem for a quantum scalar field in an inflationary background we address this issue from a classical point of view. In Section~\\ref{III} and ~\\ref{PDF properties}, we compute the probability distribution function that can be obtained in the class of models considered in this article. ",
        "conclusions": "In this paper we have explored the possibility of generating significant non-Gaussian initial metric perturbations in the context of inflationary cosmology while preserving a power spectrum of slow roll type adiabatic fluctuations. We found that the only viable mechanism is through a multiple field inflation where transverse (e.g. isocurvature) modes developed non Gaussian properties that can be subsequently transferred to the adiabatic fluctuations if the classical field trajectory is bent.   We have pointed out that quartic type coupling in the transverse modes is the most natural type of couplings for providing non-Gaussianities in the sense that in this case no fine-tuning in the value of the coupling constant has to be invoked. We stress that in the context of such a mechanism, unlike any others, the amount of non-Gaussianities that can be fuelled in the adiabatic fluctuations can be almost arbitrarily large.  We have examined in more details the case where the isocurvature mode generation (that is when they reach the Hubble size during the inflationary period) and the adiabatic-isocurvature mode mixing happen at very different time. The reason we consider this case is two-fold. First it is somewhat pedagogical since it shows that these two stages do not have to be concomitant. Second it implies that the non-Gaussian properties of the isocurvature modes developed mainly during the time they live at super-Hubble scales. It makes their computation much more simple since we can avoid a full treatment of the nonlinear field evolution at a quantum level (and it turns out that such a computation is not straightforward at all!). For modes living at super-Hubble scales we allowed ourselves to view the field evolution as the one of a classical stochastic field. In this case it is then possible to pursue the calculations to completion in the sense that it is possible to derive their whole set of correlation properties.  In particular we have been able to derive the one-point field cumulants at tree order in the weak coupling limit in a consistent way and finally build up the one-point probability distribution function of the field value. In such class of models, the statistical properties of the curvature perturbation are described by the superposition of a Gaussian and a non-Gaussian contributions with a relative weight proportional to the bending angle, $\\Delta\\theta$, of the trajectory in field space during slow-roll. The non-Gaussian component is fully characterized by a single parameter, $\\nu_3$, related to the reduced fourth order connected cumulant and has the same variance as the Gaussian contribution.  For practical purposes, we emphasize that its PDF is well approximated by Eq.~(\\ref{pdfapprox}) that reproduces accurately the numerically computed PDF within our approximation scheme. Thus, all the statistical properties can be encapsuled in two parameters.  These results give some insights on what type of non-Gaussian features can appear in future large-scale structure or CMB surveys while assuming the inflationary prejudice. Note that the kind of non-Gaussianity described here departs from that generated by the non-linear gravitational dynamics, in particular it has no skewness.  There is still however some ways between these results and their observational consequences. How non-Gaussian properties that are present at super-Hubble scales are transferred for instance to the CMB anisotropies at sub-Hubble scales is not totally straightforward. Whether such effects could be actually observed when observational aspects are taken into account demands in-depth analysis."
    },
    "0207/astro-ph0207010_arXiv.txt": {
        "abstract": "{ New near-IR long slit spectroscopic data obtained with ISAAC on VLT/ANTU (ESO/Paranal) complement and extend our previously published near-IR data (Alloin \\etal~\\cite{all01}) to produce \\Bg~and \\H2~emission line maps and line profile grids of the central 4''$\\times$4'' region surrounding the central engine of \\NGC1068. The seeing quality together with the use of an $0.3$'' wide slit and $0.3$'' slit position offsets allow one to perform 2D-spectroscopy at a spatial resolution $\\approx 0.5$''. Slit orientations (PA=102\\degr~and PA=12\\degr) were chosen so as to match respectively the equatorial plane and the axis of the suspected molecular/dusty torus in \\NGC1068. The selected wavelength range from 2.1 to 2.2\\micro~is suitable to detect and analyze the \\Bg~and \\H2~emission lines at a spectral resolution corresponding to 35\\kms. An asymmetric distribution of \\H2~emission around the continuum peak is observed. No \\H2~emission is detected at the location of the strong 2.2\\micro~continuum core (coincident within error-bars with the central engine location), while two conspicuous knots of \\H2~emission are detected at about 1'' on each side of the central engine along PA=90\\degr, with a projected velocity difference of 140\\kms: this velocity jump has been interpreted in Alloin \\etal~(\\cite{all01}) as the signature of a rotating disk of molecular material. From this new data set, we find that only very low intensity \\Bg~emission is detected at the location of the two main knots of \\H2~emission. Another knot with both \\H2~and \\Bg~emission is detected to the North of the central engine, close to the radio source C where the small scale radio jet is redirected and close to the brightest [OIII] cloud NLR-B. It has a counterpart to the South, placed almost symmetrically with respect to the central engine, although mainly visible in the \\Bg~emission. The northern and southern knots appear to be related to the ionization cone. At the achieved spectral resolution, the \\H2~emission line profiles appear highly asymmetric with their low velocity wing being systematically more extended than their high velocity wing. A simple way to account for the changes of the \\H2~line profiles (peak-shift with respect to the systemic velocity, width, asymmetry) over the entire 4''$\\times$4'' region, is to consider that a radial outflow is superimposed over the emission of the rotating molecular disk. We present a model of such a kinematical configuration and compare our predicted \\H2~emission profiles to the observed ones. Excitation of the \\H2~line is briefly discussed: X-ray irradiation from the central engine is found to be the most likely source of excitation. Given the fact that the material obscuring our direct view toward the central engine is Compton thick ($\\rm N_{H} \\geq 10^{24}\\,cm^{-2}$), the observed location of the main \\H2~knots at a distance of 70 pc from the central engine suggests that the rotating molecular disk is warped. ",
        "introduction": "Over the past decade, considerable efforts has been devoted to a better understanding of active galactic nuclei (AGN). In particular, the use of HST (UV and visible ranges), the use of adaptive optics (AO) on 4m-class telescopes (near-IR range) and the use of interferometry (millimeter and radio ranges) have brought a large gain in spatial resolution and hence many new constraints for AGN modeling. Conversely, the so-called unified model (see Krolik \\cite{kro99}) has played an important role in stimulating the search for AGN constituents: in particular for the elusive pc-scale molecular/dusty torus which is thought to surround the central engine (black hole and accretion disk) and to funnel the ionizing radiation within a cone. Models of the infrared (IR) emission of the torus published so far, have explored torus radii from 1 to 100\\pc~(Krolik \\& Begelman \\cite{kro86}, Pier \\& Krolik \\cite{pie93}, Granato \\& Danese \\cite{gra94}, Efstathiou \\& Rowan-Robinson \\cite{efs94}, Granato, Danese \\& Franceschini \\cite{gra97}), while the thickness, composition and inclination of the torus are other key-parameters in the modeling. Given the input energy distribution from the central engine, models then predict the flux spatial distribution at IR wavelengths from 2.2\\micro~up to 20\\micro~in the central 100\\,\\pc~region of the AGN. If one chooses an AGN for which the spatial resolution reachable today (typically 0.1'' with the facilities mentioned above in a range from UV to millimetric) provides an intrinsic scale relevant to the torus size, a comparison between predicted IR maps and observed ones will be an excellent opportunity for deriving stringent constraints on the model parameters. This calls for investigating a close-by and torus-inclined AGN: \\NGC1068 is an ideal target.  This galaxy is located at a distance of 14.4\\,Mpc, assuming a Hubble constant of 75\\Hubble. The corresponding scale is 70\\pc~per arcsec. Many high resolution images have indeed been obtained over the past decade: decisive contributions come from the HST in the UV, visible (Capetti \\etal~\\cite{cap97}) \\& near-IR (Thompson \\etal~\\cite{tho01}), from AO in the near-IR (Marco \\etal~\\cite{mar97}, Thatte \\etal~\\cite{tha97}, Rouan \\etal~\\cite{rou98} and Marco \\& Alloin \\cite{mar00}), using shift-and-add techniques or diffraction-limited observations on 8m class telescopes in the mid-IR (Bock \\etal~\\cite{boc00}, Tomono \\etal~\\cite{tom01}), from millimeter interferometry (Helfer \\& Blitz \\cite{hel95}, Tacconi \\etal~\\cite{tac97}, Backer \\cite{bac00}, Schinnerer \\etal~\\cite{sch00}) and radio interferometry (Gallimore \\etal~\\cite{gal97}). In this paper, we consider the compact radio source S1 (nomenclature after Gallimore \\etal~\\cite{gal96}) as the nuclear reference, that is the central engine location. Several observational facts suggest the presence of a molecular/dusty torus in \\NGC1068. The term torus is used here in the sense of a rotating disk-like distribution. From inner to outer scales, let us point out the most relevant ones: \\begin{itemize} \\item The 1\\pc-size disk of ionized material seen almost edge-on at PA=110\\degr, within S1, which has been detected at 8.4\\,GHz with the VLBA: it is thought to trace the inner walls of the torus (Gallimore \\etal~\\cite{gal97}) \\item The core and the 15\\pc-size elongated structure (at PA=102\\degr), detected at 2.2\\micro~with CFHT AO/PUEO (Rouan \\etal~\\cite{rou98}) and also present on 3.5 and 4.5\\micro~images with ESO AO/Adonis (Marco \\& Alloin \\cite{mar00}); the position of the near-IR core is found by Marco \\etal~(\\cite{mar97}) to be coincident with S1 within $\\pm 0.05$''. This configuration is consistent with thermal emission from dust particles at 1500K in the core, and 600K at a radius of 15\\pc~and could trace the body of the torus. \\item From interferometer maps (Backer~\\cite{bac00}, Schinnerer \\etal~\\cite{sch00}) one observes, in addition to a yet unresolved core of low level CO emission, two prominent CO emitting peaks located symmetrically 1'' (70\\pc) away from S1 at PA=98\\degr~and separated in velocity by 100\\kms: the authors interpret the CO observed emission in terms of a warped CO disk. \\end{itemize}  The possible presence of a molecular/dusty torus could be unveiled through the analysis of the \\H2~roto-vibrational $\\nu$=1-0\\,S(1)~emission line at rest wavelength 2.122\\micro. Indeed, excitation of molecular transitions such as \\H2~2-1\\,S(1) and \\H2~1-0\\,S(1) can occur from irradiation by UV photons or X-rays, as well as from shocks, all ingredients which are found in profusion in an AGN. In addition, as the \\H2~emission arises from warm molecular gas, the study of the \\H2~lines provides physical information complementary to that derived from CO molecular transitions.  Molecular hydrogen was detected for the first time in an extragalactic source in the central region of \\NGC1068 by Thompson, Lebofsky \\& Rieke (\\cite{tho78}). They also detected \\Bg~emission in this region. Further \\H2~line observations were carried out by Hall \\etal~(\\cite{hal81}), and Oliva \\& Moorwood (\\cite{oli90}). The first attempt at resolving spatially the \\H2~emitting region was made by Rotaciuc \\etal~(\\cite{rot91}): from their best-seeing map they found that the \\H2~line emission was very weak at the location of the central engine and was mainly arising from two knots of unequal intensity located at $\\sim$1.3'' from the nucleus along PA=70\\degr. A second attempt to get a $\\leq 1$'' resolution image in the \\H2 line emission was made by Blietz \\etal~(\\cite{bli94}). In the latter study, in contrary to Rotaciuc \\etal~(\\cite{rot91}), the \\H2 emission was found to arise in three knots, the brightest one being coincident (within an error bar of $\\pm$0.3'') with the strong near-IR continuum core (central engine location).  In this study, we analyze the \\H2~and \\Bg~emission lines, at respective rest wavelengths 2.122 and 2.166\\micro~using ISAAC on VLT/ANTU, in its short wavelength medium resolution mode. Preliminary results about the \\H2~emission within a small region (1.5''$\\times$1.5'') around the central engine were already published in a Letter (Alloin \\etal~\\cite{all01}). In the current paper we extend the size of the region explored to 4''$\\times$4'' around the nucleus and provide  \\H2~line profiles and \\H2~line emission maps over this region. We also analyze, similarly to \\H2, the \\Bg~line emission map and line profiles. From the \\H2~line profile analysis, we build a kinematical model of the warm molecular component in the AGN of \\NGC1068.  We present in Section 2 the relevant information for data collection and data processing. In Section 3 we provide complete sets of the \\H2~and \\Bg~line profiles across the entire region explored. From these data sets, emission maps in the \\H2~and \\Bg~lines have been reconstructed and are compared to the map in the near-IR continuum. Fluxes and velocities are also given in Section 3. From the \\H2~line peak velocities and profiles, we discuss in Section 4 the kinematics of the molecular material and we present in Section 5 a simple kinematical model, aiming at reproducing the profile shapes. Possible sources of excitation of the \\H2~line in the close environment of the AGN in \\NGC1068 are discussed in Section 5. Our concluding remarks are presented in Section 6. ",
        "conclusions": "We have provided in this paper a comprehensive study of the kinematics of the warm molecular material in the environment of the AGN in \\NGC1068. New runs of 2D spectroscopic observations (ISAAC on VLT/ANTU), obtained under very good seeing conditions, have complemented the data set discussed earlier by Alloin \\etal~(\\cite{all01}). The achieved spectral (35\\kms) and spatial (0.5'') resolutions allow to recover emission maps in the \\H2~and \\Bg~emission lines, at respective rest wavelengths 2.122 and 2.166\\micro~, over the central 4''$\\times$4'' region surrounding the central engine in \\NGC1068. In addition, one witnesses the spatial evolution of the \\H2~and \\Bg~line profiles over the inner 3''$\\times$3'' region with a spatial sampling of 0.3''$\\times$0.3''.\\\\ The \\H2~and \\Bg~line emission maps and profile spatial evolution have been used to analyze the distribution and kinematics of the warm molecular gas across the 200\\pc~central region. The salient results of this work are the following: \\begin{itemize} \\item There is no detectable \\H2~emission at the location of the strong 2.2\\micro~continuum core (which is known already to be almost coincident with the central engine), while two main regions of \\H2~emission, East-\\H2~and West-\\H2, are detected at about 70\\pc~on each side of the central engine along PA=90\\degr, with a projected velocity difference of 140\\kms: this velocity jump is interpreted at first order as the signature of a rotating structure. Under this assumption, the systemic velocity of the AGN is found to be 1144$\\pm$4\\kms. No strong \\Bg~emission is detected at the location of these knots. The East-\\H2~knot is about three times brighter in \\H2~line emission than the West-\\H2~knot \\item One detects as well a diffuse and extended \\H2~emission, although at low intensity level. At about 0.3'' to the North of the central engine, the \\H2~emission also coincides with a \\Bg~line emitting knot. This \\north -\\H2-\\Bg~knot is close to the radio knot C and the brightest [OIII] cloud, NLR-B, although not exactly spatially coincident with them. \\item To the South, another \\Bg~line emitting region is observed, although weaker in intensity, the South-\\Bg~knot. The \\north -\\H2-\\Bg~knot and South-\\Bg~knot are almost symmetrical with respect to the central engine and appear to be located on the inner edges of the ionizing cone.  \\item The observed changes in the \\H2~line profile across the 200\\pc~central region cannot be reproduced by the simple, disk-like, kinematical structure which matches the velocity jump between of the East-\\H2~knot and the West-\\H2~knot. A possible solution is to add an outflow at the surface of the rotating disk-like structure. The main parameters of the complete kinematical model are: a rotational velocity around 100\\kms~at a characteristic radius of 50\\pc~, an inclination angle of the disk-like structure of 65\\degr~and an outflow velocity of 140\\kms. \\item The source of excitation of the \\H2~line is briefly discussed. Both UV photons and shocks can be discarded as the main source of excitation of the \\H2~line, while X-ray irradiation from the central engine is found to be the most likely mechanism. \\item An additional consequence of an outflow at the surface of a rotating disk-like structure is the occurrence of a warp. This warp does naturally explain how the X-rays can impact the molecular disk at a distance as large as 70\\pc~and give rise to the two observed main knots emitting \\H2. \\end{itemize}"
    },
    "0207/astro-ph0207226_arXiv.txt": {
        "abstract": "We discuss the possibilities of measuring ultra-high energy cosmic rays and neutrinos with radio techniques.  We review a few of the properties of radio emission from cosmic ray air showers and show how these properties can be explained by coherent ``geosynchrotron'' emission from electron-positron pairs in the shower as they move through the geomagnetic field. This should allow one to use the radio emission as a useful diagnostic tool for cosmic ray research. A new generation of digital telescopes will make it possible to study this radio emission in greater detail.  For example, the planned Low-Frequency Array (LOFAR), operating at 10-200 MHz, will be an instrument uniquely suited to study extensive air showers and even detect neutrino-induced showers on the moon. We discuss sensitivities, count rates and possible detection algorithms for LOFAR and a currently funded prototype station LOPES. This should also be applicable to other future digital radio telescopes such as the Square-Kilometer-Array (SKA).  LOFAR will be capable of detecting air-shower radio emission from $>2\\cdot10^{14}$ eV to $\\sim10^{20}$ eV. The technique could be easily extended to include air shower arrays consisting of particle detectors (KASCADE, Auger), thus providing crucial additional information for obtaining energy and chemical composition of cosmic rays. It also has the potential to extend the cosmic ray search well beyond an energy of $10^{21}$ eV if isotropic radio signatures can be found. Other issues that LOFAR can address are to determine the neutral component of the cosmic ray spectrum, possibly look for neutron bursts, and do actual cosmic ray astronomy. ",
        "introduction": "A standard method to observe energetic cosmic rays is simply an array of particle detectors on the ground measuring either the energetic muons or electrons in the shower pancake. Since only a small fraction of the total number of particles in the shower are intercepted by the ground array, the conversion from number of particles received to primary particle energy is not really straight forward. Very useful additional information for energy calibration and particle track recovery of air showers can therefore be gained by observing radiation emitted from the secondary particles as the shower evolves. Such radiation is for example Cherenkov radiation, as observed in the CASA-MIA-DICE experiment, or fluorescence light from nitrogen molecules in the atmosphere, as seen by the Fly's Eye detector HiRes and others.  So far this emission is only detected in the optical and hence requires clear and moonless dark skies far outside major cities. This gives a duty-cycle of typically 10\\%.  Radio emission might provide an alternative method for doing such observations including detecting neutrino-induced showers at a higher duty cycle. This becomes particularly relevant since a new generation of digital radio telescopes -- designed primarily for astronomical purposes -- promises a whole new approach to measuring air showers. ",
        "conclusions": "While the investigation of the radio emission from extensive air showers has lain dormant for a rather long time there is enough information available that suggests that this field could be revived. The properties of these radio pulses from cosmic ray air showers are all consistent with it being coherent geosynchrotron emission from electrons and positrons in the air shower. This process is basically unavoidable and hence the radio emission should directly reflect the shower evolution of the leptonic component of cosmic ray air showers if properly measured.  Because the emission is highly beamed, a key to successful usage of the radio emission is the rather new possibility to build digital telescopes that combine a large field of view with the ability to form virtual beams retrospectively in the direction of transient events. For this reason, the planned radio array LOFAR (and possibly also the SKA) will become a very efficient cosmic ray detector which is sensitive to high-energy cosmic rays at all energies from $\\sim10^{14}$ to $10^{20}$ eV. The calibration and accuracy could be further improved by combining this digital radio technology with existing or upcoming air shower arrays consisting of particle counters on the ground (KASCADE/LOPES,Auger).  The combination of radio techniques and particle counters should provide a unique tool to study the energy spectrum and composition of cosmic rays over a broad range rather efficiently, simultaneously probing a parameter space never combined in a single array. Moreover, at energies around $10^{18}$ eV neutron astronomy would, for the first time, become possible. A large radio array like LOFAR could also be used to search for radio emission from neutrino induced showers in the air or from the lunar regolith, possibly opening a new window to the universe. Hence,  digital radio telescopes could provide a significant technological advantage for astronomy and astroparticle physics.  \\medskip {\\bf Acknowledgment:} Many people have contributed with their suggestions and comments to this paper. We particularly want to thank Tim Huege, Alan Roy, Peter Biermann, Karl Mannheim, Ger de Bruyn,  and the participants of the LOFAR workshop in Dwingeloo, May 2001, for intensive discussions on this and related topics. We thank Patricia Reich for providing us with a properly smoothed version of the 408 MHz sky survey.  We are also grateful to the editor, Alan Watson, an anonymous referee, and Harold Allan for helpful comments on the manuscript. Ralph Spencer provided some useful insight into the details of the early radio experiments. We thank Walter Fu\\ss{}h\\\"oller for helping to prepare some of the figures. This work has been performed in part at the Jet Propulsion Laboratory, California Institute of Technology, under contract with the U.S. National Aeronautics and Space Administration.  Funding was also provided in part by the German Ministry of Education and Research (BMBF), grant 05 CS1ERA/1 (LOPES)."
    },
    "0207/astro-ph0207156_arXiv.txt": {
        "abstract": "We present a new model for the Galactic distribution of free electrons. It (a) describes the distribution of the free electrons responsible for pulsar dispersion measures and thus can be used for estimating the distances to pulsars; (b) describes large-scale variations in the strength of fluctuations in electron density that underly interstellar scattering; (c) can be used to interpret interstellar scattering and scintillation observations of Galactic objects and of extragalactic objects, such as intrinsically compact AGNs  and Gamma-ray burst afterglows; and (d) serves as a preliminary, smooth spatial model of the warm ionized component of the interstellar gas. This work builds upon and supersedes the Taylor \\& Cordes (1993) model by exploiting new observations and methods, including (1)~a near doubling in the number of lines of sight with dispersion measure or scattering measurements; (2)~a substantial increase in the number and quality of independent distance measurements or constraints; (3)~improved constraints on the strength and distribution of scattering in the Galactic center; (4)~improved constraints on the (Galactocentric) radial distribution of free electrons; (5) redefinition of the Galaxy's spiral arms, including the influence of a local arm; (6) modeling of the local interstellar medium, including the local hot bubble identified in X-ray and \\ion{Na}{1} absorption measurements; and (7)~an improved likelihood analysis for constraining the model parameters. For lines of sight directed out of the Galactic plane, the new model yields substantially larger values for pulsar dispersion measures, expect for directions dominated by the local hot bubble. Unlike the TC93 model, the new model provides sufficient electrons to account for the dispersion measures  of the vast majority of known,  Galactic pulsars. The new model is described and exemplified using plots of astronomically useful quantities on Galactic-coordinate grids.  Software available on the Internet is also described. Future observations and analysis techniques that will improve the Galactic model are outlined. ",
        "introduction": "\\label{sec:intro}  Radio wave propagation measurements provide unique information about the magnetoionic component of the interstellar medium (ISM).  Pulsars are important probes because they emit short duration radio pulses that are modified by intervening plasma and also because they are highly spatially coherent, allowing scattering processes to significantly perturb their radiation.    Other compact sources, both Galactic and extragalactic, also serve as probes of the plasma.  In this first of a series of papers, we present a new model for the free electron density of the Galaxy.  It is called ``NE2001'' because it incorporates data obtained or published through the end of~2001.  A short history of such models prior to~1993 is given in Taylor \\& Cordes~(1993, hereafter TC93).  Our work builds upon and supersedes the TC93 model by exploiting new dispersion and scattering measurements and also by employing new techniques for modeling. Following Cordes \\etal~(1991) and TC93, we model fluctuations in electron density as well as the local mean density.  Since TC93 was written, significant developments have taken place that increase the sample of relevant measurements and indicate shortcomings of the model. Most importantly, independent distance measurements are available on about 50\\% \tmore objects.   Some of these are precise parallax measurements obtained through pulse timing or interferometric techniques.  Others result from \\ion{H}{1} absorption measurements combined with a kinematic rotation model for the Galaxy.  Still others arise from the association of pulsars with supernova remnants and globular clusters.  Distance estimates combined with dispersion measures quantify the line-of-sight average electron density. The new distance constraints indicate, in a minority of cases, that some of the distances derived from the TC93 model using pulsar dispersion measures are in error by factor of two or more.  However, we point out that some of the parallax measurements also differ from previous ones by significant amounts. Additional information is provided by the distribution of dispersion measures for the entirety of available pulsar samples. The number of such measurements is about double of that available in 1993. Constraints on the square of the electron density arise from scattering measurements such as angular broadening, pulse broadening and diffractive scintillation measurements. The number of scattering measurements has almost doubled since 1993.  The TC93 model is flawed in several respects, some of which were apparent even at the time of its development.  First, it provides insufficient column density at high Galactic latitudes so that only lower bounds on pulsar distances can be derived from it.  Second, in some directions, particularly in the fourth quadrant along tangents to the Carina-Sagittarius and Crux-Scutum spiral arms, the model provides either too many electrons for some objects (Johnston \\etal~2001) or too few for others (as discussed below), indicating that the spiral arms need redefinition.  Third, in the direction of the Gum Nebula and Vela supernova remnant, the TC93 model provides too little scattering to account for the pulse broadening of some pulsars (Mitra \\& Ramachandran 2001).  Fourth, interstellar scintillations of nearby pulsars have scintillation bandwidths (which measure the column density of electron density fluctuations) that are not well modeled (Bhat \\& Gupta~2002).  Fifth, similar to the first deficiency, the calibration between scattering of Galactic and extragalactic sources needs revision in order to ascertain the column density of scattering material toward cosmological sources as compared to that of Galactic objects outside of but near the apparent boundary of free electrons.  Our knowledge of the ISM has increased significantly from a host of other investigations across the electromagnetic spectrum. The local ISM, in particular, is now much better modeled and we incorporate that information into our model for the free electron density. The local ISM has been probed using continuum  X-ray measurements and absorption of \\ion{Na}{1} toward nearby stars.  Other models for the mean electron density have been presented since TC93.  These include the recent work of G\\'omez, Benjamin \\& Cox (2001; hereafter GBC01), who presented a two-component, axisymmetric model, not dissimilar in form to that used by Cordes \\etal\\ (1991). The two components have sech$^2(r/R)$sech$^2(z/H)$ variations with different radial and $z$ scales. GBC01 do not attempt to model electron density fluctuations relevant to scattering and their model is based solely on the 109 objects available to them that have independent distance estimates.  As we demonstrate, though their model is adequate for nearby pulsars, it grossly underpredicts the dispersion measures of many distant pulsars at low Galactic latitudes  and some at high latitudes.  Their work underscores the need, demonstrated before by others (Ghosh \\& Rao 1992; TC93), for spiral-arm structure in the free-electron distribution.  Our model uses significantly larger data samples than were available for TC93 and makes use of all available data, including independent distance constraints on pulsar distances, dispersion and scattering data on pulsars, and scattering of other Galactic as well as extragalactic sources.   We also  incorporate published, multiwavelength data that allows modeling of the local ISM and of the spiral arms of the Galaxy.   In this first paper we present the model and its usage. In a second paper (Cordes \\& Lazio 2002b; hereinafter Paper~II) we describe the input data and associated references, the likelihood analysis and modeling, alternative model possibilities, and discussions of particular lines of sight.  Future papers will apply the model to various astronomical and astrophysical applications.  The plan of this paper is as follows. In \\S\\ref{sec:observables} we describe the observable quantities that we use to constrain the NE2001 model parameters. In \\S\\ref{sec:model} we describe the various components of the NE2001 model. In \\S\\ref{sec:performance} we demonstrate the model's ability to account for the distances and scattering of pulsars and discuss briefly astronomical applications of the model.  Extensive discussion of applications of NE2001 is deferred to a later paper.  In \\S\\ref{sec:discussion} we summarize the results and outline future prospects for improving the model.  These will rest largely on usage of the Parkes Multibeam Pulsar sample (e.g. Manchester \\etal\\ 2001), improved parallax measurements using very long baseline interferometry, and incorporation of additional multiwavelength observations into the model definition and fitting. Appendix~\\ref{app:LISM} describes our model for the local interstellar medium. Appendix~\\ref{app:fortran} describes how to obtain the model as a set of Fortran subroutines and its implementation in tools available through the World Wide Web. ",
        "conclusions": "\\label{sec:discussion}  We have presented a new model, NE2001, for the Galactic distribution of free electrons and the fluctuations within it.  As observational constraints we make use of pulsar dispersion measures and distances and radio-wave scattering measurements available at the end of 2001 (hence its name), and we are guided by multi-wavelength observations of the Galaxy, particularly of the local interstellar medium.  Building on the Taylor-Cordes model (TC93), we model the free electron distribution as composed of three large-scale components, a thick disk, thin disk, and spiral arms.  Since the publication of the TC93 model, though, the number of available data have expanded greatly (e.g., the number of pulsar DMs available is now approximately double the number available to TC93).  With this larger data set, it is clear that ``smooth,'' large-scale components are insufficient to produce a realistic description of the electron density distribution.  We must take into account the distribution of electrons in the local ISM, and we require ``clumps'' and ``voids'' of electrons, mesoscale structures distributed throughout the Galaxy on a large number of lines of sight in order to produce reasonable agreement with the observations.  Tables~\\ref{tab:ne2001}--\\ref{tab:lism} summarize the model and the best-fitting parameters of the large-scale and local ISM components. Tables~\\ref{tab:clumpsX}--\\ref{tab:voids} lists the lines of sight requiring clumps or voids and the relevant parameter for each clump or void.  Figure~\\ref{fig:grayscale} shows the model electron density in the Galactic plane.  We used an iterative likelihood method to find the best-fitting model parameters.  Our focus here has been on the exposition of the model. In a companion paper (Cordes \\& Lazio 2002b) we describe details of the fitting procedure and  discuss those lines of sight that require a clump or void or are otherwise problematic. We also discuss possible alternative  fitting approaches and models.  Our model, NE2001, improves upon the TC93 model.  First, the distance estimates obtained from the model agree with available distance constraints for nearly all pulsars with such constraints.  Second, none of the parameters of the large-scale components are indeterminate (e.g., as was the case with the thick disk for TC93).  The cost of these improvements has been an increase in the complexity of the model, particularly with respect to the number and location of clumps and voids.  We believe, however, that this additional complexity is justified both by the quantity of data and because it is astrophysically reasonable.  A small number of pulsars or extragalactic sources has been known for some time to have anomalously large DMs or scattering properties or both due to intervening \\ion{H}{2} regions or supernova remnants.  While we consider NE2001 to be an improvement over TC93, we foresee a number of probable developments that will allow future generations of electron density models. We group these improvements into increases in the quantity of data and improvements in the modeling technique.  Perhaps most important will be an increase in the number of pulsar parallaxes, both from pulse timing methods applied to millisecond pulsars and from large programs  using very long baseline interferometry (VLBI). Also, using pulsar luminosities to estimate distances should become feasible once beaming of pulsar radiation becomes better understood and despite the fact that the luminosity function for radio emission from pulsars is very broad. DM-independent distances provide crucial calibration information for NE2001 or any successors, and we regard it as likely that the number of pulsars with DM-independent distances will double in the next few years.  We have made use of only the positions and DMs of pulsars discovered in the Parkes multibeam sample.  Efforts are underway to measure the scattering along the lines of sight to many of these pulsars, which could increase the number of lines of sight with measured SMs by roughly 50\\% or more.  The advent of the Green Bank Telescope and the refurbished Arecibo telescope suggest the possibility of conducting a northern hemisphere equivalent of the Parkes multibeam sample, which could increase the number of pulsars by at least another 50\\%.  Possible improvements in modeling include the adoption of a  more realistic location for the Sun.  NE2001 places the Sun in the Galactic plane at a Galactocentric distance of~8.5~kpc, though a smaller distance from the Galactic center ($\\approx 7.1$~kpc) and a modest offset from the plane ($\\approx 20$~pc) seem warranted.  Although we have provided a formalism for comparing DM and SM to EM, we have made little use of it.  Future work to include observational constraints on EM, e.g., from H$\\alpha$ surveys, has the potential of producing a better model for the local ISM.  Finally, Figure~\\ref{fig:dmvsl} suggests that the large-scale structure of the Galaxy could be determined \\textit{ab inito}, provided that a sufficient number of lines of sight exist. Rather than imposing a large-scale structure as done both here and previously, the presence and location of large-scale components, particularly the spiral arms, could be determined directly from the population of pulsars.  Future pulsar surveys may yield the number of lines of sight required to employ this approach."
    },
    "0207/astro-ph0207360_arXiv.txt": {
        "abstract": "We present the first spectroscopic survey of intrinsically low X-ray luminosity clusters at $z\\gg0$, with {\\it Hubble Space Telescope (HST)} WFPC2 imaging and spectroscopy from Calar Alto and WHT-LDSS2.  We study 172 confirmed cluster members in a sample of ten clusters at $0.23<z<0.3$, with $L_X\\lesssim 4\\times 10^{43} h^{-2}$ ergs s$^{-1}$ [0.1-2.4 keV] ($\\Omega_m=0.3,\\Lambda=0.7$). The core of each cluster is imaged with WFPC2 in the F702W filter, and the spectroscopic sample is statistically complete to $M_r\\sim-19.0+5\\log{h}$, within an 11\\arcmin\\ ($\\sim 1.8$ h$^{-1}$ Mpc) field.  The clusters are dynamically well-separated from the surrounding field and most have velocity distributions consistent with Gaussians.  The velocity dispersions range from $\\sim 350$--$850$ km s$^{-1}$, consistent with the local $L_X-\\sigma$ correlation. All ten clusters host a bright, giant elliptical galaxy without emission lines, near the centre of the X-ray emission. We measure the equivalent width of two nebular emission lines, \\oii\\ and \\ha, and the \\hd\\ absorption line to spectrally classify the cluster members.  Galaxy morphologies are measured from the {\\it HST} images, using the two-dimensional surface-brightness fitting software {\\sc gim2d}. Emission line galaxies in these clusters are relatively rare, comprising only 22$\\pm 4$\\% of the sample. There is no evidence that these emission-line galaxies are dynamically distinct from the majority of the cluster population, though our sample is too small to rule out the $\\sim 30$\\% difference that has been observed in more massive clusters. We find eleven galaxies, comprising 6\\% of the cluster members, which are disk-dominated but show no sign of emission in their spectrum.  Most of these are relatively isolated, spiral galaxies with smooth disks. We find no cluster members with a starburst or post-starburst spectrum.  The striking similarity between the spectral and morphological properties of galaxies in these clusters and those of galaxies in more massive systems at similar redshifts implies that the physical processes responsible for truncating star formation in galaxies are not restricted to the rare, rich cluster environment, but are viable in much more common environments. In particular, we conclude that ram pressure stripping or cluster-induced starbursts cannot be solely responsible for the low star formation rates in these systems. ",
        "introduction": "The evolutionary history of galaxies depends both on cosmic time and on the type of environment in which they exist. For example, recent studies have shown that the universal average star formation rate (SFR) has been declining steadily since at least $z\\sim 1$ \\citep{L96,Madau,Cowie+99,Wilson+02}. In a somewhat analogous way, star formation rates are known to monotonically decrease with increasing density at a given epoch \\citep[e.g.][]{B+97,P+99,2dF-sfr}.  In both cases, the reason for the decrease in star formation is unknown.  In particular, there is plenty of gas still available for star formation at the present day, so the sharp decline in activity over the past $\\sim 5$ Gyr is a critical issue. It is an intriguing possibility that the processes which influence the evolution of galaxies that end up in dense clusters may be more generally important to galaxies in other, more common, environments.  If this is so, it may be possible to link the decline in the universal average star formation rate  to environmental effects in a Universe which is growing hierarchically.  We now have a good empirical description for the galaxy populations of massive clusters.  Galaxies within the virial radius have, on average, lower star formation rates, and less recent ($<1$Gyr) star formation, than galaxies in the surrounding field;  this is true both locally \\citep{2dF-sfr,Sloan_sfr} and at higher redshifts \\citep[e.g.][]{CS87,B+97,PSG,P+99,PLO,C+01,A1689}.  Evidence is mounting that this deficiency in star formation activity is at least partially independent of the morphology-density relation \\citep{D+97,B+98,P+99,C+01,2dF-sfr,Sloan_sfr}. However, an explanation for this difference between cluster and field galaxies is still lacking.  Ram pressure stripping of cold gas in the disk of a galaxy  \\citep{GG,F98,QMB} is only likely to take place in the dense cores of rich clusters, and it seems unlikely that it can explain the suppression of star formation as far as several Mpc from the centre \\citep{B+97,Kodama_cl0939}.  Galaxy harassment \\citep{harass} may be effective at destroying small galaxy disks, but the effect that this will have on the star formation rate of the galaxy is not clear.  On the other hand, the observed radial and density dependences of galaxy stellar populations and morphologies are reproduced quite well by hierarchical models in which only the diffuse, hot halo gas expected to surround isolated galaxies is stripped \\citep{LTC,infall,Diaferio,Okamoto,Bekki02}.  The observed increase in activity with redshift is most likely a consequence of the higher infall rate \\citep{PSext,Erica,KB01}, though projection effects may still play a role \\citep{Diaferio}.  For most cluster galaxies, the last episode of star formation occurred many billions of years ago.  Thus, if some physical mechanism is responsible for the transition from a more active state, it must have occurred in the distant past, and will be difficult to uncover by observing galaxies in their present state \\citep[e.g.][]{TragerVI,deJ+D,K+01}.  It is therefore necessary to consider how the galaxy populations have evolved over time \\citep{BO84,D+99,KS00,KB01}.  However, in hierarchical models of galaxy formation, the progenitors of today's most massive clusters are expected to be numerous smaller structures at higher redshifts \\citep[e.g., ][]{K96}.  Thus, galaxies must be observed not only over a range of redshifts, but for a range of cluster masses as well.  For this purpose we have carried out an extensive observational campaign to obtain {\\it Hubble Space Telescope (HST)} imaging and ground-based imaging and spectroscopy for ten clusters at $z\\approx0.25$, selected to have low X-ray luminosities (hereafter referred to as the \\lowlx\\ sample). This sample can be directly compared with studies of more massive clusters, both locally \\citep{2dF-sfr,Sloan_sfr} and at higher redshift \\citep{B+97,GPS2}. For example, in \\citet[][hereafter Paper I]{lowlx-morph} we presented an analysis of the {\\it HST} data for the present cluster sample, and compared it with a similar {\\it HST} sample of high X-ray luminosity cluster cores.  We found marginal ($\\sim 2\\sigma$) evidence that the \\lowlx\\ clusters have more disk-dominated galaxies at a fixed local density.  This suggests that at least galaxy morphology is sensitive to the large-scale structure.  In the present paper we will revisit this, and other issues, in light of the spectroscopic data.  The paper is organized as follows.  In \\S~\\ref{sec-obs} we present the cluster selection, and the data acquisition, reduction and analysis. Our results are presented in \\S~\\ref{sec-results}, where we consider the dynamics and spectral properties of the galaxy population. In \\S~\\ref{sec-discuss} we compare our results with those found for more massive clusters, and consider the implications of these results for models of galaxy evolution.  We summarize our findings in \\S~\\ref{sec-conc}. We use a cosmology with $\\Lambda=0.7$, $\\Omega_m=0.3$, and parametrise the Hubble constant as $H_\\circ=100h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec-conc} We have analysed ground-based spectroscopy and {\\it HST}-based morphologies of galaxies in ten clusters at $z\\approx0.25$ with low X-ray luminosities.  The sample includes 165 galaxies brighter than $R=20.5$, which corresponds to $M_r=-19+5\\log{h}$ at $z=0.25$.  We have measured morphologies using the {\\sc gim2d} surface-brightness fitting software of \\citet{Gim2d}, and the strengths of important spectral features, in particular the \\oii\\ emission line.  The properties of the ten clusters can be summarized as:  \\begin{itemize} \\item All ten clusters host a giant elliptical galaxy near the centre of the X-ray emission. None of the nine central galaxies for which we have a reliable spectrum show emission lines. \\item Apart from the double cluster Cl\\,1444, and the close pair of clusters Cl\\,1701 and Cl\\,1702, all cluster velocity dispersions are consistent with a single Gaussian, and they appear dynamically well-separated from the surrounding field.  Thus, they appear to be evolved systems in approximate dynamical equilibrium. \\item The measured velocity dispersions range from $\\sim 350$--$850$~km~s$^{-1}$, and are consistent with the local $L_X-\\sigma$ relation observed in larger samples \\citep[e.g.][]{XW}. \\item The fraction of cluster galaxies with \\ewoii\\ $>5$\\AA\\ (2$\\sigma$ confidence limit) is 22$\\pm 4$\\%.  The mean is \\ewoii$=3.2$\\AA, and the median is \\ewoii$=0.7$\\AA. There is no evidence for a significant correlation between \\ewoii\\ and either radius or density, apart from the lack of strong emission-line galaxies in the densest, central regions ($\\lesssim 0.1$ h$^{-1}$ Mpc).  Also, we do not measure a significant difference in the dynamics of the emission-line galaxies, relative to the rest of the population. \\item Disk-dominated galaxies ($B/T<0.4$) comprise 18$\\pm5$\\% of the sample within the central 0.4 h$^{-1}$ Mpc covered by our {\\it WFPC2} images.  Less than 25\\% (2/8) of these galaxies show significant emission.  The remainder, a population of ``anemic'' disk galaxies, are relatively isolated, regular spiral galaxies near $L^\\ast$, with smooth disks.  Such galaxies are rarely found in local, field samples, but are also seen, in similar abundance, in more massive clusters \\citep{P+99}. \\item No galaxies in our sample have a spectrum characteristic of a post-starburst or of truncated star formation.  Only four galaxies have \\hd$>4$\\AA\\ with at least 1$\\sigma$ significance, and all of these show nebular emission lines.  Thus there is no evidence that these cluster environments act to enhance star formation activity, even temporarily. \\end{itemize}  The distribution of \\ewoii\\ in these clusters is similar to that measured in the sample of \\citet{B+97}, which is comprised of clusters approximately an order of magnitude more massive.  Galaxies in both systems show low star formation rates even at projected surface densities as low as $\\sim 10 h^{2}$ Mpc$^{-2}$, where the fraction of spiral and irregular galaxies expected from the morphology-density relation is $\\sim 60$\\%.  The fact that star formation rates are so low even in these low-mass structures has important implications for understanding galaxy evolution in general. The phenomenon is not likely to be driven by extreme processes, such as ram-pressure stripping, or interaction-induced starbursts, which are expected to be important only in the richest clusters.  Rather it is something that operates in more commonplace environments, possibly groups in the infall regions of clusters \\citep{Kodama_cl0939,2dF-sfr}. Tracing the evolution of galaxies in these groups, therefore, may shed light on the processes responsible for the observed decline in the globally-averaged star formation rate of the Universe."
    },
    "0207/astro-ph0207287.txt": {
        "abstract": "{We present a detailed millimeter line study of the circumstellar environment of the low-luminosity Class 0 protostar IRAM 04191~+~1522 in the Taurus molecular cloud. New line observations demonstrate that the $\\sim 14000$~AU radius protostellar envelope is undergoing both extended infall and fast, differential rotation. Radiative transfer modeling of multitransition CS and C$^{34}$S maps indicate an infall velocity $v_{\\rm inf} \\sim 0.15$ km~s$^{-1}$ at $r \\sim 1500$ AU and $v_{\\rm inf} \\sim 0.1$ km~s$^{-1}$ up to $r \\sim 11000$~AU,  as well as a rotational angular velocity $\\Omega \\sim 3.9 \\times 10^{-13}$~ rad~s$^{-1}$, %$\\Omega \\sim 12$ km~s$^{-1}$.pc$^{-1}$, strongly decreasing with radius beyond 3500~AU down to a value %$\\Omega \\sim $~0.5--1~km~s$^{-1}$.pc$^{-1}$ $\\Omega \\sim $~1.5--3$\\times 10^{-14}$~rad~s$^{-1}$ at $\\sim $~11000~AU. Two distinct regions, which differ in both their infall and their rotation properties, therefore seem to stand out: the inner part of the envelope ($r \\simlt 2000-4000$ AU) is rapidly collapsing and rotating, while the outer part undergoes only moderate infall/contraction and slower rotation. These contrasted features suggest that angular momentum is conserved in the collapsing inner region but efficiently dissipated due to magnetic braking in the slowly contracting outer region. We propose that the inner envelope is in the process of decoupling from the ambient cloud and corresponds to the effective mass reservoir ($\\sim 0.5\\, M_\\odot $) from which the central star is being built. Comparison with the rotational properties of other %prestellar and protostellar objects in Taurus suggests that IRAM~04191 is at a pivotal stage between a prestellar regime of constant angular velocity enforced by magnetic braking and a dynamical, protostellar regime of nearly conserved angular momentum. The rotation velocity profile we derive for the inner IRAM~04191 envelope should thus set some constraints on the distribution of angular momentum on the scale of the outer Solar system at the onset of protostar/disk formation. ",
        "introduction": "\\label{intro}  \\subsection{The enigmatic onset of protostellar collapse} \\label{intro_background}  Despite recent progress, the initial conditions of star formation and the first phases of protostellar collapse remain poorly known \\citep[e.g.][ for reviews]{Myers99,Andre00}. In the standard theory of isolated, low-mass star formation \\citep[e.g.][]{Shu87}, the initial conditions correspond to essentially static singular isothermal spheres (SISs, which have $\\rho = (\\as^2/2\\pi G )r^{-2}$), assumed to be in slow, solid-body rotation \\citep[][ -- TSC84]{Terebey84}. This leads to a strictly constant mass accretion rate, $\\maccm \\sim \\as^3/G$ (where $\\as$ is the isothermal sound speed), and to a growth of the centrifugal disk as $R_{disk} \\propto  t^3$ (cf. \\citeauthor*{Terebey84}) during the protostellar accretion phase ($t > 0$). Other theoretical models exist, however, that predict a time-dependent accretion history if the collapse initial conditions are either not singular or not scale-free \\citep[e.g.][]{Foster93,Henriksen97,Basu97,Ciolek98,Hennebelle02}. Starting from realistic, finite-sized prestellar cores with $\\rho \\approx cte$ in their central region \\citep[c.f.][]{Ward99,Bacmann00,Alves01}, these models yield supersonic inward velocities close to the center prior to point mass formation (i.e. at $t<0$) and result in denser, nonequilibrium density distributions with strong differential rotation at the onset of the main accretion phase, i.e., at $t=0$. In these models, the accretion rate \\macc is initially significantly larger than in the Shu model, then quickly converges toward the standard $\\sim \\as^3/G$ value, and finally declines much below $\\as^3/G$ because of the finite reservoir of mass \\citep*[see, e.g.,][]{Foster93,Henriksen97}. Conservation of angular momentum during dynamical collapse at $t<0$ produces a differential rotation profile at $t=0$ \\citep[e.g. $\\Omega \\propto r^{-1}$ in the magnetically-controlled model of][]{Basu98}. This rotation profile in turn implies a more rapid growth of $R_{disk}$ initially (i.e., $R_{disk} \\propto  t$ at small $t>0$ in the Basu model) than in the \\citeauthor*{Terebey84} model.  Getting at a better, more quantitative knowledge of protostellar collapse is crucial, e.g., to gain insight into the origin of stellar masses and disk formation. Observationally, there are two complementary approaches to estimating the initial conditions of protostar formation. The first approach consists in studying the structure and kinematics of ``prestellar cores'' such as L1544 \\citep*[e.g.][]{Ward99,Tafalla98} , representative of times $t \\simlt 0$. The second approach, adopted here, is the detailed study of Class~0 accreting protostars observed at $t \\simgt 0$, such as IRAM~04191 (see \\S~\\ref{intro_iram04191}), which should still retain detailed memory of their initial conditions.  \\subsection{IRAM~04191: A  very young Class 0 protostar in Taurus} \\label{intro_iram04191}  The massive ($M_{tot} \\sim 1.5\\, M_\\odot $) dense core/envelope of the Class~0 object, IRAM 04191~+~1522 (hereafter IRAM 04191), was originally discovered in the millimeter dust continuum with the IRAM 30m telescope in the southern part of the Taurus molecular cloud \\citep[][ -- hereafter AMB99]{Andre99}. Follow-up observations revealed a highly collimated CO bipolar outflow (see Fig.~\\ref{n2h+flow}), a weak 3.6~cm VLA radio continuum source located at its center of symmetry, and spectroscopic evidence of spatially extended infall motions in the bulk of the envelope. These are typical attributes of a Class~0 protostar \\citep[][]{Andre93,Andre00}.  \\begin{figure} [!ht] \\resizebox{\\hsize}{!}{\\includegraphics[width=3.5cm,angle=0]{fig1.ps}} %\\resizebox{\\hsize}{!}{\\includegraphics[width=3.5cm,angle=0]{/home/storage/pandre/Belloche/018_99/Otf/I4191/outflow+core_n2h+_pdbi.ps}} \\caption[]{N$_2$H$^+$(1-0) integrated intensity map of the IRAM~04191 protostellar envelope overlaid on the CO(2-1) outflow map of \\citeauthor*{Andre99}. Both maps were taken with the IRAM 30m telescope. The N$_2$H$^+$(1-0) emission is integrated over the whole seven-component multiplet from 5.7 to 22.6 km~s$^{-1}$ and the contours go from 0.5 to 4.5 K~km~s$^{-1}$ by 0.5  K~km~s$^{-1}$. The blueshifted CO(2-1) emission ({\\it dotted contours}) is integrated over the 0-5 km~s$^{-1}$ velocity range, and the redshifted CO(2-1) emission ({\\it dashed contours}) is integrated over the 8-13 km~s$^{-1}$ velocity range; the CO(2-1) contours go from 5 to 21 K~km~s$^{-1}$ by 2 K~km~s$^{-1}$.\\\\ The insert in the bottom-left corner shows a 227~GHz dust continuum map of the inner envelope obtained at the IRAM Plateau de Bure interferometer (see \\S~\\ref{obs_pdbi}), with positive (solid) contours from $+1$ to $+6$ by 1 mJy/1.9\\arcsec -beam; the dotted contour is negative at $-1$~mJy/1.9\\arcsec -beam). \\\\ In both maps, the central star symbol at (0,0) marks the position of IRAM~04191 as determined by a 2D Gaussian fit to the PdBI 227~GHz continuum image. The other star symbol indicates the position of the Class I source IRAS~04191. } \\label{n2h+flow} \\end{figure}  The very high envelope mass to luminosity ratio of IRAM 04191 ($M_\\mathrm{env}^{< 4200 \\mathrm{ AU}}/L_\\mathrm{bol}\\simgt 3$ M$_\\odot$/L$_\\odot$) and its position in the $L_\\mathrm{bol}-T_\\mathrm{bol}$ evolutionary diagram ($L_\\mathrm{bol} \\sim 0.15$ L$_\\odot$ ~and $T_\\mathrm{bol} \\sim 18$ K) suggest an age $t \\sim 1-3 \\times 10^4$ yr since the beginning of the accretion phase (see \\citeauthor*{Andre99}). This is significantly younger than all of the IRAS candidate protostars of Taurus \\citep[e.g.][]{Kenyon93}, including L1527 which has  $t \\simlt 10^5$~yr \\citep[e.g.][]{Ohashi97a}. IRAM~04191 thus appears to be the youngest accreting protostar known so far in Taurus, although the collapsing protostellar condensation MC~27 discovered by \\citet{Onishi99} may be in a comparable evolutionary state.  As IRAM~04191 is particularly young, nearby (d $= 140$ pc), and relatively isolated, the study of its velocity structure based on molecular line observations provides a unique opportunity to set constraints on collapse models. This is especially true since the viewing angle is favorable. The CO(2--1) outflow map of \\citeauthor*{Andre99} (see Fig.~\\ref{n2h+flow}) shows well separated outflow lobes with almost no overlap between blue-shifted and red-shifted emission, indicating that the flow lies out of the plane of the sky at an intermediate inclination angle \\citep[e.g.][]{Cabrit90}. In addition, \\citeauthor*{Andre99} estimate an aspect ratio of $\\sim 0.65$ for the circumstellar dust/N$_2$H$^+$ envelope, whose major axis is perpendicular to the outflow axis (cf. Fig.~\\ref{n2h+flow}). Both characteristics are consistent with an inclination angle of the outflow axis to the line of sight of $i \\sim 50 \\degr$.  Here, we present and discuss the results of a comprehensive set of molecular line observations toward IRAM~04191. The layout of the paper is as follows. Sect.~\\ref{obs_set} summarizes observational details. Sect.~\\ref{obs_ana} interprets the observations in terms of infall, rotation, and outflow motions in the envelope. We then model the observed spectra using radiative transfer simulations computed in 1D spherical geometry with radial infall motions (Sect.~\\ref{simul_1d}) and in 2D axial geometry with both infall and rotational motions (Sect.~\\ref{simul_2d}). Sect.~\\ref{discuss} compares the derived constraints on the velocity structure of the IRAM~04191 envelope with the predictions of collapse models. Our conclusions are summarized in Sect.~\\ref{concl}.  %%%%%%%%%%%%%%%%%% %%%% ",
        "conclusions": "\\label{concl}  We have carried out a detailed study of the structure and kinematics of the envelope surrounding the Class~0 protostar IRAM~04191 in Taurus. Our main results and conclusions are as follows:  \\begin{enumerate}  \\item Extended, subsonic infall motions with $v_\\mathrm{inf} \\sim 0.5\\, a_s \\sim 0.1$ km~s$^{-1}$, responsible for a marked `blue infall asymmetry' in self-absorbed CS and H$_2$CO lines, are present in the bulk of the envelope, up to at least $r_\\mathrm{i,o} \\sim 10000-12000$ AU. The observations are also consistent with larger infall velocities scaling as $v_\\mathrm{inf} \\propto r^{-0.5}$ in an inner region of radius $r_\\mathrm{i} \\approx 2000$~AU. The corresponding mass infall rate is estimated to be $\\dot{M}_\\mathrm{inf} \\sim 2-3 \\times a_s^3/G \\sim 3 \\times 10^{-6}$~M$_\\odot$~yr$^{-1}$.  \\item The protostellar envelope is differentially rotating with an angular velocity profile $\\Omega \\propto r^{-2.5 \\pm 0.5}$ between $r_\\mathrm{m} \\approx 3500$~AU and $r_\\mathrm{m,o} \\sim 7000$~AU. The rotation profile is shallower, albeit more poorly constrained, in the inner $r < r_\\mathrm{m}$ region, i.e., $\\Omega \\propto r^{-0.9 \\pm 0.4}$. The angular velocity is estimated to be $\\Omega \\sim 12$~km~s$^{-1}$~pc$^{-1}$ at $r \\sim 3500$~AU and only $\\Omega \\simlt 0.5- 1$~km~s$^{-1}$~pc$^{-1}$ at $r \\sim 11000$~AU. The present value of the centrifugal radius is estimated to be less than 400 AU.  \\item The extended infall velocity profile is inconsistent with the inside-out collapse picture of \\citet{Shu87} and only marginally consistent with isothermal collapse models starting from marginally stable equilibrium Bonnor-Ebert spheres. The latter tend to produce somewhat faster infall velocities than are observed.  \\item The contrast observed between the (steeply declining) rotation velocity profile and the (flat) infall velocity profile beyond $r_\\mathrm{m} \\approx 3500$~AU suggests that angular momentum is {\\it not} conserved in the outer envelope. This is difficult to account for in the context of non-magnetic collapse models.  \\item Based on a qualitative comparison with magnetic ambipolar diffusion models of cloud collapse (e.g. \\citeauthor*{Basu94}), we propose that the rapidly rotating inner envelope of IRAM~04191 corresponds to a magnetically supercritical core decoupling from an environment still supported by magnetic fields and strongly affected by magnetic braking. In this view, the outer ($r_\\mathrm{m} < r < r_\\mathrm{m,o}$) envelope represents a transition region between the forming protostar and the slowly rotating ambient cloud. Although published ambipolar diffusion models have difficulty explaining supercritical cores as small as $R_{crit} \\sim 3500$~AU, we speculate that more elaborate versions of these models, including the effects of MHD turbulence in the outer envelope, would be more satisfactory.  \\item Interestingly, the steepening of $\\Omega(r)$ in IRAM~04191 occurs at a radius comparable to the $\\sim 5000$~AU scale inside which the specific angular momentum of Taurus dense cores appears to be conserved \\citep[cf.][ and Fig.~\\ref{joverm}]{Ohashi97b}. Our results therefore support \\citet{Ohashi97b}'s proposal that $r \\sim 5000$~AU represents the typical size scale for dynamical collapse in Taurus. More generally, we suggest that the rotation/infall properties observed here for IRAM~04191 are representative of the physical conditions prevailing in isolated protostellar envelopes shortly ($\\sim 10^4$~yr) after point mass formation.  \\end{enumerate}"
    },
    "0207/astro-ph0207406_arXiv.txt": {
        "abstract": "We have undertaken a study of radio and infrared molecular--line emission towards several SNRs in order to investigate molecular signatures of SNR shocks, and to test models for \\oh maser production in SNRs.  Here we present results on G349.7+0.2. ",
        "introduction": "Shocks driven by a supernova remnant (SNR) into a molecular cloud compress, accelerate and heat the gas, and can contribute to disruption of the molecular cloud. Despite the close association between SNRs and molecular clouds (Huang \\& Thaddeus 1986), establishing unambiguous cases of interaction has been difficult, mostly because of confusion resulting from unrelated clouds along the line of sight to the SNR and the possibility of chance alignment. Studies of SNR/molecular cloud interactions have been re--energized by the recent discovery of an association between \\oh maser emission and SNRs (see Koralesky et al.~1998, and references therein). This maser emission requires very specific conditions: gas density of $\\sim$\\e{5}\\cm{-3}, gas temperature of 50--125 K, and OH column density of $\\sim$\\e{16}--\\e{17}\\cm{-2} (Lockett, Gauthier \\& Elitzur 1999). These conditions can be satisfied in the cooling gas behind a non--dissociative C--shock, irradiated by the X--ray flux from the SNR interior (Wardle 1999). The presence of \\oh masers in SNRs provides two crucial pieces of information for studying SNR shocks in molecular clouds: 1) the velocity of a cloud associated with an SNR, and 2) the precise location of the on--going interaction.   \\begin{figure} \\centerline{\\epsfig{file=lazendicj_1.ps,width=7cm}} \\caption{G349.7+0.2 is located at a distance of $\\sim$23 kpc, and is the third brightest galactic SNR at radio wavelengths. Five OH(1720 MHz) masers (marked with crosses) have been detected along the bright emission ridge of the SNR (Frail et al.~1996). The \\co\\ 1--0 emission obtained with SEST is shown, integrated between 10 and 20\\kms\\  (contours), and overlaid on the 18--cm radio continuum greyscale image obtained with ATCA with resolution of 9$\\arcsec\\times5\\arcsec$. The contour levels are: 16, 24, 32, 48, 63, 71, 78 K\\kms. The dots mark the grid positions of the SEST observations.} \\label{fig-candy-co} \\end{figure} ",
        "conclusions": ""
    },
    "0207/astro-ph0207630_arXiv.txt": {
        "abstract": "We have used the NRAO Very Large Array (VLA) in conjunction with the Very Long Baseline Array (VLBA) Pie Town antenna as a real-time interferometer system to measure the size of the extragalactic source \\j18 as a function of frequency from 1285 to 4885 MHz.  These observations were made in an attempt to determine the effect interstellar scattering has on the observed sizes of OH (1720 MHz) masers in the nearby (d $=$ 2.5 kpc) supernova remnant \\wtens.  The observations clearly show that \\j18 displays angular broadening due to turbulence in the Galaxy's interstellar medium.  The minimum distance of the nearby (two arcminutes from \\j18) pulsar PSR B1758$-$23 is constrained to be 9.4$\\pm$2.4 kpc.  This value is based on both the measured size of 220 mas for \\j18 at 1715 MHz and the temporal broadening of the pulsar.  A single thin scattering screen along the line of sight to the \\wte OH(1720 MHz) masers must be at 4.7$\\pm$1.2 kpc for this minimum pulsar distance. The screen may be placed closer to the Earth, but for reasonable values of the pulsar distance (i.e., the pulsar is within the Galaxy), this choice leads to a negligible scattering contribution to the sizes of the masers.  Thus the OH(1720 MHz) masers, at a distance of 2.5$\\pm$0.7 kpc, are unaffected by interstellar scattering, and the measured maser sizes must be intrinsic.  Our measured upper limits to the size of the pulsar itself are consistent with the distance estimates to the pulsar and the scattering screen. ",
        "introduction": "The \\wte supernova remnant (SNR) lies in the direction of the Galactic Center: $(l,b)=(6.8,-0.06)$. The distance to \\wte is somewhat uncertain.  Since the distance to the SNR is important to the conclusions of the current study, it is useful to examine previous estimates of the distance found in the literature. Many authors have used the $\\Sigma-D$ relation to derive a distance to \\wte near $\\sim$2 kpc (Milne 1970; Clark \\& Caswell 1976; Goudis 1976; Milne 1979); however, as discussed by Kaspi \\etal (1993), this method is extremely uncertain.  The furthest distance estimated in the literature is 3.6 kpc by Lozinskaya (1974), based on H$\\alpha$ measurements and assuming an LSR velocity near +18\\kmss for \\wtens.  Velazquez \\etal (2002) estimate a distance of 1.9$\\pm$0.3 kpc adopting +7\\kmss for the LSR velocity, and assuming a standard circular rotation model (this is the near kinematic distance; 15 kpc is the far kinematic distance). Frail, Kulkarni, \\& Vasisht (1993) adopt a distance of 3 kpc, based on HI absorption, assuming an LSR velocity of 17.6\\kms. While the weight of the evidence for distance estimates to \\wte favors a lower distance ($\\sim$ 2 kpc; the argument of Velazquez \\etal 2002 for an LSR velocity of +7\\kms is especially strong), we adopt a conservative estimate for the distance to \\wte of 2.5$\\pm$0.7 kpc.  The 60,000 year old pulsar PSR~B1578$-$23 lies in the same direction as \\wte but is located outside the SNR, approximately three arcminutes to the north of its bright radio continuum edge.  A radio continuum source, \\j18 (also known as 1758$-$231), lies within two arcminutes of this pulsar. \\j18 is assumed to be an extragalactic source, coincidentally along the line of sight to the pulsar, based on measurements of neutral hydrogen absorption (Frail, Kulkarni, \\& Vasisht 1993).  Using observations of the pulsar and the neighboring extragalactic source, Frail et al. argued in favor of the association of the pulsar and the SNR.  Kaspi et al. (1993) disagreed, suggesting that the pulsar was much more distant.  The discussion of Frail \\etal was based upon the similarity of HI absorption profiles toward the pulsar and extragalactic source. Kaspi et al. base their conclusion on the high dispersion measure of the pulsar (1074 pc cm$^{-3}$), and the argument that no compact HII region or cloud appears to lie along the line of sight to the pulsar which could account for such a high dispersion measure.  Several OH (1720 MHz) masers associated with the \\wte SNR (Frail, Goss, \\& Slysh 1994) have been studied at high angular resolution with the Very Long Baseline Array (VLBA) and MERLIN by Claussen \\etal (1999a).  Though there have been only a few OH(1720 MHz) masers in \\wte observed at 10 mas resolution, those masers are resolved and have sizes ranging from 50 to 100 milliarcseconds (mas).  Claussen \\etal discussed the effects of interstellar scattering on the sizes of the masers and the constraints that could be placed on such scattering sizes based on available information about the size of \\j18, the pulsar, and the pulse broadening of the pulsar. Interstellar scattering of an extragalactic source results in angular broadening while interstellar scattering of a pulsar's radio emission can also result in pulse broadening, an increase in the apparent width of a pulsar's average pulse profile beyond its intrinsic width.  The degree of angular broadening for the masers and the extragalactic source and the degree of pulse broadening for the pulsar all depend in different ways upon the relative geometry of the observer, scattering material, and sources.  Measurements of these scattering effects can constrain the distribution of the scattering material and can be used to estimate, for example, the unscattered sizes of the masers.  In this paper, we present angular broadening measurements of the extragalactic source \\j18 as a function of observing frequency from 1285 MHz to 4885 MHz. We used the NRAO Very Large Array (VLA) in its most extended configuration ({\\bf A}) linked, in real time, with the Pie Town (Pt) antenna of the Very Long Baseline Array (VLBA). The questions to be addressed by these measurements are (1) the relative distances of the pulsar and the supernova remnant (which we assume harbors the OH masers), (2) the effects of scattering on the observed size of the masers and the pulsar, and (3) the distance to the scattering screen. ",
        "conclusions": "We have measured the size of the extragalactic source \\j18 at six frequencies, and find that our size measurements are consistent with angular broadening due to turbulence in the interstellar medium.  Together with the size measurements of OH(1720 MHz) masers, assumed to be associated with the \\wte supernova remnant (Claussen et al. 1999a), and the temporal broadening of the pulsar B1758$-$23 (Frail, Kulkarni, \\& Vasisht 1993; Kaspi et al. 1993), the results of the new size measurements of \\j18 imply:  $\\bullet$  The minimum distance for the pulsar is 9.4$\\pm$2.4  kpc and it is {\\it not} associated with the SNR.  The distance to a single scattering screen for this minimum pulsar distance is 4.7$\\pm$1.2 kpc, and is also beyond the SNR.  $\\bullet$  The scattering screen can be moved placed closer to the Earth by increasing the pulsar distance.  Even with an extreme distance to the pulsar, this leads to a negligible scattering contribution to the size of the OH masers.  $\\bullet$  In either of the two possibilities above, the maser sizes must be little affected by interstellar scattering, and the sizes measured for the masers must therefore be intrinsic to the masers.  $\\bullet$ When the pulsar is placed at the minimum distance of 9.4 kpc, the scattering size of the pulsar, estimated from the \\j18 size measurements, is about 110 mas at 1720 MHz, consistent with our upper limit to the pulsar size of 500 mas."
    },
    "0207/astro-ph0207540_arXiv.txt": {
        "abstract": "We report our discovery of the likely near-infrared counterpart to the anomalous X-ray pulsar (AXP) 1E~1048.1$-$5937, using observations from the 6.5~m Baade (Magellan~I) telescope in Chile.  We derived a precise position for the X-ray source using archival data from the {\\em Chandra X-Ray Observatory}.  This position is inconsistent with a position reported earlier from {\\em XMM-Newton}, but we show that the originally reported {\\em XMM-Newton} position suffered from attitude reconstruction problems. Only two of the infrared objects in a 17\\arcsec$\\times$17\\arcsec\\ field containing the target have unusual colors, and one of these has colors consistent with those of the identified counterparts of two other AXPs.  The latter object is also the only source detected within the 0$\\farcs$6 {\\em Chandra} error circle, and we identify it as the counterpart to 1E~1048.1$-$5937. This is the first AXP counterpart detected in multiple infrared bands, with magnitudes $J$=21.7(3), $H$=20.8(3), and $K$=19.4(3). There is marginal evidence for spectral flattening at longer wavelengths. ",
        "introduction": "The anomalous X-ray pulsars (AXPs) are a small group of neutron stars with spin periods falling in a narrow range (6--12~s), very soft X-ray spectra, and with no evidence of binary companions (see Mereghetti et al. 2002 for a recent review).  Their X-ray luminosities greatly exceed the power available from spin-down of the pulsars.  These objects are thus believed to be isolated neutron stars either having extremely strong ($\\sim10^{14}$ G) surface magnetic fields (``magnetars'') or accreting from a residual accretion disk.  X-ray bursts detected from the AXPs 1E~1048.1$-$5937 (Gavriil, Kaspi, \\& Woods 2002) and 1E~2259+586 (Kaspi \\& Gavriil 2002) have strengthened an already suspected connection to the soft gamma-ray repeaters (SGRs).  The recent identification of optical/infrared counterparts to the AXPs 4U~0142+61 (Hulleman, van Kerkwijk, \\& Kulkarni 2000) and 1E~2259+586 (Hulleman et al. 2001) and the discovery of optical pulsations from 4U~0142+61 (Kern \\& Martin 2002) favor the magnetar scenario for these objects.  In this paper, we report on the detection of a third AXP counterpart, and the first measurement of infrared colors for one of these objects.  The AXP 1E~1048.1$-$5937 ($l$=288\\fdg3, $b$=$-$0\\fdg5) was discovered serendipitously in 1979 during an {\\em Einstein} observation of the Carina Nebula (Seward, Charles, \\& Smale 1986).  It has since been extensively observed by a series of X-ray missions including {\\em EXOSAT} (Seward et al. 1986), {\\em Ginga} (Corbet \\& Day 1990), {\\em ROSAT} (Mereghetti 1995), {\\em ASCA} (Corbet \\& Mihara 1997; Paul et al. 2000), {\\em RXTE} (Mereghetti, Israel, \\& Stella 1998; Baykal et al. 2000; Kaspi et al. 2001; Gavriil et al. 2002), and {\\em XMM-Newton} (Tiengo et al. 2002).  These observations establish the source as a 6.4~s pulsar that is spinning down, with an absorbed power-law + blackbody X-ray spectrum and occasional SGR-like bursts. The high column density to the source indicates that it lies behind the Carina Nebula, setting a lower limit on its distance of $d\\gtrsim 2.8$~kpc (Seward et al. 1986; \\\"Ozel, Psaltis, \\& Kaspi 2001). Previous searches for an optical counterpart have been unsuccessful (Seward et al. 1986; Mereghetti, Caraveo, \\& Bignami 1992) with the resultant limiting magnitudes $BVR \\gtrsim 24.3$ (Israel, Mereghetti, \\& Stella 2001). ",
        "conclusions": "We have identified the likely near-infrared counterpart to 1E~1048.1$-$5937, on the basis of coincidence with the {\\em Chandra} X-ray source position, unusual colors, and similarity to the  \\begin{figure*}[t] \\centerline{\\epsfig{file=table.eps}} \\end{figure*}  \\noindent two other known AXP counterparts. We note, however, that this conclusion relies strongly on our $i'$ non-detection.  A sufficiently faint $i'$ magnitude could push this candidate out of the expected AXP region in the bottom panel of Figure~2. This is the first AXP counterpart with infrared detections in more than one band, allowing the intrinsic infrared colors to be inferred.  Assuming $A_V=5.8$ as determined in \\S4, the dereddened infrared magnitudes of the counterpart are $J_0=20.1$, $H_0=20.0$, and $K_{s0}=18.8$.  We plot these magnitudes and limits in Figure~3.  For comparison, we plot the inferred spectral shape of the 4U~0142+61 counterpart (Hulleman et al.~2000; Hulleman 2002). While there is marginal evidence for a spectral flattening or even a turnover at the $K_s$ band, the data are also consistent with a flat $\\nu F_\\nu$ infrared spectrum at the 1$\\sigma$ level.  We note that there is clear evidence for spectral flattening in the infrared in the spectrum of 4U~0142+61 (Hulleman 2002).  It would thus be of great interest to search for emission at longer wavelengths to constrain the spectral shape of AXPs further.  While there has been relatively little theoretical work presented on the expected AXP optical/infrared emission in the magnetar scenario, we note that models involving X-ray illumination of a ``fallback'' accretion disk predict a spectral turnover leading to strong emission in the infrared and submillimeter bands (e.g., Chatterjee, Hernquist, \\& Narayan 2000; Perna, Hernquist, \\& Narayan 2000)."
    },
    "0207/astro-ph0207295_arXiv.txt": {
        "abstract": "In a recent paper, McLure \\& Jarvis (2002, astro-ph/0204473, v.1) reanalyze AGN broad-line reverberation measurements presented in Kaspi et al. (2000). They find the broad-line region size and the AGN luminosity at 5100~\\AA\\ can be related by $R_{BLR}\\propto L_{5100}^{0.5}$. This differs from the result of Kaspi et al., who found $R_{BLR}\\propto L_{5100}^{0.7}$, a departure from expectations for a constant ionization parameter. In this note I investigate the source of the discrepancy. McLure \\& Jarvis use quasar luminosities based on the single-epoch, multichannel-spectrophotometer measurements of Neugebauer et al. (1987), obtained in 1980. I show that the Neugebauer et al. fluxes are systematically higher, by a constant flux offset, compared with the multi-epoch CCD measurements of Kaspi et al., obtained concurrently with the echo-mapping radii. In addition, McLure \\& Jarvis erred when converting these fluxes to luminosities. The two effects compound to a typical factor 2 (0.3 dex) overestimate of the luminosities of the PG quasars in the sample. Since McLure \\& Jarvis did adopt the Seyfert luminosities in Kaspi et al. (2000), they obtained a slope that is flatter by 0.2. ",
        "introduction": "Reverberation mapping of active galactic nuclei (AGNs) exploits the light-travel-time delay between continuum variations and the response of broad emission lines to measure the sizes of AGN broad-line regions (BLRs). This size can then be used, together with the estimated velocities of the BLR gas to derive masses for the central black holes (See Netzer \\& Peterson 1996, for a review.) The size-luminosity and mass luminosity relations of AGNs may shed new light on understanding these objects and the connections between the black holes in AGNs and those found in normal nearby galaxies.  Kaspi et al. (2000) measured reverberation sizes, $R_{BLR}$, for 17 PG quasars between 1991--1998 using the Wise Observatory 1m telescope and the Steward Observatory 2.3m telescope, with a combination of CCD spectrophotometry and photometry. The time averaged flux for each quasar was derived from 20-70 epochs per quasar. The rest-frame 5100 \\AA\\ continuum flux densities were measured from the $\\sim 10$ \\AA\\ resolution spectra while minding potential systematics such as broad-emission-line wings and atmospheric absorptions in the higher-redshift objects. Following Galactic extinction corrections (small for most of the objects) the absolute fluxes were used to calculate 5100 \\AA\\ luminosities, $\\lambda L_{\\lambda}(5100 {\\rm \\AA})$. The sizes and luminosities of Seyfert galaxies with reverberation measurements were compiled from previously published works. A regression analysis taking into account the errors in both $R_{BLR}$ and $\\lambda L_{\\lambda}(5100{\\rm \\AA})$ showed that $R_{BLR} \\propto [\\lambda L_{\\lambda}(5100{\\rm \\AA})]^{0.70\\pm 0.03}$. This was a surprising result, since it has long been speculated that the dependence would be to the power 0.5. Such a scaling would lead to an ionization parameter (ratio of ionizing photon density to electron density) at the surface of the BLR clouds that is independent of luminosity, and would explain the similarity of AGN spectra over many orders of magnitude in luminosity.  In a recent paper, McLure \\& Jarvis (2002) examine to what degree UV observables, namely 3000\\ \\AA\\ luminosities and Mg~II line velocities, can serve the purpose of the optical observables -- 5100 \\AA\\ luminosities and H$\\beta$ widths -- used to date in AGN black hole estimates. In the course of their work, they re-analyze the data presented by Kaspi et al. (2000) and conclude that actually $R_{BLR} \\propto [\\lambda L_{\\lambda}(5100{\\rm \\AA})]^{0.50\\pm0.02}$ is the  best fit, contrary to the result of Kaspi et al., and consistent with the expectations for a constant ionization parameter.  Vestergaard (2002) has carried out an analysis along similar lines. After studying the various systematics that can affect the slope determination, she concludes that the best estimate (her equation A5) is $R_{BLR} \\propto [\\lambda L_{\\lambda}(5100)]^{0.66\\pm0.09}$, consistent with the result of Kaspi et al. (2000). The statement by McLure \\& Jarvis (2002), that Vestergaard (2002) found a slope consistent with 0.5 but decided to adopt a slope of 0.7 anyway, is incorrect.  In this Note, I investigate the source of the discrepancy between Kaspi et al. and McLure \\& Jarvis. I show that it arises, first, due to the adoption by McLure \\& Jarvis of old, single-epoch, and systematically higher fluxes, but only for the high luminosity part of the sample; and second, due to incorrect conversion from flux to luminosity. This is confirmed and corrected in a revised version of their paper (R. McLure, private communication). ",
        "conclusions": ""
    },
    "0207/astro-ph0207248_arXiv.txt": {
        "abstract": "We present the results of a simulation of a wind-driven non-linear gravity wave breaking on the surface of a white dwarf. The ``wind'' consists of H/He from an accreted envelope, and the simulation demonstrates that this breaking wave mechanism can produce a well-mixed layer of H/He with C/O from the white dwarf above the surface. Material from this mixed layer may then be transported throughout the accreted envelope by convection, which would enrich the C/O abundance of the envelope as is expected from observations of novae. ",
        "introduction": "Classical novae result from the ignition (and subsequent explosive thermonuclear burning) of a ($\\sim 10^4$ m) layer of hydrogen-rich material that has accreted from a main sequence companion onto the surface of a white dwarf \\cite{truran82,shara89,starrfield89,livio94}. Observed abundances and explosion energies estimated from observations indicate that there must be significant mixing of the heavier material of the C/O or O/Ne white dwarf into the lighter accreted material (H/He). This mixing is critical because otherwise hydrogen burning would be too slow to reproduce observed nova characteristics in outburst. Further, without this mixing it is difficult to understand the observed abundances of intermediate-mass nuclei in the ejecta. Accordingly, nova models must incorporate a mechanism that will dredge up the heavier white dwarf material~\\cite[and references therein]{rosner01}.  A recently proposed mixing mechanism is the breaking of non-linear resonant gravity waves at the C/O surface \\cite{rosner00,alexakis01,rosner01}. The gravity waves, driven by the ``wind'' of accreted material, can break, forming a layer of  well-mixed material. This mixed layer may then be transported upward by convection, thereby enriching the accreted material. Because the length scale of this mixed layer may be very small (much smaller than the length scale of convection), previous precursor simulations have not captured this effect.  In this manuscript, we present a simulation of a wind-driven non-linear gravity wave breaking on the surface of a white dwarf. The simulation was performed with FLASH, a parallel, adaptive-mesh simulation code for the compressible, reactive flows found in many astrophysical environments~\\cite{fryxell00,calder00}. This simulation is part of an ongoing study of this mechanism to assess its efficacy for mixing white dwarf material with envelope material. \\begin{figure} \\includegraphics[height=.3\\textheight]{f1.eps} \\caption{Results from the simulation of a wind-driven gravity wave. The left panel shows the potential energy of the wave vs.\\ time. Also shown is the potential energy predicted by the linear theory. The right panel shows the mixed C/O mass per unit area vs.\\ time. \\label{fig:mixpe}} \\end{figure} ",
        "conclusions": "The simulation presented in this manuscript demonstrates the proposed breaking wind-driven non-linear gravity wave mixing mechanism. The results show the development of well-mixed zone just above the surface of the white dwarf. This simulation will be one part of a study of this mixing mechanism investigating effects of wind profiles and speeds. Complete details of the study will appear in \\citet{alexakis02}.  The expectation is that the breaking of non-linear gravity waves on the surface of the white dwarf will lead to a thin well-mixed layer of material that may then be transported throughout the the envelope by convection. This study investigating the mixing mechanism should provide quantitative information about the mixing rate that will allow for the development of subgrid models that may be applied to multidimensional convection simulations to study the enrichment of the envelope. A preliminary simulation of this kind is also presented in this volume ~\\cite{dursi02}. \\begin{figure} \\includegraphics[height=.27\\textheight]{f3.eps} \\caption{Images of the log of density with velocity vectors at later times during the simulation. The units of density are g/cm$^3$ and the length of the velocity arrows is proportional to the magnitude of the velocity, with a maximum of $2 \\times 10^8$ cm/s. The left panel shows the simulation at t = 0.030 s, and the right panel shows the simulation at t = 0.045 s. \\label{fig:late}} \\end{figure}  \\begin{theacknowledgments} This work is supported in part by the U.S. Department of Energy (DOE) under Grant No. B341495 to the Center for Astrophysical Thermonuclear Flashes at the University of Chicago. J.~W.~Truran acknowledges partial support from DOE grant DE-FG02-91ER40606.  L.~J.~Dursi is supported by the Krell Institute CSGF. K.~Olson acknowledges partial support from NASA grant NAS5-28524. M. Zingale acknowledges support from the Scientific Discovery through Advanced Computing (SciDAC) program of the DOE, grant number DE-FC02-01ER41176. Additional details about the project and information about requesting a copy of FLASH may be found at \\url{http://flash.uchicago.edu}.  \\end{theacknowledgments}"
    },
    "0207/astro-ph0207281_arXiv.txt": {
        "abstract": "The neutrino flux close to a supernova core contributes substantially to neutrino refraction so that flavor oscillations become a nonlinear phenomenon. One unexpected consequence is efficient flavor transformation for anti-neutrinos in a region where only neutrinos encounter an MSW resonance or vice versa.  Contrary to previous studies we find that in the neutrino-driven wind the electron fraction $Y_e$ always stays below~0.5, corresponding to a neutron-rich environment as required by r-process nucleosynthesis. The relevant range of masses and mixing angles includes the region indicated by LSND, but not the atmospheric or solar oscillation parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207554_arXiv.txt": {
        "abstract": "We present a new grid of ionizing fluxes for O and Wolf-Rayet stars for use with evolutionary synthesis codes and single star H\\two\\ region analyses. A total of 230 expanding, non-LTE, line-blanketed model atmospheres have been calculated for five metallicities (0.05, 0.2, 0.4, 1 and 2\\,Z$_\\odot$) using the WM-basic code of Pauldrach et al. (2001) for O stars and the {\\sc cmfgen} code of Hillier \\& Miller (1998) for W-R stars. The stellar wind parameters are scaled with metallicity for both O and W-R stars. We compare the ionizing fluxes of the new models with the CoStar models of Schaerer \\& de Koter (1997) and the pure helium W-R models of Schmutz, Leitherer \\& Gruenwald (1992).  We find significant differences, particularly above 54\\,eV, where the emergent flux is determined by the wind density as a function of metallicity.  The new models have lower ionizing fluxes in the He\\one\\ continuum with important implications for nebular line ratios.  We incorporate the new models into the evolutionary synthesis code Starburst99 (Leitherer et al. 1999) and compare the ionizing outputs for an instantaneous burst and continuous star formation with the work of Schaerer \\& Vacca (1998; SV98) and Leitherer et al. (1999).  The changes in the output ionizing fluxes as a function of age are dramatic. We find that, in contrast to previous studies, nebular He\\two\\ $\\lambda4686$ will be at, or just below, the detection limit in low metallicity starbursts during the W-R phase. The new models have lower fluxes in the He\\one\\ continuum for $Z \\ge 0.4$\\,Z$_\\odot$ and ages $\\le 7$\\,Myr because of the increased line blanketing.  We test the accuracy of the new model atmosphere grid by constructing photoionization models for simple H\\two\\ regions, and assessing the impact of the new ionizing fluxes on important nebular diagnostic line ratios.  For the case of an H\\two\\ region where the ionizing flux is given by the WM-basic dwarf O star grid, we show that He\\one\\ $\\lambda5786$/H$\\beta$ decreases between 1 and 2\\,Z$_\\odot$ in a similar manner to observations (e.g. Bresolin et al. 1999).  We find that this decline is caused by the increased effect of line blanketing above solar metallicity. We therefore suggest that a lowering of the upper mass limit at high abundances is not required to explain the diminishing strength of He\\one\\ $\\lambda5786$/H$\\beta$, as has been suggested in the past (e.g. Shields \\& Tinsley 1976; Bresolin et al. 1999). For an H\\two\\ region where the ionizing flux is provided by an instantaneous burst of total mass $10^6$\\,M$_\\odot$, we plot the softness parameter $\\eta^\\prime$ against the abundance indicator $R_{23}$ for ages of 1--5\\,Myr. The new models are coincident with the observational data of Bresolin et al. (1999), particularly during the W-R phase, unlike the previous models of SV98 which generally over-predict the hardness of the ionizing radiation.  The new model grid and updated Starburst99 code can be downloaded from http://www.star.ucl.ac.uk/starburst. ",
        "introduction": "Evolutionary synthesis codes are commonly used to derive the properties of young, unresolved stellar populations from observations at various wavelengths.  In the satellite ultraviolet, the stellar wind spectral features of a population of massive stars can be synthesized to provide information on the star formation rates, the slope of the initial mass function (IMF), and ages (e.g. Robert, Leitherer \\& Heckman 1993; Leitherer, Robert \\& Heckman 1995; de Mello, Leitherer \\& Heckman 2000). In the optical region, the properties of integrated stellar populations are usually derived indirectly from their total radiative energy outputs using nebular diagnostic line ratios (e.g. Garc\\'\\i a-Vargas, Bressan \\& D\\'\\i az 1995; Stasi\\' nska \\& Leitherer 1996; Stasi\\' nska, Schaerer \\& Leitherer 2001).  This method uses the theoretical ionizing fluxes from a population of massive stars as a function of age as input into a photoionization code.  The ability of evolutionary synthesis models to predict correctly the properties of a young stellar population from nebular emission line ratios therefore depends heavily  on the accuracy of the evolutionary and atmospheric models developed for single massive stars.  Early attempts to model stellar populations using evolutionary synthesis coupled with photoionization codes relied mainly on Kurucz (1992) plane-parallel LTE model atmospheres (e.g.  Garc\\'\\i a-Vargas et al. 1995). Gabler et al. (1989), however, showed that the presence of a stellar wind has a significant effect on the emergent ionizing flux in the neutral and ionized helium continua of O stars. Non-LTE effects depopulate the ground state of He\\two, leading to a decrease in the bound-free opacity above 54\\,eV, and hence a larger flux in the He\\two\\ continuum by up to 3--6 orders of magnitude. The emergent spectrum is flattened in the region of the He\\one\\ continuum and has a higher flux due to the presence of the stellar wind.  Stasi\\' nska \\& Leitherer (1996) were the first to use expanding non-LTE atmospheres for stars with strong winds to construct photoionization models for evolving starbursts. They used the grid of pure helium, unblanketed non-LTE W-R atmospheres from Schmutz, Leitherer \\& Gruenwald (1992) to represent evolved stars with strong winds, and Kurucz (1992) models for hot stars close to the main sequence. Line-blanketed, expanding non-LTE models covering the main sequence evolution of O stars were first introduced by Stasi\\' nska \\& Schaerer (1997) who studied the effect of using the CoStar models of Schaerer \\& de Koter (1997) on the ionization of single star H\\two\\ regions. They found that higher ionic ratios were obtained in comparison to Kurucz models with the same stellar temperatures. Schaerer \\& Vacca (1998) used the CoStar and Schmutz et al. (1992) grids to construct evolutionary synthesis models for young starbursts. They predicted strong nebular He\\two\\ $\\lambda4686$ in low metallicity starbursts containing W-R stars. The same evolutionary synthesis models were used by Stasi\\' nska, Schaerer \\& Leitherer (2001) combined with photoionization models to analyse the emission line properties of H\\two\\ galaxies.  One major question arises: how realistic are the ionizing fluxes being used in the evolutionary synthesis studies outlined above? To address this, there have been numerous studies aimed at empirically testing the accuracy of hot star model atmospheres by analysing H\\two\\ regions containing single stars with well defined spectral types and effective temperatures.  Schaerer (2000) reviews the success of non-LTE model atmospheres in reproducing nebular diagnostic line ratios.  Esteban et al. (1993) examined the accuracy of the W-R grids by photoionization modelling of nebulae ionized by single W-R stars. They found that the lack of line blanketing was most important for the coolest W-R stars. In recent years, new computational techniques have allowed line-blanketed W-R atmospheres to be calculated for a few stars (e.g. Schmutz 1997; de Koter, Heap \\& Hubeny 1997; Hillier \\& Miller 1998).  Crowther et al. (1999) tested these models by seeking to find a consistent model that reproduced the stellar and nebular parameters of a cool WN star.  They found that line blanketing plays a significant role in modifying the W-R ionizing output although the models of de Koter et al. (1997) and Hillier \\& Miller (1998) predict quite different ionizing flux distributions below the He\\one\\ edge.  Bresolin, Kennicutt \\& Garnett (1999) modelled extragalactic H\\two\\ regions using Kurucz and non-LTE models. They find that the temperatures of the ionizing stars decrease with increasing metallicity, and suggest that this can be explained by lowering the upper mass limit for star formation. In a further study, Kennicutt et al. (2000) have analysed a large sample of H\\two\\ regions containing stars of known spectral types and effective temperatures.  They confirm the dependence of stellar temperature on metallicity, although they note that this relationship depends on the correctness of the input ionizing fluxes, and particularly the amount of line blanketing. Oey et al. (2000) have presented a detailed comparison of spatially resolved H\\two\\ region spectra to photoionization models to test how well the CoStar models and Schmutz et al. (1992) W-R models reproduce the observed nebular line ratios. They find that overall the agreement is within 0.2\\,dex but the nebular models appear to be too hot by $\\sim 1000$\\,K, and suggest that the ionizing flux distribution may be too hard in the 41--54\\,eV range.  On the other hand, Kewley et al. (2001) model a large sample of infrared starburst galaxies, and find that the Schmutz et al. (1992) W-R atmospheres do not produce sufficient flux between 13.6--54\\,eV to match their observations.  It is clear that the use of expanding non-LTE model atmospheres in the analysis of single and unresolved H\\two\\ regions is a major improvement over using static LTE models to represent hot massive stars. The empirical studies described above show that, in addition, it is essential that the atmospheres are line-blanketed, particularly for studies of young stellar populations at solar or higher metallicities.  The CoStar models used in more recent studies do incorporate line blanketing but the lack of any line blanketing in the W-R atmosphere grid of Schmutz et al. (1992) is a serious deficiency. With the recent advances in computing and the development of large codes to calculate expanding, non-LTE line blanketed atmospheres, it is now feasible to compute a grid of realistic ionizing fluxes for OB and W-R stars.  In this paper, we present such a grid for five metallicities from 0.05--2\\,Z$_\\odot$ for use with the Starburst99 (Leitherer et al. 1999) evolutionary synthesis code and analyses of single star H\\two\\ regions.  In Section~\\ref{grid}, we present the details of the model atmosphere grid, and in Section~\\ref{comp}, we compare the predicted ionizing fluxes with previous single star models used in synthesis codes. In Section~\\ref{sb99}, we incorporate the new grid into Starburst99, and in Section~\\ref{evolcomp}, we examine the changes in the ionizing fluxes as a function of age for an instantaneous burst and continuous star formation. In Section~\\ref{photo}, we use the photoionization code {\\sc cloudy} (Ferland 2002) to assess the differences in the nebular diagnostic line ratios of H\\two\\ regions ionized by single O stars and synthetic clusters as a function of age. We discuss our results in Section~\\ref{discuss} and present the conclusions in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc} We have presented a large grid of non-LTE, line-blanketed models for O and W-R stars covering metallicities of 0.05, 0.2, 0.4, 1 and 2\\,Z$_\\odot$. The grid is designed to be used with the evolutionary synthesis code Starburst99 (Leitherer et al. 1999) and in the analysis of H\\two\\ regions ionized by single stars. We have computed 110 models for O and early B stars using the WM-basic code of Pauldrach et al. (2001).  The OB stellar parameters are defined at solar metallicity and are based on the most recent compilations of observational data.  The mass loss rates and terminal velocities have been scaled according to metallicity by adopting power law exponents of 0.8 and 0.13 (Leitherer et al. 1992).  For the W-R grid, we used the model atmosphere code {\\sc cmfgen} of Hillier \\& Miller (1998) and computed 60 WN and 60 WC models. The new grid is based on what we consider to be the most realistic parameters derived from recent observations and individual model atmosphere analyses. In particular, the upper temperature limit of the grid is 140\\,000\\,K since the vast majority of W-R stars that have been analysed fall below this value. For the W-R mass loss rates at solar metallicity, we used the relationships derived by Nugis \\& Lamers (2000) which are corrected for inhomogeneities in the W-R winds. For the first time, we have introduced a W-R wind/metallicity dependence and adopted the same power law exponents used for the O star models.  We argue in Sect.~\\ref{W-Rgrid} that recent theoretical and observational work indicate that the strengths of W-R winds must depend on metallicity. We stress that the new W-R grid excludes the few known examples of individual hot W-R stars with weak winds that have significant ionizing fluxes above 54\\,eV. Since they are so rare (in the Galaxy and LMC, at least), we expect them to make a negligible contribution to the ionizing flux of a young starburst.  We find significant differences in the emergent fluxes from the WM-basic models compared to the CoStar models of Schaerer \\& de Koter (1997). For supergiants, the wind density determines the transparency below 228\\,\\AA.  Generally, we find a lower flux in the He\\one\\ continuum with important implications for nebular line diagnostic ratios.  We believe that the CoStar models over-predict the number of He$^0$ ionizing photons through the neglect of photon absorption in line transitions and their re-emission at longer wavelengths (Crowther et al.  1999).  We compared the new W-R model emergent fluxes with the pure helium W-R models of Schmutz et al. (1992) which have very different parameters (particularly wind densities and temperatures). We find that at $\\sim 45\\,000$\\,K for solar metallicity WN models, the emergent flux below 504\\,\\AA\\ is much lower than the SLG92 models because of the inclusion of line blanketing. At 60\\,000\\,K, we find that blanketing is less important, and at 90\\,000\\,K, the wind density controls the emergent flux below 228\\,\\AA.  The WM-basic model grid has been integrated into Starburst99 (Leitherer et al. 1999) by directly replacing the Lejeune et al. (1997) LTE model library. For the W-R grid, we have used a new method of matching the models to the evolutionary tracks by taking a weighted mean of the uncorrected hydrostatic temperature $T_{\\rm {hyd}}$ and the corrected hydrostatic temperature $T_{2/3}$. This is necessary because of our lower, more realistic temperatures, particularly for the WC stars.  We next compared the output ionizing fluxes of the new grid integrated into Starburst99 with the evolutionary synthesis models of Leitherer et al. (1999) and Schaerer \\& Vacca (1998) for an instantaneous burst and continuous star formation. The changes in the ionizing outputs are dramatic, particularly during the W-R phase, with the details depending on metallicity because of line-blanketing and wind density effects. For an instantaneous burst, we find that the number of He$^+$ ionizing photons ($Q_2$) emitted during the W-R phase is negligible at $Z_\\odot$ and higher. At lower metallicities, $Q_2/Q_0$ is softer by a factor of $\\sim 20$ compared to the SLG92 models. In contrast to Schaerer (1996), we predict that nebular He\\two\\ will be at, or just below, the detection limit in low metallicity starbursts during the W-R phase.  We also find lower He$^0$ ionizing fluxes for $Z\\ge 0.4$\\, Z$_\\odot$ and ages of $\\le 7$\\,Myr compared to SV98 because of the smaller contributions of the line-blanketed models. The ionizing fluxes of the continuous star formation models emphasize the differences found for the single star models.  We have tested the correctness of the new model grid by computing nebular line diagnostic ratios using {\\sc cloudy} and single star and evolutionary synthesis models as inputs. For the former, we calculated the nebular He\\one\\ $\\lambda 5786$/H$\\beta$ ratio as a function of $Z$ since this is a sensitive diagnostic of $Q_1/Q_0$ for the temperature range where He is partially ionized. Observations indicate that this ratio decreases with increasing metal abundance, leading to the suggestion that the effective temperature decreases with increasing $Z$, and thus that the upper mass limit may be $Z$-dependent (e.g. Bresolin et al. 1999). The WM-basic models for dwarf O stars show a decrease in He\\one\\ $\\lambda 5786$/H$\\beta$ above $Z_\\odot$ in contrast to Kurucz LTE models which are essentially independent of $Z$. This decrease is in the same sense as the observations and suggests that the observed decline in He\\one\\ $\\lambda 5786$/H$\\beta$ with increasing $Z$ is simply due to the effect of line blanketing above $Z_\\odot$, and is not caused by a lowering of the upper mass limit.  The ionizing fluxes of the instantaneous burst models have been tested by comparing them to the predictions of SV98 and the observations of Bresolin et al. (1999) for extragalactic H\\two\\ regions.  In plots of the softness parameter $\\eta^\\prime$ against the abundance indicator $R_{23}$ for ages of 1--5\\,Myr and $\\log U=-2$ and $-3$, we find that the new models cover the same parameter space as the data points in contrast to the SV98 ionizing fluxes which are generally too hard at all ages.  We therefore conclude that the ionizing fluxes of the new model grid of O and W-R stars represent a considerable improvement over the model atmospheres that are currently available in the literature for massive stars. To prove their worth, they should be tested rigorously against specific observations of single H\\two\\ regions and young bursts of star formation.  The grid of ionizing fluxes for O and W-R stars and the updated Starburst99 code can be obtained from: {\\tt http://www.star.ucl.ac.uk/starburst}."
    },
    "0207/astro-ph0207624_arXiv.txt": {
        "abstract": "{\\small As described in the review by Marek Abramowicz (these proceedings), accretion theory is at something of a crossroads.  Many theoretical descriptions of observations have centered on the ``hydrodynamic'' approach of the ``standard'' Shakura \\& Sunyaev $\\alpha$-disk \\cite{shakura:73a}; however, recent MHD simulations of (non-radiative) accretion flow have called into question the validity of this approach \\cite{krolik:02a,hawley:01a}.  There has been a great deal of optimism that these simulations give direct insight into current observations as similar phenomena appear in both; e.g., jets \\cite{hawley:02b,hawley:01c}, rapid high-amplitude fluctuations \\cite{hawley:02a}, etc.  In comparison to real data from black hole candidates (BHC), however, these similarities are in many ways only superficial.  In some aspects, simple theories agree with the observations quite well.  The observed rapid variability is much more highly structured than found in simulations, and shows interesting correlations with spectra that have been interpreted in terms of simple models.  On the other hand, debates have arisen over even the most basic phenomenological issues concerning, e.g., the geometry and dominant radiation mechanisms of the accretion flow.  In this article I present my views of those observational properties of BHC states that most urgently need to be addressed, and I briefly discuss some of the models currently being debated.} ",
        "introduction": "BHC states can be broadly classified into two categories: hard with high variability (typically low luminosity), and soft with weak to moderate variability (typically high luminosity).  Historically, these have been classified as ``low'' and ``high'' states, respectively, with radio observations showing the former to be radio loud, and the latter to be (mostly) radio quiet (e.g., \\cite{fender:99b}). Examples of transitions between states are shown in Fig.~\\ref{fig:trans}.  The ``off'' state appears to be a low luminosity hard state \\cite{kong:00a}, and not necessarily a distinct state.  The ``very high state'' is a high luminosity soft state wherein the flux of the hard tail, the X-ray variability, and the radio emission are elevated from the very low levels observed in moderate luminosity soft states. Compared to the hard state, the very high state radio emission is more episodic and its spectrum is steeper, evocative of the ``ejection'' behavior seen in GRS~1915+105 (e.g., \\cite{pooley:97a,belloni:97b}). It is interesting to note that the very high state of \\gx, based solely on the X-ray observations, was suggested to be a jet-ejection state \\cite{miyamoto:91b}.  The ``intermediate state'' is somewhat ill-defined, and is similar to a very high state (i.e., increased hard tail and X-ray variability, with variable radio).  This label has been applied to spectra intermediate between ``high'' and ``very high'' states, as well as to spectra occurring inbetween ``hard/low'' and ``soft/high'' states.  The behavior exhibited by \\cyg\\ during its state transitions and ``failed'' transitions \\cite{pottschmidt:02a} may also qualify as an ``intermediate state''.  (The canonical soft state of \\cyg, however, is itself in many ways reminiscent of definitions of the ``intermediate state''.)  In this review, I will primarily refer to states as being ``hard'' or ``soft'', rather than use the labels of off/low/intermediate/high/very high.  \\begin{figure}[htb] \\centering \\psfig{file=mnowak01_fig1.eps,width=13cm} \\caption{{\\it Left:} State transitions in the BHC \\gx, showing a simultaneous rise in soft flux with a quenching of the radio \\protect\\cite{fender:99b}. RXTE observations of state transitions in \\lmc, showing a clear spectral softening as the flux rises \\protect\\cite{wilms:01a}.} \\label{fig:trans} \\end{figure}  X-ray spectra of low variability soft states are in fact remarkably well-described by simple disk models (with fitted inner disk radii comparable to the expected marginally stable orbit radii), and do not show any of the variability seen in MHD simulations.  Nor do they exhibit the instability behavior normally associated with $\\alpha$-disks \\cite{piran:78a}, despite the fact that their luminosities can be $\\approx 5$--$10\\%$ of their Eddington luminosity. Although transitions between extremes of the soft state and extremes of the hard state appear to track accretion rate, the luminosity at which these transitions occur is not fixed, and in fact exhibits evidence of hysteresis \\cite{miyamoto:95a}. That is, a state tends to remain soft to lower luminosity levels, or hard to higher luminosity levels, if it began in a soft or hard state, respectively.  This is observed in recent state transitions of \\gx\\ \\cite{nowak:02a}, as well as for numerous other BHC and neutron star X-ray binaries (\\cite{muno:02a,barret:02a}; Maccarone et al., in prep.).  Very complex behavior can be seen (e.g., the ``comb-like'' color-intensity diagrams of XTE J1550$-$564; \\cite{homan:01a}), and it is unlikely that accretion rate, even accounting for hysteresis, is the sole-determinant of state transition behavior.  As I further discuss below, even for a given luminosity \\emph{and} spectral hardness, it is not clear that any remaining spectral features (e.g., line strengths or widths) remain uniform from one instance of a given state to another. ",
        "conclusions": ""
    },
    "0207/astro-ph0207138_arXiv.txt": {
        "abstract": "{We report on Chandra ACIS-S observations of five type\\,I X-ray bursters with low persistent emission: \\SAXI, \\SAXII, \\SAXIII, \\SAXIV, and \\SAXV. We designate candidate persistent sources for four X-ray bursters. All candidates are detected at a persistent luminosity level of 10$^{32-33}$ $\\ergs$, comparable to soft X-ray transients in quiescence. From the number of bursters with low persistent emission detected so far with the Wide Field Cameras, we estimate a total of such sources in our Galaxy between 30 and 4000. ",
        "introduction": "Many low-mass X-ray binaries show bursts of X-rays which are characterized by a rapid rise and exponential decay, and by a black body spectrum with spectral softening during the decay i.e.\\ the emitter cools. Such type\\,I X-ray bursts are interpreted as thermonuclear flashes on surfaces of neutron stars, and thus effectively identify the emitting source as a neutron star as opposed to a black hole. The theory of these bursts predicts a relation between the accretion rate onto the neutron star, as measured by the persistent X-ray luminosity, and the properties of the X-ray burst. Briefly, for very low and very high accretion rates, no X-ray bursts are expected, because thermonuclear fusion is steady (Fujimoto et~al. 1987). At intermediate accretion rates, hydrogen/helium fusion occurs sporadically in bursts, and the burst frequency is a function of the accretion rate per square meter on the neutron star. Because the effectively accreting area of the neutron star is also a function of the accretion rate, the burst frequency is a non-monotonic function of the persistent X-ray luminosity. Recent reviews of burst theory are given by Bildsten (1998, 2000).  Low-mass X-ray binaries are discovered as either persistent sources or transient sources. The transient sources with neutron stars show outbursts lasting for weeks, sometimes up to years, at luminosities above $10^{36}$ $\\ergs$. During their quiescent state their luminosity drops to a level of $10^{32-33}$ $\\ergs$ (e.g. Campana et~al. 1998), and the time averaged luminosities are $\\ltap10^{36}$ $\\ergs$ (e.g. White et~al. 1984). Most bursts are emitted by systems at luminosities $\\gtap10^{36}$ $\\ergs$, e.g.\\ the transients Aql\\,X-1 and Cen\\,X-4 emitted X-ray bursts when they were in outburst (Koyama et~al. 1981, Matsuoka et~al. 1980).  The Wide Field Cameras (WFC) on board the Italian-Dutch Satellite BeppoSAX discovered sporadic type\\,I bursts from nine previously unknown burst sources, which had persistent X-ray fluxes below the WFC detection limit of a few times $10^{-10}$ $\\ergcms$ (2-28 keV). At 8 kpc, the distance of the Galactic center, these flux limits correspond to luminosities of $\\sim10^{36}$ $\\ergs$. Four of the nine previously unknown burst sources were detected with other instruments at fluxes well below the WFC detection limit (see Table\\,\\ref{previous}). The five other bursters are listed in Table\\,\\ref{overzicht}. In this article we present Chandra observations which we obtained in order to determine the flux levels of these five burst sources.  The persistent luminosities of the nine previously unknown burst sources are (possibly far) below $10^{36}$ $\\ergs$, i.e.\\ below the level X-ray bursts are usually observed. This is the reason why Cocchi et~al.\\ (2001) suggested that these sources are members of a new class of bursters with low persistent emission (see also Cornelisse et~al. 2002).  \\begin{table}[t] \\caption{Overview of the detection of four of the low persistent emission bursters. For each source we list the instrument which detected the source, the date of observation and the persistent flux in $10^{-11}$ $\\ergcms$ plus passband in keV. References: a. Kaptein et~al. 2000, b. Cocchi et~al. 1999, c. Pavlinsky et~al. 1994,  d. Cornelisse et~al. 2002, f. Antonelli et~al. 1999, g. in 't Zand et~al. 2002 (in preparation). \\label{previous}} \\begin{tabular}{l@{\\hspace{0.20cm}}c@{\\hspace{0.20cm}}c@{\\hspace{0.20cm}}c@{\\hspace{0.15cm}}c@{\\hspace{0.15cm}}c@{\\hspace{0.15cm}}} \\hline source & instrument & date & $F$ & range &ref\\\\ \\hline 1RXS\\,J1718.4$-$4029 & ROSAT/P & 1990 & 1    & 2-10 &a\\\\ 1RXS\\,J1718.4$-$4029 & ROSAT/H  & 1994 & 0.4  & 2-10 &a\\\\ \\GRS                 & GRANAT     & 1990 & 19   & 4-30 &b,c\\\\ SAX\\,J1828.5$-$1037  & ROSAT/P & 1993 & 0.19 & 0.5-2.5 &d\\\\ SAX\\,J2224.9$+$5421$^e$ & SAX/NFI & 1999 & 0.013 & 2-10&f,g\\\\ \\hline \\multicolumn{6}{l}{$^e$ Observation a few hours after burst.}\\\\ \\end{tabular} \\end{table}  The nine sources can be used to explore the low end of the relation between luminosity and  burst properties. The long waiting  times between type\\,I bursts, compared to brighter burst sources, plus the low persistent emission level make these sources difficult to discover. Its large field of view makes the WFC an efficient instrument for the detection of such rare events.  In Sect. 2 we describe the Chandra observations and data analysis and in Sect. 3 we discuss which of the detected sources are the most likely candidates for each burster. In Sect. 4 we briefly present unpublished but relevant observations with other instruments of \\SAXIV\\ and GRS\\,1741.9$-$2853. In Sect. 5 we discuss the implications for the class of low persistent emission bursters. ",
        "conclusions": "Three of the nine burst sources with low persistent emission discovered with the WFC were observed during ROSAT observations at luminosities of $\\simeq$$10^{34-35}$ $\\ergs$ a few years prior to the X-ray burst (Kaptein et~al. 2000; Cornelisse et~al. 2002; this paper). If the five burst sources of Table~\\,\\ref{overzicht} had similar luminosities and spectra, their countrate with Chandra would be several orders of magnitude higher than the countrates of the sources listed in Table\\,\\ref{result}.  Instead, the luminosities of the burst sources are at $\\simeq10^{33}$ $\\ergs$, comparable to the BeppoSAX/NFI observations of SAX\\,J2224.9$+$5421 (Antonelli et~al. 1999; in 't Zand 2002, in preparation) and in the range of quiescent soft X-ray transients with neutron stars (e.g. Campana et~al. 1998).  With the interstellar hydrogen column and upper limits to the distances listed in Table\\,\\ref{overzicht} and the spectrum described in Sect. 3, we compute the unabsorbed flux and upper limits to the luminosities between 0.5 and 7.0 keV. For the brightest candidate counterparts of \\SAXI, \\SAXII, \\SAXIV\\ and \\SAXV\\ we obtain upper limits to the unabsorbed persistent luminosity of 4$\\times$10$^{32}$, 3$\\times$10$^{32}$, 2$\\times$10$^{32}$ and 4$\\times$10$^{32}$ $\\ergs$, respectively. From the upper-limit derived from the observation of \\SAXIII\\ we get a luminosity of $<4\\times10^{32}$ $\\ergs$ (0.5-7 keV). These luminosities are indeed in the range expected for quiescent soft X-transients with a neutron star.  Because our Chandra observations give more than one possible counterpart, several possible counterparts must be chance coincidences. This is in agreement with known $\\log N-\\log S$ distributions, which predict $\\simeq$5 sources in the field of view (see e.g. Rosati et~al. 2002). This raises the question whether {\\em all} Chandra sources are chance coincidences, i.e.\\ whether we have not detected the actual counterparts for the bursters. Given that these systems are neutron star low mass X-ray binaries, we compare them to known other systems, i.e. the soft X-ray transients in quiescence. The lowest X-ray luminosities detected for quiescent soft X-ray transients are $\\sim10^{32}$ $\\ergs$ (e.g. Cen X-4, Campana et al.\\ 1998). This is around the detection limit for the Chandra observations discussed in this paper. We therefore consider it possible that we actually {\\em have} detected the persistent flux of the bursters, and that they are soft X-ray transients in quiescence, for which no outburst has as yet been detected. If so, this implies that their actual distances are not much less than the upper limits listed in Table\\,\\ref{overzicht}.  The persistent luminosities of the bursters observed with Chandra is well below the limit set with the WFC observations. This means that we cannot exclude that the persistent luminosity during the WFC observations was $\\sim$10-100 times higher than detected with Chandra, and that it was this higher flux level which triggered the burst.  The detections with ROSAT of 1RXS\\,J171824.2$-$402934 and SAX\\,J1828.5$-$1037, and of GRS\\,1741.9$-$2853 with GRANAT combined with non-detections at other epochs, show that the persistent flux level of these sources is variable.  The energy released during a burst due to nuclear fusion is about 1\\% of the accretion energy of the matter accreted onto the neutron star (see e.g. Lewin et al. 1993). Dividing the fluence of the bursts detected with the WFC by 1\\% of the persistent emission detected by Chandra we estimate burst intervals of $\\sim$10 years.  If only 1/6th of the persistent flux is due to accretion, the remainder being due to the cooling of the neutron star (i.e.  if only the power-law component is due to accretion, see Sect. 3) the estimated burst intervals rise to $\\sim$60 years. It is also suggested that the power-law component during quiescence is not due to accretion (see e.g. Campana et~al.  1998), and this means that the waiting time derived above is an under-limit. This explains why these events are so rare, and why we have only seen one burst for most of these sources.  This raises the question how many of these burst sources with low persistent emission exist in our Galaxy. With the WFC the Galactic Center region is observed every half year since 1996, for a total observation time of $5.5\\times10^6$ s up to end 2001. If we assume the Galactic distribution of low-mass X-ray binaries derived by van Paradijs \\& White (1995), $\\simeq$50\\% of the population is in the field of view (40$^\\circ$$\\times$40$^\\circ$) of the WFC. During all Galactic center observations 5 bursters with low persistent emission have been detected, i.e. \\SAXII, \\SAXIII, \\SAXIV, 1RXS\\,J171824.2$-$402934, and GRS\\,1741.9$-$2853 (the other four are outside the Galactic center region). This gives an average waiting time between the detection of these bursters of $1.1\\times10^6$ s. If we also assume that the waiting time between burst of one source is 60 years ($1.9\\times10^9$ s) we expect $\\simeq2\\times10^3$ sources in the Galactic center region, giving $4\\times10^3$ sources in the whole Galaxy. If on the other hand these sources are extensive periods of time at a persistent luminosity of $10^{34}$ $\\ergs$, as the detections of SAX\\,J1828.5$-$1037, 1RXS\\,J171824,2$-$402934, and \\GRS\\ suggests, the waiting time drops to 0.5 year (see Table\\,\\ref{previous}). This gives a number of 30 sources in our Galaxy. We conclude that the estimates for the total number of X-ray bursters with low persistent fluxes range from 0.5 to 60 times the number of known bursters ($\\simeq70$).  In this respect it is interesting to note that the first Chandra observations of globular clusters indicate that these systems harbour more quiescent soft X-ray transients than bursters with high persistent fluxes. For example, Liller\\,1, NGC\\,6440 and NGC\\,6652 all contain such quiescent sources in addition to the bright source (Homer et al.\\ 2001, Pooley et al.\\ 2002b, Heinke et al. 2001); and 47\\,Tuc, $\\omega$\\,Cen, NGC\\,6752 and NGC\\,6397 contain quiescent sources but no bright source (Grindlay et al.\\ 2001a, 2001b, Rutledge et al.\\ 2001c, Pooley et al.\\ 2002a). The formation mechanism for low-mass X-ray binaries in globular clusters (tidal capture or exchange encounter; see review by Hut et al.\\ 1992) is different from the formation mechanism in the galactic disk (evolution of a primordial binary). If the ratio of quiescent to bright X-ray bursters depends on the formation mechanism, we do not necessarily expect comparable ratios in the cluster and in the Galactic disk.  If bursts can arise from quiescent systems, we must consider the possibility that a burst from a globular cluster is due to a dim source, rather than to the bright source in it. This would undermine the argument that a burst from a cluster proves that the bright source in it is a neutron star. Nonetheless, we think that the argument holds in all eleven cases where it has been applied so far, as bursts from dim sources are extremely rare. For example, we have detected $\\sim2200$ X-ray bursts in our WFC observations of the Galactic center region; only five of these are from dim sources. Indeed, bursts from the globular cluster NGC\\,6440 were detected only when the transient in this cluster was active (in 't Zand et~al. 2001)."
    },
    "0207/astro-ph0207412_arXiv.txt": {
        "abstract": "A moderate resolution spectroscopic survey of Fleming's sample of 54 X-ray selected M dwarfs with photometric distances less than 25 pc is presented. All the objects consist of one or two dMe stars, some being doubles or spectroscopic binaries. Radial and rotation velocities have been measured by fits to the H$\\alpha$ profiles. Radial velocities have been measured by cross correlation. Artificial broadening of an observed spectrum has produced a relationship between H$\\alpha$ FWHM and rotation speed, which we use to infer rotation speeds for the entire sample by measurement of the H$\\alpha$ emission line.  We find 3 ultra-fast rotators (UFRs, $v \\sin i \\geq 100 $ km s$^{-1}$), and 8 stars with $30 $ km s$^{-1} \\leq v \\sin i < 100 $ km s$^{-1}$. We find that the UFRs have quite variable emission and should be observed for photometric variability. Cross-correlation velocities measured for ultra-fast rotators (UFRs) are shown to depend on rotation speed and the filtering used. The radial velocity dispersion of the sample is $17 $ km s$^{-1}$. A new double emission line spectroscopic binary with a period of 3.55 days has been discovered, RX~J1547.4+4507, and another known one is in the sample, the Hyades member RX~J0442.5+2027. Three other objects are suspected spectroscopic binaries, and at least six are visual doubles.  The only star in the sample observed to have significant lithium happens to be a known TW Hya Association member, TWA 8A. These results all show that there are a number of young ($< 10^{8}$ yr) and very young ($< 10^{7}$ yr) low mass stars in the immediate solar neighbourhood.  The H$\\alpha$ activity strength does not depend  on rotation speed. Our fast rotators are  less luminous than similarly fast rotators in the Pleiades. They are either younger than the Pleiades, or gained angular momentum in a different way. ",
        "introduction": "New technologies and space astronomy have led to the accumulation of catalogs of objects observed in widely different parts of the electromagnetic spectrum. Cross-correlating these catalogs has become a fruitful pursuit. \\citet{f98} published photometry of a sample of 54 M dwarf stars which were selected on the basis of detection in X-rays as part of the ROSAT All-Sky Survey and being red in photographic sky surveys, and which had apparent photometric parallaxes placing them closer than 25 parsecs from the Sun. Few of these stars had been studied spectroscopically, so we set out to observe their H$\\alpha$ and Li lines and to measure their radial velocities. We wanted to see whether their X-ray brightness indicated strong chromospheric activity: what fraction of Fleming's stars are dMe stars?  \\citet{f98} had shown that these stars probably have small proper motions, so we were interested in seeing whether other indications of youth are present. It has long been known that single M dwarfs decline in activity with age. The discovery of the TW Hydrae Association \\citep{kz97,wz99} shows that very young stars can be found in the immediate solar neighbourhood. Not only H$\\alpha$ emission, but also Li absorption and rotation as well as space motion are used as diagnostics of youth.  By taking spectra over some period of time, we hoped to be able to find spectroscopic binaries among these stars and to estimate their binary frequency. \\citet{fm92} have done this thoroughly for well-defined samples of relatively bright M dwarfs, including a few dMe stars. More recently,  the discovery of brown dwarfs has made searching for companions to low-mass stars a desirable goal, as reviewed by \\citet{b00}. ",
        "conclusions": "\\label{sec4}  We set out to do a spectroscopic investigation of Fleming's  sample of X-ray bright M dwarfs with apparent photometric distances of less than 25 pc \\citep{f98}. This was by far the faintest sample of stars  ever to be studied spectroscopically with the David Dunlap Observatory's 1.88 m reflector, an ambitious goal for an undergraduate practical astronomy class; many of the important papers in this field in recent years have involved 10-metre class telescopes at sites well away from and above the sort of environment we have near Toronto.  Coverage of the sample was uneven due to distribution around the sky and the weather. Nevertheless, our sensitivity, resolution and instrumental stability were sufficient to draw some interesting conclusions.  All of the X-ray sources studied by \\citet{f98} are dMe stars, though in one or possibly two cases the H$\\alpha$ emission comes from a fainter companion unresolved by Fleming. Emission equivalent widths are generally less than 10 \\AA. Observations by \\citet{hgr96} and more recent work by \\citet{mm01} suggest that dMe stars are more luminous than dM stars with the same color, and therefore larger. This means that Fleming's photometric distances are under-estimated, particularly since some of the stars may be quite young and therefore well above the main sequence.  At this stage, a full proper motion study of these stars remains to be completed, although \\citet{f98} has shown some preliminary results. The radial velocity dispersion of the sample, about $17 $ km s$^{-1}$, agrees with the single-component fits to the H$\\alpha$ emission sample of stars in the {\\it Third Catalogue of Nearby Stars} \\citep{gj91}, observed by \\citet{rhg95}. Further studies of motions should identify streams of different ages, as dramatically shown in the discovery of the TW Hya Association \\citep{kz97,wz99}.  Several of these stars show fast or ultra-fast rotation. Out of 54 $systems$ in Fleming's sample, 3 are UFRs with $v \\sin i \\gtrsim 100 $ km s$^{-1}$. About 10 more stars are fast rotators with $30 $ km s$^{-1} \\lesssim v \\sin i \\lesssim 60 $ km s$^{-1}$, including the only star in the sample with strong lithium absorption. While the star with lithium (RX~J1132.7-2651) is probably a member of the TW Hydrae association and therefore less than 10 Myr in age, the other fast and ultra-fast rotators are probably also quite youthful, especially when compared with the approximately 110 Myr age of the Pleiades and its rapid rotators \\citep{ts00}. \\citet{rm00} showed that $v \\sin i$ of M dwarfs in the Hyades was at most 35 km s$^{-1}$, with a maximum of around 100 km s$^{-1}$ in the Pleiades. If M dwarf rotation slows down with age as in solar-type stars, the rapid rotators in the field must be similar to the Pleiades in age, or even younger. We therefore have a significant population of young low-mass stars in the solar neighbourhood, generally with low proper motions, and as we see, quite low radial velocities relative to the Sun. This means that significant star formation has taken place within a few tens of parsecs of the Sun in the last few tens or hundred megayears \\citep{b00}.  An interesting question is whether the fast rotators have stronger H$\\alpha$ emission than the slower rotators, or whether saturation sets in at a relatively slow rotation speed \\citep{jj00}. \\citet{ts00}, in a careful study using much larger telescopes and spectrographs, showed that in the Pleiades there is a slight positive correlation between rotation speed and H$\\alpha$ equivalent width, but not for $v \\sin i > 20 $ km s$^{-1}$, which is our lower threshold. \\citet{rm00} show that $L_{H\\alpha}/L_{bol}$ has no correlation with rotation in either the Hyades or Pleiades, but there does appear to be a correlation of activity with age, with the Pleiades showing a factor of ten greater activity. This is extended to other clusters and ages in Figures 5.19 and 5.20 of \\citet{rh00}.  Figure~\\ref{fig12} shows the activity measure $L_{H\\alpha}/L_{bol}$ as a function of $v sin i$ for our sample, averaged for each star. There is at most a slight positive dependence of H$\\alpha$ flux on rotation speed, but given the wide range of rotation speeds in our sample, it is interesting that there is so little dependence, if any, on rotation. The ``Skumanich Law\" \\citep{sk72}, relating activity to rotation in solar-type dwarfs, does not seem to apply to mid to early M dwarfs.  Our results and those of \\citet{ts00} show that any saturation effect in H$\\alpha$ emission must take effect below $v \\sin i = 20 $ km s$^{-1}$. Combined with the X-ray results of \\citet{jj00}, we see that activity indicators in mid to early M dwarfs are strongly saturated at rotation rates above some speed below $v \\sin i = 20 $ km s$^{-1}$; \\citet{rh00} infer that the rotation threshold for activity is below 2 km s$^{-1}$. The results of this paper further confirm the idea that rotation and activity are not correlated in the mid and earlier M dwarfs.  \\begin{figure} \\figurenum{12} \\plotone{fig12.eps} \\figcaption[fig12.eps]{\\label{fig12} Activity strength $L_{H\\alpha}/L_{bol}$ versus rotation speed. Plain crosses represent apparently single stars, the star symbols  represent components of double-lined binaries, while filled squares represent members of visual doubles. The half-filled squares are suspected spectroscopic binaries.} \\end{figure}  The two SB2 dMe systems and the suspected binaries all have normal H$\\alpha$ emission strength. They have low rotation rates (the width of H$\\alpha$ in RX~J0102.4+4101 is probably not caused by rotation). While high resolution measurements are needed to accurately measure their rotation, it appears that the emission associated with tidally induced rotation in close dMe binaries is the same as in single dMe stars.  \\begin{figure} \\figurenum{13} \\plotone{fig13.eps} \\figcaption[fig13.eps]{\\label{fig13} Activity strength as a function of absolute bolometric magnitude. Symbols have same connotation as in Fig.~\\ref{fig12}. The {\\it mean relations} for the Pleiades, and the Hyades and field, are also shown. } \\end{figure}   Figure~\\ref{fig13} shows activity strength $L_{H\\alpha}/L_{bol}$ as a function of bolometric absolute magnitude M$_bol$. The mean relation for the Pleiades is estimated from Figure 9 of \\citet{rm00}, while the mean value of -3.9 for the Hyades and the field has been obtained by \\citet{rm00} and \\citet{hgr96}, respectively. In this diagram, our sample has the same distribution as the field and Hyades samples, and lies below the Pleiades distribution, where the less massive stars appear to be more active.  While the activity distribution as a function of mass looks just like the field and the Hyades, rotation as a function of mass is somewhat different, as seen in Figure~\\ref{fig14}. Here the lines represent the {\\it upper envelopes} of rotation in the Hyades and Pleiades. The UFRs of our X-ray selected sample are cooler (less massive) than the Pleiades, and rotate more rapidly than the Hyades. This suggests that the stars in our sample either acquired their angular momentum in a different manner than the M dwarfs in the Pleiades, or that our rapid rotators are significantly younger than the Pleiades and have not undergone as much braking. It is also possible that the sources of Pleiades rotation data as compiled by \\citet{rm00} are incomplete at the faint end. These diagrams convincingly show that for these stars the activity strength does not depend on rotation speed.  \\begin{figure} \\figurenum{14} \\plotone{fig14.eps} \\figcaption[fig14.eps]{\\label{fig14} Rotation speed as a function of absolute bolometric magnitude. Symbols as in Fig.~\\ref{fig12}. The {\\it upper envelopes} of the Pleiades and Hyades distributions are shown.} \\end{figure}  An interesting exercise will be to photometrically monitor the fast and ultra-fast rotators to obtain their rotation periods. We expect spots and photometric variability, as well as flares, given the H$\\alpha$ variability we have observed, and the presence of some of our stars in the catalog of flare stars \\citep{gk99}.  Are the true rotation speeds even higher than the observed $v \\sin i$ values? Trigonometric parallaxes combined with photometry will allow estimates of radii to be made. How much larger are these stars than ``main sequence\" M dwarfs? Are they still contracting to the main sequence (certainly the case for RX~J1132.7-2651A = TWA 8A), or do internal magnetic fields suppress convection in the core and thereby increase the radii of these stars \\citep{mm01}? What sort of magnetic braking takes place in these stars?"
    },
    "0207/gr-qc0207105_arXiv.txt": {
        "abstract": "An implicit fundamental assumption in relativistic perturbation theory is that there exists a parametric family of spacetimes that can be Taylor expanded around a background. The choice of the latter is crucial to obtain a manageable theory, so that it is sometime convenient to construct a perturbative formalism based on two (or more) parameters.  The study of perturbations of rotating stars is a good example: in this case one can treat the stationary axisymmetric star using a slow rotation approximation (expansion in the angular velocity $\\Omega$), so that the background is spherical. Generic perturbations of the rotating star (say parametrized by  $\\lambda$) are then built on top of the axisymmetric perturbations in $\\Omega$.  Clearly, any interesting physics requires non--linear perturbations, as at least terms $\\lambda\\Omega$ need to be considered. In this paper we analyse the gauge dependence of non--linear perturbations depending on two parameters, derive explicit higher order gauge transformation rules, and define gauge invariance. The formalism is completely general and can be used in different applications of general relativity or any other spacetime theory. ",
        "introduction": "An implicit fundamental assumption in relativistic perturbation theory is that there exists a parametric family of spacetimes such that the perturbative formalism is built as a Taylor expansion of this family around a background.  The perturbations are then defined as the derivative terms of this series, evaluated on this background~\\cite{wald}. In most cases of interest one deals with an expansion in a single parameter, which can either be a formal one, as in cosmology~\\cite{BMMS,MMB,BMT} or in the study of quasi--normal modes of stars and black holes~\\cite{chandrabook,kokkomodes}, or can have a specific physical meaning, as in the study of binary black hole mergers via the close limit approximation~\\cite{price,gleisclose}, or in the study of quasi--normal mode excitation by a physical source (see \\cite{kokkomodes,my123} and references therein). Typically the perturbative expansion stops at the first order, but recent interesting developments deal with second order perturbations~\\cite{BMMS,MMB,BMT,cl,gleiser}.  In some physical applications it may be instead convenient to construct a perturbative formalism based on two (or more) parameters, because the choice of background is crucial in having a manageable theory.  The study of perturbations of stationary axisymmetric rotating stars (see \\cite{kojimarev,laf,rsk} and references therein) is a good example. In this case, an analytic stationary axisymmetric solution is not known, at least for reasonably interesting equations of state. A common procedure is to treat axisymmetric stars using the so--called slow rotation approximation, so that the background is a star with spherical symmetry~\\cite{hartle,hartlethorne}. In this approach the first order in $\\Omega$ discloses frame dragging effects, with the star actually remaining spherical; $\\Omega^2$ terms carry the effects of rotation on the fluid. This is intuitive from a Newtonian point of view, as rotational kinetic energy goes like $\\Omega^{2}$. This approximation is valid for angular velocities $\\Omega$ much smaller than the mass shedding limit $\\Omega_K\\equiv\\sqrt{M/R_{star}^3}$, with typical values for neutron stars $\\Omega_K\\sim 10^3Hz$. Therefore the slow rotation approximation, despite the name, can still be valid for large angular velocities. In practice, the perturbative approach up to $\\Omega^2$ is accurate for most astrophysical situations, with the exception of newly born neutron stars (see~\\cite{STER} and references therein).  Given that the differential operators appearing in the perturbative treatment of a problem are those defined on the background, the theory is considerably simplified when the latter is spherical.  Generic time dependent perturbations of the rotating star (parametrized by a dummy parameter $\\lambda$ and describing oscillations) are then built on top of the stationary axisymmetric perturbations in $\\Omega$.  Clearly, in this approach any interesting physics requires non--linear perturbations, as at least terms of order $\\lambda\\Omega$ need to be considered.  A similar approach could be used to study perturbations of the slowly rotating collapse, even if in specific cases~\\cite{CPM,seidel,gundlach} the perturbative expansion depends by one parameter only.  Classical studies in the literature have not analysed in full the gauge dependence and gauge invariance of the non-linear perturbation theory.  For example, in~\\cite{CPM} the second order perturbations  are treated in a gauge invariant fashion on top of the first order perturbation in a given specific gauge. The perturbation variables used are therefore non gauge invariant under a complete second order gauge transformation~\\cite{BMMS,FW}, but only invariant under  ``first order transformations acting at second order''~\\cite{CPM}. While this may be perfectly satisfactory from the point of view of obtaining physical results, one may wish to convert results in a given gauge to a different one~\\cite{BMMS,MMB}, to compare results obtained in two different gauges, or to construct a fully gauge invariant formalism.  To this end one needs to know the gauge transformation rules and the rules for gauge invariance, either {\\it up to} order $n$~\\cite{BMMS} or {\\it at} order $n$ {\\it only}, as in~\\cite{CPM}. The situation is going to be more complicated in the case of two parameters, as we shall see.\\footnote{The concept of perturbation theory with more than one parameter has already been introduced, for first-order perturbations, in~\\cite{MAIA}, where the standard definition of spacetime perturbations~\\cite{stewart} is extended by using a (4+n)-dimensional flat space in which space-times are embedded.  The main aim of these works is to re-examine the gauge invariance of the metric. }  In this paper we keep in mind the above practical examples, but we do not make any specific assumption on the background spacetime and the two--parameter family it belongs to. As in \\cite{BMMS,SB,BS}, we do not even need to assume that the background is a solution of Einstein's field equations: the formalism is completely general  and can be applied to any spacetime theory.  We analyse the gauge dependence of perturbations in the case when they depend on two parameters, $\\lambda$ and $\\Omega$, derive explicit gauge transformation rules up to fourth order, i.e.\\ including any term $\\lambda^k\\Omega^{k'}$ with $k+k'\\leq 4$, and define gauge invariance. This choice of keeping fixed the total perturbative order is due to the generality of our approach.  In practical applications one would be guided by the physical characteristics of the problem in deciding where to truncate the perturbative expansion.  For example, in the case of a rotating star one could consider first order oscillations, parametrized by $\\lambda$, on top of a stationary axisymmetric background described up to $\\Omega^2$, neglecting therefore $\\lambda^2\\Omega$ terms.  Or instead, one could decide that $\\lambda^2\\Omega$ terms are more interesting than the $\\lambda\\Omega^2$ ones in certain cases.  From a practical point of view, our aim is to derive the effects of gauge transformations on tensor fields $T$ up to order $k+k'=4$. It is indeed reasonable to assume that in a practical example like that of rotating stars, at most one will want to consider second order oscillations~$\\sim\\lambda^2$ on top of a slowly rotating background described up to $O(\\Omega^2)$, in order to take into account large oscillations and fluid deformations due to rotation.  We will show that the coordinate form of a two--parameter gauge transformation can be represented by: \\beq \\fl {\\tilde x}^{\\mu} & = & x^{\\mu}+\\lambda\\xi^{\\mu}_{(1,0)}+ \\Omega\\xi^{\\mu}_{(0,1)} \\nn \\\\ \\fl & & +\\frac{\\lambda^2}{2}\\left(\\xi^{\\mu}_{(2,0)}+\\xi^{\\nu}_{(1,0)} \\xi^{\\mu}_{(1,0),\\nu}\\right) +\\frac{\\Omega^2}{2}\\left(\\xi^{\\mu}_{(0,2)}+\\xi^{\\nu}_{(0,1)} \\xi^{\\mu}_{(0,1),\\nu}\\right) \\nn \\\\ \\fl & & +\\lambda\\Omega\\left(\\xi^{\\mu}_{(1,1)}+\\epsilon_0\\xi^{\\nu}_{(1,0)} \\xi^{\\mu}_{(0,1),\\nu}+\\epsilon_1\\xi^{\\nu}_{(0,1)} \\xi^{\\mu}_{(1,0),\\nu}\\right)+O^3(\\lambda,\\Omega)\\,, \\label{coord2} \\eeq where the full expression is given in Eq.~(\\ref{coordtransf}).  Here $\\xi^\\mu_{(1,0)}$, $\\xi^\\mu_{(0,1)}$, $\\xi^\\mu_{(2,0)}$, $\\xi^\\mu_{(1,1)}$, and $\\xi^\\mu_{(0,2)}$ are independent vector fields and $(\\epsilon_0,\\epsilon_1)$ are any two real numbers satisfying $\\epsilon_0+\\epsilon_1=1$. Coupling terms like the $\\lambda\\Omega$ in (\\ref{coord2}) are the expected new features of the two--parameter case, cf.\\ \\cite{BMMS,SB}. Our main results are the explicit transformation rules for the perturbations of a tensor field $T$ and the conditions for the gauge invariance of these perturbations.  The paper is organized as follows: in Section~\\ref{taylorexp} we develop the necessary mathematical tools, deriving Taylor expansion formulae for two--parameter groups of diffeomorphisms and for general two--parameter families of diffeomorphisms.  In Section~\\ref{gauges} we set up an appropriate geometrical description of the gauge dependence of perturbations in the specific case of two--parameter families of spacetimes.  In Section~\\ref{gaugeiandt} we apply the tools developed in Section~\\ref{taylorexp} to the framework introduced in Section~\\ref{gauges}, in order to define gauge invariance and formulas for gauge transformations, up to fourth order in the two--parameter perturbative expansion. Section~\\ref{conclusions} is devoted to the conclusions. We follow the notation  used previously in \\cite{BMMS,SB,BS} for the case of one parameter perturbations. ",
        "conclusions": "Many astrophysical systems (in particular, oscillating relativistic rotating stars) can be well described by perturbation theory depending on two parameters. A well--founded description of two--parameter perturbations can  be very useful for such applications, specially in order to handle properly perturbations at second order and beyond. For example, one may wish to compare results derived in different gauges.  In this paper we have studied the problem of gauge dependence of non--linear perturbations depending on two parameters, considering perturbations of arbitrary order in a geometrical perspective, and generalizing the results of the one--parameter case~\\cite{BMMS,SB} to the case of two parameters. We  have constructed a geometrical framework in which a {\\it gauge choice} is a two--parameter {\\it group} of diffeomorphisms, while a {\\it gauge transformation} is a two--parameter {\\it family} of diffeomorphisms. We have shown that any \\mbox{two--parameter} family of diffeomorphisms can be expanded in terms of Lie derivatives with respect to vectors $\\xi^{\\mu}_{(p,q)}$. In terms of this expansion, which can be deduced order by order, we have derived general expressions for transformations of coordinates and tensor perturbations, and the conditions for gauge invariance of tensor perturbations.  We have computed these expressions up to fourth order in the perturbative expansion, i.e.\\ up to terms $\\lambda^k\\Omega^{k'}$ with $k+k'=4$.  The way in which the expansion of a two-parameter family of diffeomorphisms was derived in this paper is order by order, constructing derivative operators that can be rewritten as Lie derivatives with respect some vector fields.  The development of an underlying geometrical structure, analogous to the knight diffeomorphisms introduced in the one--parameter case~\\cite{BMMS}, would be interesting for two reasons: first, in order to have a deeper mathematical understanding of the theory, and second, in order to derive a close formula, valid at all orders, for gauge transformations and gauge invariance conditions. The present paper has been devoted to the derivation of the useful formulae for practical applications. In particular, our expressions will be useful to compare results derived in different gauges, and can form the basis for the construction of a gauge invariant theory of two-parameter systems in the line of works done for the one-parameter case like as for example~\\cite{moncrief,gerlach,bardeen,be,cb}. We leave the development of a more formal framework for future work.     \\appendix"
    },
    "0207/astro-ph0207142_arXiv.txt": {
        "abstract": "In this paper, we study heliospheric Ly-$\\alpha$ absorption toward nearby stars in different lines of sight.  We use the Baranov-Malama model of the solar wind interaction with a two-component (charged component and H atoms) interstellar medium. Interstellar atoms are described kinetically in the model.  The code allows us to separate the heliospheric absorption into two components, produced by H atoms originating in the hydrogen wall and heliosheath regions, respectively.  We study the sensitivity of the heliospheric absorption to the assumed interstellar proton and H atom number densities. These theoretical results are compared with interstellar absorption toward six nearby stars observed by the Hubble Space Telescope. ",
        "introduction": "The solar system is traveling in the surrounding Local Interstellar Cloud (LIC).  In the 1960s, it was realized [e.g., {\\it Patterson et al.}, 1963; {\\it Fahr}, 1968; {\\it Blum and Fahr}, 1970] that interstellar atoms penetrate deep into the heliosphere and, therefore, can be observed.  The interstellar atoms of hydrogen have been detected by measurements of the solar backscattered Ly-$\\alpha$ irradiance [{\\it Bertaux and Blamont}, 1971; {\\it Thomas and Krassa}, 1971].  Later, interstellar atoms of helium were also measured directly [{\\it Witte et al.}, 1996] and indirectly as backscattered solar irradiance [{\\it Weller and Meier}, 1981; {\\it Dalaudier et al.}, 1984]. At present, there is no doubt that inside the heliosphere, properties of He atoms such as temperature and velocity are different from those of H atoms. In particular, the H atoms are decelerated and heated compared with atoms of interstellar helium inside the heliosphere [{\\it Lallement et al.}, 1993; {\\it Costa et al.}, 1999; {\\it Lallement}, 1999]. The reason for this is a stronger coupling of H atoms with plasma protons in the heliospheric plasma interface through charge exchange.  Many works developed the concept of the heliospheric plasma interface over more than four decades after pioneering papers by {\\it Parker} [1961] and {\\it Baranov et al.} [1970].  \\begin{figure} \\noindent\\includegraphics[width=\\hsize]{fig1_new.eps} \\caption{ The heliospheric interface is the region of the solar wind interaction with LIC. The heliopause is a contact discontinuity, which separates the plasma wind from interstellar plasmas. The termination shock decelerates the supersonic solar wind. The bow shock may also exist in the interstellar medium. The heliospheric interface can be divided into four regions with significantly different plasma properties: 1) supersonic solar wind; 2) subsonic solar wind in the region between the heliopause and termination shock (i.e., the heliosheath) ; 3) disturbed interstellar plasma region (or \"pile-up\" region) around the heliopause; 4) undisturbed interstellar medium. \\label{fig1}} \\end{figure}  The heliospheric interface is formed by the interaction of the solar wind with the charged component of the interstellar medium (see Figure 1). The heliospheric interface is a complex structure having a multi-component nature.  The solar wind and interstellar plasmas, interplanetary and interstellar magnetic fields, interstellar atoms, Galactic and anomalous cosmic rays (GCRs and ACRs), and pickup ions all play prominent roles.  Interstellar hydrogen atoms interact with the heliospheric interface plasma by charge exchange.  This interaction significantly influences both the structure of the heliospheric plasma interface and the flow of the interstellar H atoms.  In the heliospheric interface, atoms newly created by charge exchange have the properties of local protons.  Since the plasma properties are different in the four regions of the heliospheric interface shown in Figure 1, the H atoms in the heliosphere can be separated into four populations, each having significantly different properties.  For example, population 3 consists of the atoms created by charge exchange with relatively hot protons in the region of disturbed interstellar plasma around the heliopause (region 3 in Figure 1).  It was realized theoretically by {\\it Baranov et al.} [1991] that the atoms of population 3 form a so-called ``hydrogen wall'' around the heliopause. The hydrogen wall is a significant enhancement of the density of interstellar atoms of hydrogen around the heliopause compared with the number density of the undisturbed local interstellar medium. The self-consistent models of the heliospheric interface predict that the secondary interstellar atoms (or atoms of population 3) are decelerated and heated compared with the original interstellar H atoms.  The Ly-$\\alpha$ transition of atomic H is the strongest absorption line in stellar spectra.  Thus, the heated and decelerated atomic hydrogen within the heliosphere produces a substantial amount of Ly-$\\alpha$ absorption. This absorption, which must be present in all stellar spectra since all lines of sight go through the heliosphere, has been unrecognized until recently, because it is undetectable in the case of distant objects characterized by extremely broad interstellar Ly-$\\alpha$ absorption lines that hide the heliospheric absorption.  We know now that in the case of nearby objects with small interstellar column densities, it can be detected in a number of directions. The heliospheric absorption will be very broad, thanks to the high temperature of the heliospheric H, and its centroid will also generally be shifted away from that of the interstellar absorption due to the deceleration, allowing its presence to be detected despite the fact that it remains blended with the interstellar absorption.  The hydrogen wall absorption was first detected by {\\it Linsky and Wood} [1996] in Ly-$\\alpha$ absorption spectra of the very nearby star $\\alpha$ Cen taken by the Goddard High Resolution Spectrograph (GHRS) instrument on board the Hubble Space Telescope (HST).  Since that time, it has been realized that the absorption can serve as a remote diagnostic of the heliospheric interface, and for stars in general, their ``astrospheric'' interfaces.  The $\\alpha$ Cen line of sight lies 52$^{\\circ}$ from the upwind direction of the interstellar flow through the heliosphere.  An additional detection of heliospheric H I absorption only 12$^{\\circ}$ from the upwind direction was provided by HST observations of 36 Oph [{\\it Wood et al.}, 2000a].  For downwind lines of sight, the heliospheric absorption is also not negligible [Williams et al., 1997]. Analysis of Ly-$\\alpha$ absorption toward Sirius shows that absorption of atoms of population 2 created in the heliosheath (see Figure~1) is needed in addition to interstellar and ``hydrogen wall'' absorption components in order to explain observations [{\\it Izmodenov et al.}, 1999].  In addition to heliospheric absorption, ``astrospheric'' absorption towards Sun-like stars was detected by {\\it Wood et al.} [1996], {\\it Dring et al.} [1997], and {\\it Wood and Linsky} [1998]. These observations might be used to infer properties of these stars and their interstellar environments [{\\it Wood et al.} 2001; {\\it M\\\"{u}ller et al.}, 2001].  A theoretical model of the heliospheric interface should be employed to interpret observations and put constraints on the local interstellar parameters and the heliospheric interface structure.  This model should self-consistently take into account plasma and H atom components.  Since the mean free path of H atoms is comparable with the size of the heliospheric interface, the H atom flow needs to be treated kinetically through the velocity distribution function, which is not Maxwellian [{\\it Baranov et al.}, 1998; {\\it M\\\"{u}ller et al.}, 2000; {\\it Izmodenov et al.}, 2001]. Recently, {\\it Wood et al.} [2000b] compared observations of Ly-$\\alpha$ toward six nearby stars with model-predicted absorption.  Both a Boltzmann mesh code [{\\it Lipatov et al.}, 1998] and a multi-fluid approach [{\\it Zank et al.}, 1996] were used to compute H atom distributions in the heliosphere.  In comparing these models with the data, it was found that the kinetic models predict too much absorption.  Models created assuming different values of the interstellar temperature and proton density fail to improve the agreement.  Surprisingly, it was found that a model that uses a multifluid treatment of the neutrals rather than the Boltzmann particle code is more consistent with the data [{\\it Wood et al.}, 2000b].  In this paper, we continue to study heliospheric absorption toward these stars.  In our study, we use the Baranov-Malama model [{\\it Baranov and Malama}, 1993, 1995, 1996] of the heliospheric interface, which is described in the next section.  This model uses a Monte Carlo code with splitting of trajectories, which allows very precise computation of H atom distributions.  Another advantage of the model is the possibility of separating of heliospheric H atoms into several populations, as discussed above. This model advantage allows us to consider separately two types of heliospheric absorption, hydrogen wall absorption and heliosheath absorption.  \\begin{figure} \\noindent\\includegraphics[width=\\hsize]{fig2_new.eps} \\caption{Density distributions of H atoms of population 3, atoms originated in the disturbed interstellar medium (region 3 in figure 1), and population 2, atoms originated in the heliosheath (region 2 in figure 2). Density distributions are shown as functions of heliospheric orientation relative to different line of sights. Curves 1, 2, 3, 4, 5, and 6 correspond to $\\theta = 12^{\\circ}$, $52^{\\circ}$, $73^{\\circ}$, $112^{\\circ}$, $139^{\\circ}$, and $148^{\\circ}$, respectively, where $\\theta$ is the angle relative to the upwind direction of the interstellar flow. The figure shows the number densities for model 3 (see Table 1). \\label{fig2}} \\end{figure}   \\begin{figure} \\noindent\\includegraphics[width=\\hsize]{fig3_new.eps} \\caption{ Number densities (upper row), velocities (middle row), and effective temperatures (lower row) of population 3 (left column) and population 2 (right column) toward 36 Oph. Dot-dash lines correspond to model 4, dot-dot-dot-dash lines correspond to model 5 and long dashes correspond to model 6 (see Table 1).} \\end{figure} ",
        "conclusions": "As seen from the previous section, absorption spectra toward five of six stars can be explained by taking into account the hydrogen wall absorption only.  Only the Sirius line of sight requires a detectable amount of heliosheath absorption to fit the data. Unfortunately, the hydrogen wall absorption is not very sensitive to such interstellar parameters as the interstellar proton and H atom number densities.  The hydrogen wall absorption is most apparent and detectable in upwind directions.  The small differences between the models and observations in these directions can be eliminated by small alterations of the stellar profiles (see Figure 7).  The poorest agreement is for model 2.  After including heliosheath absorption, model 4 seems to predict too much heliosheath absorption compared with the upwind data.  Therefore, model 2 and model 4 produce the worst agreement with the data.  This conclusion is also consistent with the analysis of absorption in crosswind directions (31 Com, $\\beta$ Cas).  Downwind lines of sight (Sirius and $\\epsilon$ Eri) can serve as good diagnostics of the heliosheath absorption.  For the assumed stellar line profiles, all models predict too much absorption toward the most downwind line of sight, $\\epsilon$ Eri.  However, since most of the discrepancy is away from the base of the absorption, we can try to correct this problem by modifying the shape of the assumed stellar Ly-$\\alpha$ profile.  Figure 8 demonstrates how the stellar profile of $\\epsilon$ Eri would have to be modified in order to improve agreement with the data.  We do not display the central part of the profile, because the absorption is saturated, providing no information on the stellar Ly-$\\alpha$ profile there.  Because the models predict too much absorption, the fluxes of the assumed stellar profile must be increased to improve the fit.  The revised profiles in Figure 8 are plausible in that they generally do not contain extremely steep slopes or fine structure. However, it is questionable if the tallest of these profiles (for models 4 and 5) are truly realistic. The revised stellar profiles in Figure 8 would suggest that the stellar Ly $\\alpha$ profile of $\\epsilon$ Eri is significantly taller and narrower than the line profiles of similar stars like $\\alpha$ Cen B and 36 Oph (see Figure 7), and the Sun.  There is some subjectivity in deciding what is plausible and what is not, but we would conclude from Figure 8 that models 4 and 5 probably require unrealistic modification to the stellar line profile to fit the data, but the other four models that require less extreme alterations of the assumed stellar profile may be all right.  In truth, however, the situation is actually worse than this, because as mentioned above our models do not extend far enough downwind and therefore underestimate the amount of heliosheath absorption.  Since a larger grid size would significantly worsen agreement with the data, it is uncertain whether any of the six models is truly consistent with the $\\epsilon$ Eri data.  Furthermore, since the kinetic models presented by {\\it Wood et al.} [2000b] had the same problem with the $\\epsilon$ Eri line of sight, it is possible that inaccuracies in the physical assumptions involved in current kinetic models may be resulting in overpredictions of absorption in downwind directions, although additional theoretical work is required to test this interpretation.  The models that have the least disagreement with the $\\epsilon$ Eri line of sight are models 1 and 6.  The common feature of these models is that both predict small number densities of H atoms of population 2.  In model 1 this is due to a low assumed interstellar H density, while in model 6 this is due to the small size of the heliosheath.  Model 6 in fact predicts the narrowest heliosheath region of our six models.  However, the interstellar proton and H atom number densities of this model seem to be unrealistically high, because the model predicts that the termination shock should be at 70 AU in the upwind direction.  As of August 2001, Voyager 1 was at 82 AU and Voyager 2 was at $\\sim$65 AU, and neither has yet encountered the termination shock.  Due to its low assumed interstellar H density, model 1 predicts the smallest amount of H atoms penetrating through the heliopause into the heliosphere of our six models. The number density of H atoms at the termination shock is 0.056 cm$^{-3}$ for this model.  This value is very close to the number density derived from the observed deceleration of the solar wind due to interaction with interstellar H atoms.  Comparison of the solar wind speed measured on Ulysses and Voyager shows a decrease of about 40 km~s$^{-1}$, or 10\\% in radial speed near 60 AU.  Wang and Richardson [2001] have shown that this speed decrease implies an interstellar neutral density at the termination shock of 0.05 cm$^{-3}$, which is close to our model 1 value.  Consideration of the Sirius line of sight complicates matters further. Model 2 and especially model 4 fit the observed profile very well, once again focusing our attention at the base of the absorption (see Figure 7), while the other models predict too little absorption. Models 2 and 4 assume a small interstellar ionization fraction in the vicinity of Sun, and recent studies of the LIC ionization state are in favor of such small ionization [{\\it Lallement}, 1999]. However, model 4 does not fit the other lines of sight well, and it must be noted once again that the models will underestimate the amount of heliosheath absorption in downwind directions due to the limited extent of the current model grid.  The Ly-$\\alpha$ absorption profile of the Sirius spectrum has been a controversial subject for the last few years and will probably remain so.  First interpreted as due to absorption by a stellar wind from Sirius A [{\\it Bertin et al.}, 1995a], the extra absorption on the blue side of the line has later been explained as resulting from an astrosphere around Sirius and shown to be compatible with a very crude model of such an astrosphere. However, the comparison between the Sirius B and Sirius A spectra later obtained by {\\it Hebrard et al.} [1999] has shown that a large fraction (if not the totality) of the extra absorption is specific to Sirius A and must arise closer to the star than predicted by an astrospheric model, favoring again a stellar wind as the source of the absorption.  On the other hand, the extra absorption on the red side, first interpreted as due to interstellar hot gas [{\\it Bertin et al.}, 1995b], has later been compared satisfactorily with heliospheric absorption by the heliosheath [{\\it Izmodenov et al.}, 1999], an interpretation which appears plausible based on Figure 7.  Unfortunately, the comparison with Sirius B does not help here to disentangle the sources of the absorption, since the Sirius B spectrum shows a broader and deeper absorption on the red wing, which encompasses the Sirius A line, and is identified by {\\it Hebrard et al.} [1999] as due to the intrinsic photospheric Ly-$\\alpha$ profile. {\\it Hebrard et al.} [1999] also note that the excess absorption on the red side of the line can be removed if one relaxes the assumption that both interstellar components detected toward Sirius have ${\\rm D/H}=1.6\\times 10^{-5}$.  The heliosheath or interstellar gas remain possible sources for the excess absorption seen towards Sirius A on the red side of the line.  If the absorption is due to the heliosheath, this is the only line of sight in which heliosheath absorption is detected.  We now suggest one possible cause of our difficulties with downwind lines of sight.  The process of pickup ion assimilation into the solar wind is a very complicated phenomenon.  Currently, it is believed that pickup ions form a separate population of protons, which is co-moving with the solar wind proton plasma. The process of energy exchange between the populations is quite slow and it is expected that the relaxation time is large compared with the time that is needed for a proton to reach the termination shock distance.  How the non-equilibrium distribution of pickup ions is affected by the termination shock is still an open question, although some scenarios have been developed [e.g., {\\it Fichtner}, 2001]. Our model considers the solar protons and pickup ion protons as one fluid.  This approach is based on fundamental questions of mass, momentum, and energy conservation. Therefore, it is quite natural to expect that the model predicts the positions of the shocks and the heliopause reasonably well. At the same time, measurements of speed and temperature of the solar wind at large heliocentric distances by Voyager are in favor of the two populations being co-moving, but thermally different. This suggests a different scenario of the heliosheath plasma flow than is assumed in our model.  This new scenario would result in different distributions of H atoms created in the heliosheath, which could possibly solve our difficulties in modeling the heliospheric absorption in downwind directions.  {\\it A priori}, we cannot exclude a scenario that provides the greatest absorption at heliocentric angles of about $130-150^{\\circ}$ from upwind.  For such a scenario, perhaps all six lines of sight in Figure 7 could be explained simultaneously. Calculating energetic neutral atom (ENAs) fluxes, {\\it Gruntman et al.} [2001] consider different plasma scenarios of heliosheath plasma flow.  One of the models considered in the paper suggests a maximum of ENA fluxes in directions of $130-150^{\\circ}$ from upwind.  This model assumes that the pickup proton population is carried through the shock without thermalization, i.e., preserving its velocity distribution filling the sphere in velocity space.  Since ENAs originate in the heliosheath, it supports the idea to consider different scenarios of the heliosheath plasma in order to interpret all six observable spectra simultaneously.  It remains to be seen whether such models could both reproduce the absorption seen toward Sirius and solve our problem of overpredicting absorption toward $\\epsilon$~Eri.  Finally, we have to note that an axisymmetric model of the heliospheric interface was used to calculate the heliospheric absorption. The actual heliosphere is not axisymmetric due to both the latitudinal variation of the solar wind and possible effects of the interstellar magnetic field.  Note, however, that the upwind direction of the ISM flow is only about $8^{\\circ}$ from the ecliptic, so variations of solar wind properties with latitude are unlikely to produce major differences from the axisymmetric approximation.  In any case, asymmetry in the heliosheath or hydrogen wall could result in the heliospheric absorption being different from axisymmetric.  A 3D model is needed to assess the inaccuracies introduced by the axisymmetric assumption, but unfortunately there is as yet no adequate 3D model including the interstellar H atom population that can address this issue. Such a model is currently being developed by us.  At the same time, our results can be interpreted as indirect evidence that there is no strong deviation of the heliospheric interface from being axisymmetric, since all four upwind/crosswind lines of sight discussed here can be well interpreted on the base of axisymmetric model. Taking into account the discussion above, one may conclude that our difficulties with the downwind lines of sight have other interpretations. However, this statement needs to be verified by 3D modeling of the heliospheric interface and by studying the sensitivity of heliospheric absorption to the 3D structure.  \\begin{figure} \\noindent\\includegraphics[width=\\hsize]{fig8_new.eps} \\caption{Modifications to the assumed stellar Ly-$\\alpha$ profile of $\\epsilon$ Eri that lead to better fits to the data for our 6 models. The lines are identified with the models the same as in Figure 7.} \\end{figure}"
    },
    "0207/astro-ph0207468_arXiv.txt": {
        "abstract": "We report an analysis of the dynamical structure of clusters of galaxies from a survey of photometric and spectroscopic observations in the fields of southern Abell Clusters. We analyze the galaxy velocity field in extended regions up to $ 7 h^{-1}$ Mpc from cluster centers and we estimate mean velocity dispersions and their radial dependence. Only one from a total number of 41 Abell clusters does not correspond to a dynamically bound system. However, four of these bound objects are double clusters.   We estimate that 20 \\% (7 clusters) of the 35 remaining are subject to serious projection effects. Normalizing the clustercentric distances by means of the overdensity radius $r_{200}$, and the velocity dispersion profiles (VDPs) by the corresponding mean cluster velocity dispersion, we computed the average VDP. Our results indicate a flat behavior of the mean VDP at large distances from the cluster center. Nevertheless, we found that for the inner part of the clusters ($r/r_{200}\\leq 1$) the VDP is up to a 10\\% smaller than at larger radii. ",
        "introduction": "Analysis of large scale structure formation may greatly benefit from studies of the dynamics of clusters of galaxies. Measurements of galaxy velocity dispersions in clusters provide reliable estimates of cluster masses and a direct normalization of the primordial mass power spectrum (Eke, Cole \\& Frenk, 1996). Moreover, the velocity field in the extended halos of clusters may set additional important constraints to the formation of structure as well on the mean density parameter of the universe.  There have been several recent studies on the dynamics of clusters of galaxies,  see for instance Girardi et al (1993), Zabludoff et al (1993), Collins et al (1995), Mazure et al (1996), Fadda et al. (1996) Alonso et al. (1999). The resulting distribution function of velocity dispersions from the ESO Nearby Abell Cluster Survey (ENACS) given by Mazure et al (1996) is in agreement with the distribution of cluster X-ray temperatures, suggesting $\\beta=\\sigma \\mu m_h/( k T_X) \\simeq 1$. The velocity dispersion profiles (hereafter VDP) may provide a useful tool for the study of the dynamics of clusters of galaxies. The analysis by Fadda et al. 1996 is consistent with a tendency of flat VDP in rich Abell clusters. Jing and B\\\"orner 1996 investigated the VDPs of clusters for several cosmological models. They found that on average, VDPs decrease with the cluster radius in every model up to $1h^{-1} Mpc$ from the cluster center. Also, these  authors found that the slope of the profiles are different in different models, being steeper in lower-$\\Omega$ models than in higher-$\\Omega$ models.   In the hierarchical scenario of structure formation, galaxy systems grow by aggregation of smaller structures formed earlier. Therefore, we expect a significant degree of substructure in clusters of galaxies if the remnants of the accretion of groups in the recent past has not been erased by the dynamical relaxation of the clusters. The substructure in rich clusters has been extensively analyzed in recent years (Dressler \\& Shectman (1988), West \\& Bothun (1990), Zabludoff et al. (1993), Girardi et al. (1997), Solanes et al. (1999)). The results are consistent with substructure in most of the cases studied, irrespective of the samples and method of analyses used. West \\& Bothun, (1990) made an analysis of substructure in clusters of galaxies and their surroundings. The authors developed a technique that is sensitive to correlations between galaxy positions and local kinematics finding little evidence for substructure in the inner regions, and significant departures from a relaxed substructure-free systems in the external regions.  More recently, Biviano et al. 2002 realized a detailed analysis of the consequences of substructure on luminosity and morphological segregation. These authors find that the number of galaxies in substructures decrease markedly toward the cluster center and report differences in the properties of galaxies depending whether they belong to substructures or not. These differences are also present in the the dynamical properties of galaxies.   Escalera et al. (1994) provide an extensive discussion of the presence of substructure in clusters of galaxies by using galaxy positions and redshifts. In their studies a multi-scale analysis is adopted that considers the kinematics as well as the wavelet transform providing estimators of the degree of substructure. Other works (see for instance Fadda et al. 1996) consider velocity gradients and anisotropy of galaxy orbits. Extensions of the different methods of analysis can provide new useful quantitative estimates of substructure, essential for a better understanding of the dynamics of clusters of galaxies.  The dynamics of clusters of galaxies may also be studied from information in the X-ray band. X-ray emission detected in a large fraction of clusters of galaxies provides an invaluable observational material. Several properties of the clusters and the intra-cluster medium may be addressed with this information. For example, the global mass distribution, the dynamical state and the evolution with redshift , the composition of the intra-cluster medium, etc. White (1999) presents an elegant methodology to recover the spatial properties of the intra-cluster gas from X-ray observations. From the deconvolution of ASCA satellite X-ray data, he finds a large fraction (90\\%) of clusters consistent with isothermality. These results are in conflict with the Markevitch et al. (1998) analysis from a sample of 30 clusters where most show steeply-declining intra-cluster temperature profiles. In their analysis of ASCA resolved spectroscopic data these authors obtained projected temperature profiles and in many cases two-dimensional temperature maps concluding that the gas temperature varies within a factor 1.3-2 or greater within the clusters.  The conflicting evidence for isothermality of the intra-cluster medium show that the information on the VDP for clusters may add important information to the subject. On the other hand, the tendency for subclustering to occur at large distances from cluster centers encourages us to explore the outer regions of clusters of galaxies. In this paper, we analyze the radial velocity distribution in regions extending up to $7 h^{-1}$Mpc  in projection from the Abell cluster center (we adopt $H_0$=100$hkms^{-1}$ and $q_0$ =0.5). We provide a detailed analysis of each individual cluster providing the degree of substructure and an estimate of the VDP at large distances from the cluster center. In section 2 we describe the method of analysis of substructure. Section 3 deals with the identification of the clusters and the projection effects which significantly affect the measurements of velocity dispersions. Section 5 provides the estimates of mean velocity dispersion and the correlation with richness counts as well as the velocity dispersion profile of several clusters. ",
        "conclusions": "A significant degree of substructure in clusters of galaxies is expected in the hierarchical scenario of structure formation. This is because of the large time-scale for the remnants of the accretion of groups onto clusters in the recent past to be erased by dynamical relaxation.  On the other hand, the identification of clusters in two dimensions may be strongly biased by spurious systems due to projection effects as shown in numerical simulations (van Haarlem, Frenk \\& White 1997). These two issues heavily complicate a detailed analysis of the dynamical properties of clusters of galaxies.   The evidence for substructure in rich clusters have been extensively explored in different studies of the galaxy distribution in cluster fields where the most accepted view is the relevant presence of substructure. The X-ray observations also contribute to our knowledge of the spatial properties of the intra-cluster gas. However, recent analyses  provides conflicting results regarding the isothermality of the gas or the existence of steeply-declining temperature profiles. This issue, and the fact that the degree of substructure increases at large distances from cluster centers motivated the present study of the outer regions of clusters of galaxies.  Our work is mainly centered in the analysis of the radial velocity distribution of galaxies in extended regions of Abell clusters, focusing on the existence of gradients in the velocity dispersion profiles.  We have carried out an analysis of the velocity field of galaxies in extended regions up to $7 h^{-1}$ Mpc from the cluster centers. We have applied several methods to remove contamination by projection effects and analyzed the presence of sub-clustering. We have obtained suitable estimates of the mean velocity dispersions and its radial dependence using the ROSTAT routines. Our analysis can be compared to  Fadda et al. 1996 for a fraction of common objects. It is clear from our analysis that the larger differences arise in those clusters with more contamination and a smaller number of measured redshifts. We also find that the correlation between mean velocity dispersion $\\sigma$ and richness number counts $\\cal N$ is strongly affected by projection effects. There is some evidence of  correlation between $\\sigma$ and $\\cal N$ for a sub-sample restricted to systems with no significant degree of contamination.  From our original sample of 41 Abell clusters we found that 40 are real clusters although 4 of these appear as double systems. These results are similar to those found by Mazure et al. 1996. These authors found that almost all ACO clusters with richness class 1 or greater correspond to real systems in the redshift space and about 10\\% of the ACO clusters appears to be the result of a superposition of two similar poorer systems.  Beside  the double systems we also found that  7 of the clusters in our sample are subject to serious projection effects. The  fraction of clusters with a high degree of contamination in our sample compares well with the results of such effects in the mock catalogs from the numerical simulations of van Haarlem, Frenk \\& White 1997) where 1/3 of Abell-type clusters are expected to arise from projection of groups along the line of sight.  From the resulting sample of 44 clusters, four are poorly defined and more data are needed in order to better establish their properties. In spite of projection effects, 28 of the 44 clusters are well defined in redshift space and have a velocity distribution consistent with a gaussian function.  Our results show that the average VDP is flat at large distances from the cluster center. This behavior is found for 19 clusters (65\\% of the 29 with VDP estimates) indicating that  an isothermal hypothesis can be assumed even at radii well beyond the virial radius. Nevertheless, we found that on average, the normalized velocity dispersion is about 10\\% smaller in the inner region of the clusters ($r/r_{200} \\leq 1$). A possible interpretation for the decay of VDPs in central regions can be related to the morphological segregation in clusters. Fadda et al. 1996 found kinematical segregation in the sence that early type galaxies show smaller values of $\\sigma$ than late types and Mazure et al. 1996 found that the brightest cluster galaxies (typically of early type morphology) move slower than other galaxies and Ram\\'{\\i}rez et al. 1998  found that the differences between the velocity distributions of elliptical and spiral galaxies are associated with the shape of their orbit families. Since early type objects are dominat within $r_{200}$ the results shown in figure 7 are to be expected.  \\placefigure{fig-7} \\def\\figureps[#1,#2]#3.{\\bgroup\\vbox{\\epsfxsize=#2 \\hbox to \\hsize{\\hfil\\epsfbox{#1}\\hfil}}\\vskip12pt \\small\\noindent Figure#3. \\def\\par{\\endgraf\\egroup\\vskip12pt}} \\figureps[fig7.ps,1.00\\hsize] 7. Mean VDP for the total sample of clusters with VDP estimate. Each cluster has been normalized using the corresponding mean $\\sigma$.  The shape of VDP profiles are of fundamental importance for their implications on cluster properties and cosmology since detailed theoretical predictions from different cosmological scenarios could be used to set restrictions to current models of structure formation. Jing and B\\\"orner 1996 analysis of VDPs of clusters for several cosmological models show an average decline of VDPs with the distance to cluster centers. Nevertheless, this analysis was restricted to the very inner  region of the clusters ($\\leq 1h^{-1} Mpc$. Therefore, new numerical simulations must be analyzed to test our findings of flat VDPs at very large distances from the cluster center.  Under the assumption that clusters are in global dynamical equilibrium (even beyond the virialization radius) our results can be compared to temperature radial distributions derived from the X-ray emission of the intra-cluster medium.  Flat VDP at large clustercentric distances may shed light on the recent  controversy  on the nature, either flat or  declining, of intracluster temperature radial profiles (see White 2000).   \\noindent Acknowledgments  We thank the Referee for useful comments and suggestions which greatly improved the original version of this paper. This work was partially supported by the Consejo de Investigaciones Cient\\'{\\i}ficas y T\\'ecnicas de la Rep\\'ublica Argentina, CONICET; SeCyT, UNC and Agencia C\\'ordoba Ciencia, Argentina. L. Infante and H. Quintana acknowledge FONDECYT and  FONDAP, Chile for partial support. H. Quintana acknowledges the support of a Catedra Presidencial en Ciencia, Chile."
    },
    "0207/astro-ph0207232_arXiv.txt": {
        "abstract": "We present the follow-up of three medium redshift galaxy clusters from the SHARC survey observed with {\\sc Xmm--newton}. We studied RX~J0256.5+0006 which shows two components which are very likely in interaction. The smallest component exhibits a comet-like structure indicating ram pressure stripping as it falls onto the main cluster. The second cluster, RX~J2237.0-1516 is an elliptical cluster with a gas temperature of 3.0$\\pm$0.5\\,keV. The third cluster, RX~J1200.8-0328 seems to be in a relaxed state because its shape is regular and we do not see obvious temperature gradient. Its mean temperature is 5.1$^{+0.7}_{-0.5}$\\,keV. ",
        "introduction": "Catalogs of galaxy clusters with a large range of redshifts and cluster {\\sc x}-ray luminosities are an ideal basis for the test of cosmological parameters (e.g. Henry 2000~; Borgani et al. 2001) and the study of cluster formation and evolution. The Bright and Southern Serendipitous High Redshift Archival {\\sc Rosat} Cluster (SHARC) surveys provide a sample of clusters detected in {\\sc Rosat} observations over two decades of {\\sc x}-ray luminosities (10$^{43}\\ <$ L$_{\\mathrm{x}}\\ <\\ 10^{45}$\\,erg/s) with redshifts between 0.2 and 0.8 (Romer et al. 2000~; Collins et al. 1997~; Burke et al. 1997).  We present here preliminary results based on the follow-up observations of three of these selected galaxy clusters with {\\sc Xmm--newton}~: RX~J0256.5+0006, RX~J2237.0-1516 and RX~J1200.8-0328. For the f\\mbox{}irst time, the new generation of {\\sc x}-ray observatory, like {\\sc Xmm-newton} and {\\em Chandra} give capabilities to do precise spectroscopic and imaging analysis at the same time (see e.g Arnaud et al. 2002a).  Our analysis is based on the three {\\sc epic} cameras, {\\sc mos}1\\&2 and pn. Throughout the paper, we use a cosmology with  H$_{0}=50$\\,km/s/Mpc, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. The error bars are given with a conf\\mbox{}idence level of 90\\%. ",
        "conclusions": ""
    },
    "0207/astro-ph0207004_arXiv.txt": {
        "abstract": "{We report on XMM-Newton observations of the galactic supersoft X-ray source \\rx . The RGS spectrum exhibits a wealth of spectral features from iron and oxygen. XMM-Newton data confirm the finding of previous Chandra HETGS/MEG observations that NLTE models of hot white dwarf atmospheres fail to represent the complex spectrum. There are clear evidences for P Cygni profiles with wind velocities of up to 2000 km\\,s$^{-1}$. Small flux variations with time scales larger than 1000\\,s are present. The strongest power is at $\\sim$ 0.21\\,d, a period close to that seen in V band optical light curves. A detailed analysis of the associated changes in the RGS and EPIC pn spectra hint at a mostly grey mechanism suggesting a variation of the visibility of the white dwarf due to occulting material in the accretion disk. Finally, we detect radial velocity changes of 173 $\\pm$ 47 km\\,s$^{-1}$ between two RGS observations obtained half an orbital cycle apart. The amplitude of the RGS velocity shift is consistent with that of the optical \\ion{He}{ii} $\\lambda$ 4686 and thus supports the idea that most of the \\ion{He}{ii} optical line emission arises from the accretion disk.} ",
        "introduction": "Supersoft X-ray sources have been discovered by the Einstein satellite and extensively studied by ROSAT. They usually consist of a close binary system in which the accreting component is a very hot white dwarf responsible for the soft thermal-like radiation (van den Heuvel et al. \\cite{vdh}). At the low spectral resolution of the ROSAT PSPC, the energy distributions look like black bodies with temperatures in the range of 20 to 70 eV. Higher resolution spectra with ASCA (e.g. Ebisawa et al. \\cite{ebisawa96}) started to reveal complex structures later emphasized by Chandra and XMM-Newton grating spectroscopy (Paerels et al. \\cite{pae01}, Bearda et al. \\cite{bea02}). Fitting model atmospheres of hot white dwarfs usually imply bolometric luminosities below but close to Eddington. Stable or quasi stable hydrogen burning at the surface of the white dwarf can adequately explain most of the observational features of supersoft X-ray sources (van den Heuvel et al. \\cite{vdh}). The high accretion rates required by stable thermonuclear burning ($\\dot{\\rm M}$ $\\sim$ 10$^{-7}$ \\Msol yr$^{-1}$, Iben  \\cite{iben82}) can be obtained through thermally unstable Roche lobe overflow when the mass donor star is more massive than the white dwarf. However, this picture  is probably not valid for many sources since short periods systems cannot accommodate a massive mass-donor star. In these cases, strong irradiation by the very luminous central source could play a major role in driving the mass transfer mechanism (Van Teeseling \\& King \\cite{tee98}).  Because stable burning allows matter to pile up at the surface of the white dwarf, contrary to novae where most matter is expelled in the runaway event, supersoft sources are considered as possible SN Ia progenitors. A sizeable fraction of SN Ia could occur in these systems when the white dwarf approaches the Chandrasekhar limit (Rappaport et al. \\cite{rap94}).  We report in this paper on a long XMM-Newton observation of \\rx \\ involving the EPIC pn, RGS and OM instruments. \\rx \\ is one of the few galactic supersoft source known (Motch et al. \\cite{mhp94}). It is identified with a V $\\sim$ 17.2 reddened object exhibiting Balmer and high excitation lines of \\ion{He}{ii}, \\ion{O}{vi} and \\ion{C}{iv}. The relatively high interstellar absorption towards the source (\\nh \\ $\\sim$ 10$^{22}$ cm$^{-2}$) only allows observation in the highest energy range (above $\\sim$ 0.4 keV). Among supersoft sources, \\rx \\ is unusual by its relatively high X-ray temperature and long orbital period of 4.0287 d (Schmidtke \\&  Cowley \\cite{schmi01}). Previous X-ray observations have been reported from a number of satellites (ASCA; Ebisawa et al. \\cite{ebisawa01}, BeppoSAX; Hartmann et al. \\cite{hh99} and recently using the HETG on board Chandra; Bearda et al. \\cite{bea02}). ",
        "conclusions": "XMM-Newton RGS spectra have revealed a wealth of spectral features, some of them being apparently resolved by the RGS. Emission and also maybe absorption lines due to various iron ions are seen as well as Lyman $\\alpha$ \\ion{O}{viii}. The RGS spectra confirm all the features seen in the Chandra HETGS by Bearda et al. (\\cite{bea02}) and illustrate further the difficulty to correctly model the emission of hot white dwarf atmospheres. Clearly, more sophisticated models including wind effects are required.  The X-ray spectrum of \\rx \\ display little changes with both X-ray flux level and orbital phase. At the start of the first observation, the X-ray source is receding with maximum positive \\ion{He}{ii} velocity while two days later the source moves towards us with maximum negative \\ion{He}{ii} velocities. This means that the geometrical configuration of the system is not very different between the two observations. In the absence of wake material, the column density due to the wind of the mass donor star does not change drastically and this is at first order consistent with the lack of large phase dependence of the RGS spectrum.  Very few lines show evidences for possible intensity related variations and the EPIC pn clearly indicates that absorption by cold material is not the mechanism responsible for the X-ray modulations. Interestingly, the 0.21\\,d period for X-ray oscillations is very close to the 0.23\\,d preferred time scale reported by Schmidtke et al. (\\cite{schmi00}) in the optical light curve. The full amplitude of sine fit to the soft X-ray flux modulation at 0.21\\,d is $\\sim$ 11\\% while that of the optical V band modulation is between two and three times less ($\\sim$ 4\\% Schmidtke et al. \\cite{schmi00}). Since we do not have simultaneous optical V band and X-ray observations we do not know the exact flux ratio nor phasing of the modulations. OM data are not of high enough S/N to reveal a modulation at the level of a few percents. Observations at the Dutch telescope were too short to constrain the amplitude of a contemporaneous 0.2\\,d optical periodicity. However, such an optical to X-ray amplitude ratio is expected when optical emission is dominated by X-ray heating of the companion star or of the accretion disk (see e.g. Van Paradijs \\& Mc. Clintock \\cite{vm94}, Matsuoka et al. \\cite{ma84}). A change in the visible emitting area of the accreting white dwarf by the internal rim of the accretion disc could thus explain both the X-ray and optical 0.2\\,d modulation. Simultaneous X-ray and optical photometric observations may reveal beating effects at the 4\\,d orbital period and could thus constrain the amount of optical light resulting from X-ray heating of the mass donor star and/or of the accretion disk bulge.  The RGS spectra also confirm the existence of a strong wind clearly affecting the Lyman $\\alpha$ \\ion{O}{viii} emission line and probably many other lines. The wind velocity of $\\sim$ 2000 km\\,s$^{-1}$ is more than two times slower than the velocity of the transient jet detected by Motch (\\cite{m98}) at 5200 km\\,s$^{-1}$ or of that of the candidate receding $\\lambda$ 6680.3 \\AA \\ component proposed by Schmidtke et al. (\\cite{schmi00}) at 5350 km\\,s$^{-1}$. In the absence of realistic model atmosphere including wind effects, it is difficult to strictly rule out the presence of weak satellite lines at $\\pm$ 5200 km\\,s$^{-1}$ in the RGS spectrum. The autocorrelation of the RGS spectra do show some signal at similar velocity shifts. However, a more likely explanation is the presence of several iron lines with the adequate spacing in wavelength.  Amazingly, the amplitude and direction of the RGS radial velocity changes with orbital phase (173 $\\pm$ 47 km\\,s$^{-1}$) is fully compatible with that expected from the \\ion{He}{ii} $\\lambda$ 4686 optical line. This observation thus supports the idea that in this system at least, the \\ion{He}{ii} $\\lambda$ 4686 velocity represents the orbital motion of the X-ray source.  This may be one of the first observation measuring orbital Doppler shifts in high energy spectral features from an X-ray source. Our coverage in orbital phase is not dense enough to derive the amplitude and phasing of the X-ray lines with respect to the optical \\ion{He}{ii} line. As for the H$\\alpha$ line, some of the \\ion{He}{ii} emission may arise from the heated hemisphere of the companion star or from the bulge of the accretion disc. This additional component may reduce the amplitude of the \\ion{He}{ii} velocity change with respect to K$_{\\rm X}$ and add some phase lag. A long RGS observation of \\rx \\ may thus be able to measure with relatively high accuracy the true orbital velocity of the accreting object. A reliable K$_{\\rm X}$ estimate could help constraining the mass ratio and check whether the high mass transfer in \\rx \\ can be driven by thermally unstable Roche lobe overflow."
    },
    "0207/hep-th0207142_arXiv.txt": {
        "abstract": "\\baselineskip=12pt  This paper comprises the written version of the lectures on string theory delivered at the 31st British Universities Summer School on Theoretical Elementary Particle Physics which was held in Manchester, England, August 28 -- September~12~2001. ",
        "introduction": "}  These notes comprise an expanded version of the string theory lectures given by the author at the 31st British Universities Summer School on Theoretical Elementary Particle Physics (BUSSTEPP) which was held in Manchester, England in 2001. The school is attended mostly by Ph.D. students in theoretical high-energy physics who have just completed their first year of graduate studies at a British university. The lectures were thereby appropriately geared for this level. No prior knowledge of string theory was assumed, but a good background in quantum field theory, introductory level particle physics and group theory was. An acquaintance with the basic ideas of general relativity is helpful but not absolutely essential. Some familiarity with supersymmetry was also assumed because the supersymmetry lectures preceeded the string theory lectures at the school, although the full-blown machinery and techniques of supersymmetry were not exploited to any large extent.  The main references for string theory used during the course were the standard books on the subject~\\cite{GSW,PolBook} and the more recent review article~\\cite{JohnsonRev}. The prerequisite supersymmetry lectures can be found in~\\cite{BUSSTEPPJMF}.\\footnote{\\baselineskip=12pt Further references are cited in the text, but are mostly included for historical reasons and are by no means exhaustive. Complete sets of references may be found in the various cited books and review articles.} The lectures were delivered in the morning and exercises were assigned for the tutorial sessions which took place in the afternoons. These problems are also included in these notes. Many of them are intended to fill in the technical gaps which due to time constraints were not covered in the lectures. Others are intended to give the student a better grasp of some ``stringy'' topics. The present paper has expanded on many aspects of string theory that were addressed during the school, mainly to make the presentation clearer.  There were six one-hour lectures in total. Since string theory is nowadays such a vast and extensive subject, some focus in the subject material was of course required. The lectures differ perhaps from most introductory approaches since the intent was to provide the student not only with the rudiments of perturbative string theory, but also with an introduction to the more recently discovered non-perturbative degrees of freedom known as ``D-branes'', which in the past few years have revolutionalized the thinking about string theory and have brought the subject to the forefront of modern theoretical particle physics once again. This means that much of the standard introductory presentation was streamlined in order to allow for an introduction to these more current developments. The hope was that the student will have been provided with enough background to feel comfortable in starting to read current research articles, in addition to being exposed to some of the standard computational techniques in the field. The basic perturbative material was covered in roughly the first three lectures and comprises sections~\\ref{History}--\\ref{Superstrings}. Lecture~4 (section~\\ref{RRCharge}) then started to rapidly move towards explaining what D-branes are, and at the same time introducing some more novel stringy physics. Lectures~5 and 6 (sections~\\ref{DBraneGauge} and \\ref{DBraneDyn}) then dealt with D-branes in detail, studied their dynamics, and provided a brief account of the gauge theory/string theory correspondence which has been such an active area of research over the past few years. ",
        "conclusions": ""
    },
    "0207/astro-ph0207096_arXiv.txt": {
        "abstract": "We use FRIIb radio galaxy redshift-angular size data to constrain cosmological parameters in a dark energy scalar field model. The derived constraints are consistent with but weaker than those determined using Type Ia supernova redshift-magnitude data. ",
        "introduction": "The last half-a-dozen years have seen a remarkable increase in the quality of some cosmological data. No less remarkable, but perhaps less heralded, has been the continuing acquisition of new types of data. These have been very useful developments in the on-going process of determining, through the cosmological tests, how well current cosmological models approximate reality: many independent and tight constraints on cosmological-model parameters allow for consistency checks on the models (see, e.g., Maor et al. 2002; Wasserman 2002).  For example, there is now much more higher-quality Type Ia supernova redshift-magnitude data. Recent applications of the redshift-magnitude test based on this data (see, e.g., Riess et al. 1998; Perlmutter et al. 1999; Podariu \\& Ratra 2000; Waga \\& Frieman 2000; Leibundgut 2001) indicate that the energy density of the current universe is dominated by a cosmological constant, $\\Lambda$, or by a term in the stress-energy tensor that only varies slowly with time and space and so behaves like $\\Lambda$.\\footnote{ See, e.g., Peebles (1984), Ratra \\& Peebles (1988), Efstathiou, Sutherland, \\& Maddox (1990), Ratra et al. (1997, 1999), Steinhardt (1999), Sahni \\& Starobinsky (2000), Brax, Martin, \\& Riazuelo (2000), and Carroll (2001) for discussions of such models.} Supporting evidence for $\\Lambda$ or a $\\Lambda$-like term is provided by a combination of low dynamical estimates for the non-relativistic matter density parameter $\\Omega_0$ (see, e.g., Peebles 1993) and evidence for a vanishing curvature of spatial hypersurfaces from cosmic microwave background anisotropy measurements (see, e.g., Podariu et al. 2001; Baccigalupi et al. 2002; Scott et al. 2002; Mason et al. 2002).  Evidence against the large value of the cosmological constant density parameter $\\omegal$ favored by the above tests comes from estimates of the observed rate of multiple images of radio sources or quasars, produced by gravitational lensing by foreground galaxies (see, e.g., Ratra \\& Quillen 1992; Helbig et al. 1999; Waga \\& Frieman 2000; Ng \\& Wiltshire 2001).  An improvement in data quality, as well as data from other cosmological tests, will be needed to resolve this situation. In the near future the redshift-counts test appears to be promising (see, e.g., Newman \\& Davis 2000; Huterer \\& Turner 2001; Podariu \\& Ratra 2001; Levine, Schulz, \\& White 2002).  Present redshift-angular size data provide a useful consistency check. The redshift-angular size relation is measured by Buchalter et al. (1998) for quasars, by Gurvits, Kellermann, \\& Frey (1999) for compact radio sources, and by  Daly \\& Guerra (2002) for FRIIb radio galaxies. Vishwakarma (2001), Lima \\& Alcaniz (2002), and Chen \\& Ratra (2003) use the Gurvits et al. (1999) data to set constraints on cosmological parameters. Guerra, Daly, \\& Wan (2000), Guerra \\& Daly (1998), Daly, Mory, \\& Guerra (2002), and Daly \\& Guerra (2002) examine FRIIb radio galaxy redshift-angular size cosmological constraints in various models using the modified standard yardstick method proposed by Daly (1994).  Here we use the FRIIb radio galaxy redshift-angular size data of Guerra et al. (2000) to derive constraints on the parameters of a spatially-flat model with a dark energy scalar field ($\\phi$) with scalar field potential energy density $V(\\phi)$ that at low redshift is $\\propto \\phi^{-\\alpha}$, $\\alpha > 0$ (Peebles \\& Ratra 1988). The energy density of such a scalar field decreases with time, behaving like a time-variable $\\Lambda$.  We adopt the analysis technique of Guerra et al. (2000), marginalizing over their parameter $\\beta$ to account for the uncertainty in the linear size of the ``standard rod\" used in the redshift-angular size test, to derive the likelihood (probability distribution) of the scalar field model parameters, $L(\\Omega_0, \\alpha)$. This likelihood function is used to determine confidence regions for the model parameters. We compute $L(\\Omega_0, \\alpha)$ over the ranges $ 0.05 \\leq \\Omega_0 \\leq 0.95$ and $ 0 \\leq \\alpha \\leq 8$. The radio galaxy 3C 427.1 (of the twenty used in the analysis) is a disproportionate contributor to $\\chi^2$, so we present results both including and excluding this radio galaxy. When we exclude 3C 427.1 we renormalize the error bars to make the best-fit reduced $\\chi^2$ unity. ",
        "conclusions": "Figures 1 and 2 show the constraints on $\\Omega_0$ and $\\alpha$ in the dark energy scalar field model with $V(\\phi) \\propto \\phi^{-\\alpha}$, including and excluding 3C 427.1. In both cases the constraints shown here are consistent with, but tighter than, those derived using the Gurvits et al. (1999) compact radio source redshift-angular size data (Chen \\& Ratra 2003, Fig. 3). They are also consistent with, but not as constraining as, those derived from the Riess et al. (1998) and Perlmutter et al. (1999) Type Ia supernova redshift-magnitude data (Podariu \\& Ratra 2000; Waga \\& Frieman 2000). Consistent with these analyses, the analysis here also does not rule out large values of $\\alpha$ when $\\Omega_0$ is small.  The radio galaxy 3C 427.1 can easily be identified as an outlier. This can be seen in Table 2 and Figures 2b, 8b, and 8c of Guerra \\& Daly (1998), and the effective Hubble diagrams shown in Figure 7 of Guerra et al. (2000) and Figure 8 of Daly \\& Guerra (2001), for example.  In the fits presented in each of these papers and those presented by Daly \\& Guerra (2002), as well as in the fits presented here, this one source contributes about one-half of the total $\\chi^2$, and the total $\\chi^2$ for the best fit parameters is always about 16 (see Figure 1 of Daly \\& Guerra 2002, for example). Each fit has 16 degrees of freedom since there are 20 radio galaxy points, the cosmological model has 2 free parameters, and the radio galaxy modified standard yardstick model has 2 free parameters.  To date, this point has been included in all of the fits. It is not clear whether it should be included in the fits, or whether it should be flagged as an outlier and removed from the data set, as was done for supernovae Type Ia by Perlmutter et al. (1999).  If it is flagged as an outlier and removed from the data set, the low reduced $\\chi^2$ of about one-half that would result for the best fitting parameters in any of the cosmological models considered would indicate that the error bars on each radio galaxy data point have been over-estimated by a factor of about 1.4.  It is certainly possible for these error bars to have been overestimated; the determination of the error bars is discussed in detail in the Appendix of Guerra et al. (2000).  This decrease in the error bar per point tightens the constraints on all parameters, including cosmological parameters and radio galaxy model parameters, as can be seen by comparing Figures 1 and 2. However, the only empirical basis for removing 3C 427.1 from the data set is it's position on the radio galaxy effective Hubble diagram.   It's radio structure is not unusual or remarkable (see Leahy, Muxlow, \\& Stephens 1989).  The use of this radio source to determine the ambient gas density (Wellman et al. 1997a), the ambient gas temperature (Wellman et al. 1998b), and the synchrotron aging independent ambient gas pressure, the beam power, and the total AGN-jet lifetime of this source (Wan, Guerra, \\& Daly 2000) are all unremarkable and quite in line with sources with similar redshifts, sizes, and radio powers.  The determination of the source redshift seems secure (Spinrad, Stauffer, \\& Butcher 1985). It does not appear to be an especially bright infrared source; upper bounds on its infrared luminosity are presented by Meisenheimer et al. (2001).  Thus, it is unlikely to be an especially highly obscured quasar that should have been removed from this sample of radio galaxies.   And, it's radio morphology and environmental properties are not unusual (Harvanek \\& Stocke 2002).  Thus, there does not seem to be any empirical basis to remove it from the radio galaxy sample other than the small value of its predicted average size $D_*$ relative to the average size of the full population of FRIIb radio galaxies $\\langle D \\rangle$ at similar redshifts.  Additional radio data could substantially improve the constraints presented here on the dark energy scalar field model, and on other models.  Of the 70 sources in the parent population, 20 have values of $D_*$ determined.  The radio data for these 20 radio sources were available in the published literature and in the VLA archive, and were originally obtained for studies other than those presented here.  Ten additional sources will be observed at the VLA this year; these data are optimized to deterime cosmological parameters and to test and constrain the underlying radio galaxy model.  The 40 remaining sources will then be observed, which will improve the radio galaxy constraints enormously.  The improvement will come from the additional number of data points and, with the very high quality data expected, the error bar per point may be smaller than that obtained using published data.  It is encouraging that the FRIIb radio galaxy redshift-angular size data constraints are consistent with and not much weaker than those derived from Type Ia supernova redshift-magnitude data. Future higher-quality redshift-angular size data is eagerly anticipated.  \\bigskip  We acknowledge helpful discussions with Joel Carvalho, Megan Donahue, Eddie Guerra, Philip Mannheim, Chris O'Dea, Adam Reiss, and the referee Dave Helfand.  We acknowledge support from NSF CAREER grant AST-9875031, NSF NYI grant AST-0096077, NSF grant AST-0206002, and Penn State University."
    },
    "0207/astro-ph0207575_arXiv.txt": {
        "abstract": "{ Using Fourier frequency resolved X-ray spectroscopy  we study short term spectral variability in luminous LMXBs. With RXTE/PCA observations of 4U1608--52 and GX340+0 on the horizontal/normal branch of the color-intensity diagram we show that aperiodic and quasiperiodic variability  on  $\\sim$ sec -- msec time scales  is caused primarily by variations of the luminosity of the boundary layer. The emission of the accretion disk is  less variable on these time scales and its power density spectrum follows $P_{\\rm disk}(f)\\propto f^{-1}$  law, contributing to observed flux variation at low frequencies and low energies only. The kHz QPOs have the same origin as variability at lower frequencies, i.e. independent of the nature of the \"clock\", the actual luminosity modulation takes place on the neutron star surface,    The boundary layer spectrum remains nearly constant in the course of the luminosity variations  and is represented to certain accuracy by the  Fourier frequency resolved spectrum. In the considered range $\\dot{M}\\sim (0.1-1) \\dot{M}_{\\rm Edd}$ it depends weakly on the global mass accretion rate and in the limit $\\dot{M}\\sim \\dot{M}_{\\rm Edd}$ is close to Wien spectrum with $kT\\sim 2.4$ keV (in the distant observer's frame). The spectrum of the accretion disk emission is significantly softer and in the 3--20 keV range is reasonably well described by a relativistic disk model with a mass accretion rate consistent with the value inferred from the observed X-ray flux. ",
        "introduction": "\\label{sec:intro}  It is commonly accepted that in non-pulsating neutron star X-ray binaries the magnetic field of the neutron star is  weak enough and the accretion disk can extend close to the surface of a neutron star. If the neutron star rotation frequency is smaller than the Keplerian frequency at the inner edge of the  disk, a boundary layer will be formed near the surface of the neutron star in which the accreting matter decelerates and spreads over star's surface \\citep{ss86,kluzniak,inogamov99,popham01}. For a non-rotating neutron star, in Newtonian approximation  half of the energy release due to accretion  would take place in the boundary/spread layer. The effects of the general relativity can  increase this fraction, e.g. up to $\\sim 2/3$  in the case of a neutron star with radius $R_{\\rm NS}=3 R_{\\rm g}$ \\citep{ss86,sibg00}. Rotation of the neutron star and deviations of the space-time geometry from Schwarzschild metric further modify the fraction of the energy released on the star's surface. Consequently, a luminous spectral component, corresponding to the  boundary layer emission is expected be present in the X-ray spectrum of a neutron star X-ray binary   X-ray observations of neutron star LMXBs in the high luminosity state reveal rather soft composite X-ray spectra.  Based on theoretical expectations they are usually represented as a sum of two components attributed to the optically thick emission of the accretion disk \\citep{ss73} and of the boundary layer/neutron star  surface \\citep{mitsuda84, white88}. The spectra of these two components are rather similar to each other and decomposition of the X-ray emission into the boundary layer and accretion disk components is often ambiguous, especially when based on the spectral information alone. Not surprisingly, the best fit parameters derived from the data of different instruments and, correspondingly, the inferred values of the physically meaningful quantities  are often in contradiction to each other \\citep[e.g.][]{mitsuda84, disalvo01, done02}. A robust and sufficiently model independent  method of separating the boundary layer and disk emission is of interest.    \\citet{mitsuda84} and \\citet{mitsuda86} studied pattern of spectral variability on the timescale of $\\sim 10^3$ sec in luminous LMXBs and found that the observed spectra can be represented as a sum of two components, having drastically different variability properties: a strongly variable $\\sim 2$ keV nearly blackbody component and a stable softer component. They interpreted the hard and soft components as emission from the neutron star surface and from the optically thick accretion disk respectively and concluded that the optically thick disk is stable on the timescales considered. The latter conclusion is in accord with the finding of \\citet{chur01}, who showed that the same is true in the high luminosity state of the black hole binary Cyg X-1. On  time scales $\\la 10^2$ sec the fractional amplitude of variations of the disk emission is at least an order of magnitude lower than that of the hard Comptonized component. These results indicate that a low level of variability might be an intrinsic property of the optically thick accretion disk, independent of the nature of the compact object.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{h3881f1.ps} \\caption{The color--intensity diagram of GX340+0. Dashed line polygons show the regions at the  Horizontal and upper half of the Normal Branch used for the frequency resolved analysis. \\label{fig:cid}} \\end{figure}   \\begin{figure} \\includegraphics[width=0.5\\textwidth]{h3881f2.ps} \\caption{Power spectrum of GX340+0 at the upper part of the Horizontal Branch of the color-intensity diagram. \\label{fig:pds_gx340}} \\end{figure}   As is well known \\citep[see][for review]{vdk86,vdk00}, aperiodic variability of X-ray flux from LMXBs can be broadly divided into two main phenomena -- continuum noise and quasi-periodical oscillations (QPO) with frequencies ranging from several mHz to more than  a thousand Hz \\citep[e.g.][]{has_vdk_89, vdk00, mikej_mhz}. \\citet{mitsuda84} studied the difference between the spectra averaged at different intensity levels -- that restricted the range of accessible time scales to $\\ga 10^3$ sec. In this paper we will exploit the technique of Fourier frequency resolved spectroscopy \\citep{freq_res99} to study spectral variability of luminous LMXBs on a broad range of time scales, including kHz QPO.      As defined in \\citet{freq_res99}, the Fourier frequency resolved spectrum is the energy dependent rms amplitude in a selected frequency range, expressed in absolute  (as opposite to fractional) units. A similar approach was used by \\citet{mendez0614} to study the energy spectrum of kHz oscillations in  4U0614+09. One of it's advantages over simple fractional rms--vs.--energy dependence is the possibility to use conventional (i.e. response folded) spectral approximations in order to describe the energy dependence of aperiodic variability. One should keep in mind, however, that the interpretation of the frequency resolved spectra often is not straightforward and might strongly depend on a priori assumptions. Nevertheless, several applications of this technique to variability of black hole binaries gave meaningful results \\citep[e.g.][]{freq_res99, gilfanov_cygx1}.    \\begin{figure*} \\hbox{ \\includegraphics[width=0.5\\textwidth]{h3881f3l.ps} \\includegraphics[width=0.5\\textwidth]{h3881f3r.ps} } \\caption{The low ({\\em left}) and high ({\\em right}) frequency parts of the power spectrum of 4U1608 averaged over all data used for analysis. The power spectrum of the high frequency part, showing two kHz QPO peaks was obtained by ``shift-and-add''  method. \\label{fig:pds_1608}} \\end{figure*}   The simplest situation, when frequency resolved spectra can be easily interpreted is illustrated by the following example. Consider a two-component spectrum, in which one component is stable and the normalization of the other varies while it's spectral shape is unchanged. In this case the shape of the frequency resolved spectrum would not depend on Fourier frequency and would be identical to the spectrum of the variable component. The spectrum of the non-variable component could, in principle, be  determined by subtracting the frequency resolved spectrum from the average spectrum with appropriate renormalization. Importantly, in this example the X-ray flux in all energy channels  will vary coherently and with zero time/phase lag between different energies. Presence of significant phase lag  and/or Fourier frequency dependence of the frequency resolved spectra would indicate that a more complex pattern of spectral  variability is taking place and interpretation is then less obvious. With few exceptions\\footnote{ \\label{page:cygx2_fotnote} $\\sim 150\\degr$ phase lag was detected in the normal branch QPO of Cyg X-2 with the pivot energy $\\sim 5$ keV, but  no such lags were found in a similar spectral state of Sco X-1 \\citep{dieters00} }, phase lag between light curves in different energy bands in luminous LMXBs (e.g. Sco X-1, GX5-1, 4U1608--52, 4U0614+091 etc.) is usually small, $\\Delta\\phi\\la{\\rm few}\\times 10^{-2}$, coherence is consistent with unity,  \\citep[e.g.][]{vaughan94, vaughan99, dieters00} and the fractional rms-vs-energy dependence similar at different Fourier frequencies \\citep{vdk86}. This suggests, that Fourier frequency resolved spectral analysis can be applied and its interpretation is sufficiently straightforward and model-independent.    The structure of the paper is as follows. We briefly describe the data in Sect.~\\ref{sec:data}. In Sect.~\\ref{sec:freqres_data} we present the results of the observations, show that the frequency resolved spectra do not depend upon the Fourier frequency, and constrain the time   lags between different energies. The initial observational results are summarized in Subsect.~\\ref{sec:obs_sum}. In Sect.~\\ref{sec:freqres_theory} we show that a particularly simple form of the spectral variability is required in order to satisfy the observational constraints -- the flux variations in different energy channels must be related by a simple linear transformation. In Sect.~\\ref{sec:freqres_applications} we compare the expected spectra of the  disk and boundary layer emission with the frequency resolved spectra. We show that the observed aperiodic and quasiperiodic variability is primarily caused by variations of the luminosity of the boundary layer and its energy spectrum can be represented, to certain accuracy, by the frequency resolved spectra. In Sect.~\\ref{sec:discussion} we discuss the boundary layer emission spectrum, it's dependence on the mass accretion rate and implications of these results for  disk and boundary layer models. The results are summarized in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}   The best fit values of the mass accretion rates obtained from the spectral fits agree, within a factor of $\\la 1.7$ with the values predicted from the observed energy flux and the  accretion efficiency expected for a neutron star with spin frequency of $\\sim 500$ Hz (Eq.~(\\ref{eq:mdot})). This fact is especially encouraging, as the two sources have accretion rates different by a factor of $\\sim 10$ (Table \\ref{tb:fit}). For such difference in the accretion rate the expected disk temperatures should differ by a factor for $\\sim 1.7-1.8$, resulting in the different shape and total flux of the disk spectra (Fig.~\\ref{fig:disk_bl}, lower panel). However, after the contribution of the boundary  layer is accounted for, the relativistic disk emission model is capable of reproducing both spectral shape and normalization.  Despite large difference in the mass accretion rate in the two sources, the energy spectra of the boundary layer emission are very similar to each other. This is in line with the finding that variations of the boundary layer luminosity in the broad range of time scales from  $\\sim$sec to $\\sim$msec are not accompanied by  significant variations of the spectral shape. Similar behavior was found by \\citet{mitsuda84} and \\citet{mitsuda86} on longer time scales of $\\sim 10^3$ sec.  Due to short light travel time of the accretion disk in the vicinity of the neutron star, $\\sim$msec, the reflected component, if originating in the inner disk, could contribute to the variable emission and cause  deviation of the frequency resolved spectra from the true boundary layer spectrum. This can not be directly verified with the present data --  upper limit on the equivalent width of the 6.4-6.7 keV line is in the case of both sources $\\approx 110$ eV (90\\% confidence). However, for the observed shape of the spectrum of the boundary layer, contribution of the reflected component with $\\Omega/2\\pi\\sim 0.2-0.3$, if any, would not exceed $10\\%$ in the $\\sim 10-20$ keV energy range, i.e. is comparable or smaller than other uncertainties involved.     Further along the Z-track of GX340+0 in the color-intensity diagram, on  normal and flaring branches,  the fractional rms of the X-ray variability decreases significantly, by a factor of $\\sim 5-10$. However, the statistics is sufficient to place meaningful constrains on the behavior of the frequency resolved spectra at the first half of the normal branch (Fig.~\\ref{fig:cid}). The data indicates that the behavior of the  frequency resolved spectra does not change its character -- at sufficiently high frequency, $f\\ga 1$ Hz, their shape does not depend upon the Fourier frequency and is significantly harder than the average spectrum and expected spectrum of the accretion disk. Therefore, in the same line of arguments as above, it is representative of the spectrum of the boundary layer spectrum. Fit to the frequency resolved spectrum by Comptonization model requires infinitely large values of the  Comptonization parameter. Correspondingly, the boundary layer spectrum in the normal branch  can be well fit by Wien or blackbody spectrum (which are close to each other in the 3-20 keV range)  with the best fit temperature of $kT\\approx 2.4$ keV. As evident from Fig.~\\ref{fig:disk_bl} and, especially, from Fig.~\\ref{fig:bl_alongz} the high energy part, E$\\ga 8$ keV of the spectrum of 4U1608--52 and horizontal branch of GX340+0 also follows Wien spectrum with temperature in the range $kT\\sim 2.1-2.3$. The composite fit of the total spectrum with the disk + boundary layer spectrum, the same as in Subsect.~\\ref{sec:fit}, gives a best fit value of the mass accretion rate of $\\dot{M}\\approx 4.6 \\cdot 10^{18}$ g/s, i.e. higher than in the horizontal branch (cf. Table \\ref{tb:fit}). This is consistent with the commonly accepted interpretation that the mass accretion rate increases along the Z-track on the  color-intensity diagram.  The  frequency resolved spectra ($\\approx$boundary layer spectra) on the normal and horizontal branch of the color-intensity diagram are plotted in Fig.~\\ref{fig:bl_alongz} along with Wien spectrum with $kT=2.4$ keV. Combined with the middle panel in Fig.~\\ref{fig:disk_bl} this plot shows trend in the dependence of the boundary layer spectrum upon the mass accretion rate in the range $\\dot{M}\\sim (0.1-1.0) \\dot{M}_{\\rm Edd}$. We can tentatively conclude that with increase of the mass accretion rate up to a value close to critical Eddington rate the boundary layer spectrum in the 3--20 keV energy  range approaches a Wien spectrum. Interestingly, at lower values of $\\dot{M}$ the character of the deviations of the boundary layer spectrum from the Wien spectrum is similar to that expected in the  situation when Compton scatterings are important factor of the spectral formation in the media with inhomogeneous temperature distribution \\citep{ross96}. In particular they are qualitatively  similar to the numerical results of \\citet{greb} on the  formation of the spectrum of the boundary layer.   \\begin{figure} \\includegraphics[width=0.5\\textwidth, clip]{h3881f12.ps} \\caption{The absorption corrected frequency resolved spectra of QPO ($\\approx$ boundary layer emission) in GX340 in the horizontal branch, Z=0--1 (lower $\\dot{M}$) and upper half of the normal branch, Z=1--1.5 (higher $\\dot{M}$). See Fig.~\\ref{fig:cid} for specification of the regions in the color-intensity diagram. The horizontal branch data is same as in Fig.~\\ref{fig:freqres_gx340}, \\ref{fig:blspe} and \\ref{fig:disk_bl}. The solid line shows  Wien spectrum with $kT=2.4$ keV. \\label{fig:bl_alongz}} \\end{figure}    The relatively weak dependence of the shape of the boundary layer spectrum upon the mass accretion rate (Fig.~\\ref{fig:disk_bl} and \\ref{fig:bl_alongz}) and the relative constancy of the Wien temperature is somewhat surprising. It implies that in the considered range of $\\dot{M}\\ga 0.1 \\dot{M}_{\\rm Edd}$ the plasma temperature at Comptonization depth in the boundary layer weakly depends upon the mass accretion rate, i.e. increase of the $\\dot{M}$ does not change significantly vertical temperature structure in the boundary layer.   The fact that kHz QPO show the same behavior as other components of the aperiodic variability indicates, that they have the same origin, i.e. are caused by the variations of the luminosity of the boundary layer. Although the kHz ``clock'' can be in the disk or due to it's interaction with the neutron star, the actual modulation of the X-ray flux  occurs on the neutron star surface. \\citet{mendez0614} suggested similar interpretation of kHz QPO in 4U0614+09. In particular they showed, that  energy spectrum of the kHz QPO can be approximately described by a blackbody spectrum with $kT\\sim 1.5-1.6$ keV. Note, however, that they found different energy depedence of continuum aperiodic variability at lower frequencies. That can possibly be explained by the fact that the source was in significantly lower luminosity state, $\\dot{M}\\sim 10^{-2}\\dot{M}_{\\rm Edd}$ and it's energy spectrum had a distinct power law component which could dominate variability at lower frequencies.  \\begin{figure} \\includegraphics[width=0.5\\textwidth, clip]{h3881f13.ps} \\caption{The  power density spectra of GX340+0 in the 3--7 keV energy range (horizontal branch, same data as in Fig.~\\ref{fig:freqres_gx340}, \\ref{fig:blspe} and \\ref{fig:disk_bl}). The histogram with the error bars show power spectra of total, disk and boundary layer emission, obtained using the  procedure described in the Sect.~\\ref{sec:discussion}; the straight solid line shows a power law $P(f)=3.5e-5\\times f^{-1}$. The power density is shown in units of fractional rms per Hz and all three spectra are normalized to the total count rate in the 3--7 keV band. Note, that the lower-most frequency bin is affected by the windowing effects, suppressing the power. \\label{fig:pds_disk_bl}} \\end{figure}      The disk emission is significantly less variable and does not contribute significantly to the variability of the X-ray flux at $f\\ga 0.5-1$ Hz. This conclusion is in agreement with results of \\citet{chur01} for the high state of Cyg X-1, indicating that stability of the X-ray emission might be a common property of the optically thick accretion disk, independently on the nature of the compact object. The origin of the variable component, however, is different in the case of Cyg X-1 (and presumably in the soft state of other black hole binaries). Indeed, the spectrum of the variable component in the soft state of Cyg X-1 is identical to the time average spectrum of the hard spectral component and is adequatly represented by unsaturated Comptonization in hot ($kT_{\\rm e}\\sim 50-100$ keV) and optically thin ($\\tau_{\\rm T}\\la 1$) coronal flow with possible contribution of non-thermal Comptonization. Variable component in luminous LMXBs considered in this paper, if interpreted in the framework of the Comptonization model,  requires saturated Comptonizaion (Comptonization parameter $y\\sim 1$) in the relatively low temperature ($kT_{\\rm e}\\sim 2-3$ keV) plasma and is inconsistent both qualitatively and quantitatively  with the corona models usually applied to black hole binaries.      In the bright LMXB systems, the contribution of the disk variability becomes noticeable at  lower frequencies, below $\\sim 0.5$ Hz, where the frequency resolved spectrum changes it's shape and becomes softer (Fig.~\\ref{fig:freqres_gx340} and \\ref{fig:freqres_gx340_ratios}; cf. results of \\citealt{vdk86} for GX5--1). This can be used to estimate the contribution of the disk to the observed variability of the X-ray flux. The result depends on the character  of the disk variability. In order to make a crude estimate we assume that the disk variations also obey a simple linear relation described by Eq.~(\\ref{eq:lc}): \\begin{eqnarray} F(e,t)=F_{\\rm disk}(E,t)+F_{\\rm BL}(E,t)\\approx\\\\ \\approx S_{\\rm disk}(E)\\times f_{\\rm disk}(t)+ S_{\\rm BL}(E)\\times f_{\\rm BL}(t) \\nonumber \\label{eq:lc_disk_bl} \\end{eqnarray} If disk and boundary layer variations were uncorrelated, the power density of the total signal would be: \\begin{eqnarray} P(E,\\omega)\\propto S_{\\rm disk}(E)^2 \\times |\\hat{f}_{\\rm disk} (\\omega)|^2+ S_{\\rm BL}(E)^2 \\times |\\hat{f}_{\\rm BL} (\\omega)|^2 \\nonumber \\label{eq:pds_disk_bl} \\end{eqnarray} where within the accuracy of this consideration one can assume that $S_{\\rm disk}(E)$ and $S_{\\rm BL}(E)$ are the disk and boundary layer spectra determined in Sect.~\\ref{sec:fit}. The functions $|\\hat{f}_{\\rm disk} (\\omega)|^2$ and  $|\\hat{f}_{\\rm BL} (\\omega)|^2$ after appropriate renormalization represent power density spectra of the disk and boundary layer and can be determined from linear fit to the square of the frequency resolved spectra in each frequency interval.  The power density spectra thus computed are shown in Fig.~\\ref{fig:pds_disk_bl} along with the total power spectrum  of GX340+0 in the soft 3--7 keV energy band. The power spectrum of the accretion disk flux variations is different from that of the boundary layer, does not extend significantly to the high frequency domain, and is consistent with a power law with slope of $-1$: $$ P_{\\rm disk}(f)\\propto f^{-1} $$ The excess power seen at low frequencies, $F\\la 0.5-1$ Hz in the soft energy band (cf. Fig.~\\ref{fig:pds_gx340}) can be explained as the contribution of the disk variations. At higher frequencies the variability is dominated by the boundary layer emission, giving primary contribution to quasi-periodic oscillations  and the so called band limited noise compoinent \\citep[e.g.][]{vdk86}.         The initial observational results are listed in Subsect.~\\ref{sec:obs_sum}. Below we summarize the constrains on the character of the spectral variability and implications for the boundary layer and accretion disk models.  \\subsection{Constrains on the pattern and origin of the spectral variability in luminous LMXBs}  \\begin{enumerate}  \\item Using RXTE/PCA observations of two luminous low mass X-ray binaries GX340+0 (on the horizontal/normal branch of the color-intensity diagram) and 4U1608-52 we show that the shape of the  Fourier frequency resolved spectra on $\\sim$ second -- millisecond time scales does not depend on Fourier frequency (Fig.~\\ref{fig:freqres_gx340}, \\ref{fig:freqres_gx340_ratios} and \\ref{fig:freqres_1608}). \\ The range of investigated  timescales includes the band limited continuum noise,  the kHz QPO and lower frequency QPOs observed at few tens Hz (Fig.~\\ref{fig:pds_gx340}, \\ref{fig:pds_1608}). Combined with the negligibly small phase lags, $\\Delta\\phi\\la 10^{-2}$ (Fig.~\\ref{fig:lags_gx340}), this restricts significantly the possible pattern of spectral variability of X-ray flux and requires linear relation between flux variations at different energies (Eq.~(\\ref{eq:lc})).   Considering significant difference in the expected spectra of the accretion disk and boundary layer the observed variations should be associated with either one of these two major components of the accretion flow. The X-ray variability is caused either by variations of it's luminosity under constant spectral shape, or by small variations of a spectral parameter (e.g temperature or optical depth) -- Eq.~(\\ref{eq:taylor}).   \\item We compared the Fourier frequency resolved spectra with the  expected spectra of the accretion disk and of the boundary layer. The predicted spectra were based on the observed energy flux/spectrum and very generic system parameters such as the source distance and neutron star spin frequency. The frequency resolved spectra are well consistent with the range of the boundary layer spectra expected for plausible range of the system parameters (Fig.~\\ref{fig:blspe}, Table \\ref{tb:disk_range}). On the other hand, they are significantly harder than the expected spectrum of the accretion disk. It is unlikely that the observed variations are associated with variations of the disk luminosity or spectral shape unless the disk temperature is $\\sim 3-4 $ kev, i.e. current accretion disk models are inapplicable to the neutron star binaries.   \\item The above suggests that the major part of aperiodic and quasiperiodic variability observed in luminous LMXBs above $\\sim 0.5$ Hz is caused by variations of the luminosity of the boundary layer. Its spectral shape remains nearly constant in the course of the luminosity variations. This interpretations receives additional support from the constancy of the fractional rms with energy at $E\\ga 10$ keV, where expected accretion disk emission vanishes, found in case of 4U1608--52 (Fig.~\\ref{fig:rms_1608}).  \\end{enumerate}  \\subsection{Implications for the models of the boundary layer and disk emission}  The frequency resolved spectrum is representative of the energy spectrum of the boundary layer emission. This can be used for a more precise decomposition of the spectra of luminous LMXBs into accretion disk and boundary layer components and for quantitative comparison with  predictions of the theoretical models. In the following we shall assume that boundary layer spectrum is identical to the  frequency resolved spectrum, bearing in mind that this is true to certain accuracy.  \\begin{enumerate}  \\item In the considered range of the mass accretion rate $\\dot{M}\\sim (0.1-1)\\dot{M}_{\\rm Edd}$, the boundary layer spectrum in the 3--20 keV energy range  depends weakly on $\\dot{M}$. Its shape is remarkably similar in GX340+0 and 4U1608--52 (Fig.~\\ref{fig:disk_bl}), despite the fact that the two sources have a factor of $\\sim10$ difference in the mass accretion rate (Table \\ref{tb:fit}).  In the limit of high $\\dot{M}$, of the order of $\\sim\\dot{M}_{\\rm Edd}$ (normal branch of GX340+0), the boundary layer spectrum in the 3--20 keV energy range  can be adequately  represented by the Wien spectrum with temperature  $kT\\approx 2.4$ keV (Fig.~\\ref{fig:bl_alongz}). At lower values of $\\dot{M}$ (4U1608--52 and horizontal branch of GX340+0) the spectra are better described by Comptonization model with electron temperature of $\\sim 2-4$ keV and Comptonization parameter $y\\sim 1$ (Table \\ref{tb:fit}). Their high energy part, $E\\ga 10$ keV, is well represented by Wien spectrum with temperature of $\\approx 2.1-2.3$ keV.   \\item The average spectra can be adequately described by the sum of the renormalized frequency resolved spectrum and the accretion disk emission (Fig.~\\ref{fig:disk_bl}). The spectrum of the latter is well described by the general relativistic accretion disk model. The other parameters, such as source distance and disk inclination angle being fixed at fudicial but plausible values, the best fit value of the mass accretion rate coincides, within a factor of $\\la 1.7$ with that inferred from the observed X-ray flux and accretion efficiency appropriate for a 1.4$M_{\\sun}$ neutron star with spin frequency of $\\sim 500$ Hz (Table \\ref{tb:fit}). This agreement is especially remarkable, given the luminosity and mass accretion rate in the two sources differ by the factor of $\\sim 10$.  \\item The accretion disk emission is significantly less variable than the boundary layer emission at Fourier frequencies $f\\ga 0.5-1$ Hz. The power density spectrum of the disk appears to follow a power law $P_{\\rm disk}(f)\\propto f^{-1}$ and contributes to the overall variability in the soft energy band and in the low frequency domain only (Fig.~\\ref{fig:pds_disk_bl}).  \\item The  kHz QPOs apear to have the same origin as aperiodic and quasiperiodic variability at lower frequencies. The msec flux modulations originate on the surface of the neutron star although the kHz ``clock''  might reside in the disk or be determined by the  disk -- neutron star interaction.  \\item Finally we point out that in the case of GX340+0 and presumably other Z--sources, the above results apply to the normal and horizontal branches of the color-intensity diagram. The source behaviour on the flaring branch, believed to correspond to super-Eddington accretion is more complex and is beyond the scope of this paper.  \\end{enumerate}"
    },
    "0207/astro-ph0207433_arXiv.txt": {
        "abstract": "We investigate the relationship between faint X-ray and 1.4~GHz radio source populations detected within 3$\\arcmin$ of the Hubble Deep Field North using the 1~Ms {\\it Chandra} and 40~$\\mu$Jy VLA surveys. Within this region, we find that $\\approx$42\\% of the 62 X-ray sources have radio counterparts and $\\approx$71\\% of the 28 radio sources have X-ray counterparts; thus a 40~$\\mu$Jy VLA survey at 1.4~GHz appears to be well-matched to a 1~Ms {\\it Chandra} observation. Among the different source populations sampled, we find that the majority of the 18 X-ray detected emission-line galaxies (ELGs) have radio and mid-infrared {\\it ISOCAM} counterparts and appear to be luminous star-forming galaxies at $z=0.3$--1.3. Importantly, the radio-detected ELGs make up $\\approx$35\\% of the X-ray source population at 0.5--8.0~keV X-ray fluxes between $\\approx(1$--$5)\\times10^{-16}$ \\xflux\\ and signal the emergence of the luminous, high-$z$ starburst galaxy population in the X-ray band. We find that the locally-determined correlation between X-ray luminosities and 1.4~GHz radio luminosity densities of the late-type galaxies can easily be extended to include the luminous intermediate-redshift ELGs, suggesting that the X-ray and radio emission processes are generally associated in star-forming galaxies. This result implies that the X-ray emission can be used as an indicator of star formation rate for star-forming galaxies. Finally, we show that there appear to be two statistically distinct types of {\\it ISOCAM}-detected star-forming galaxies: those with detectable radio and X-ray emission and those without. The latter type may have stronger mid-infrared emission-line features that increase their detectability at mid-infrared wavelengths. ",
        "introduction": "At bright X-ray fluxes ($\\ga10^{-15}$ \\xflux), almost all extragalactic X-ray point sources are found to be active galactic nuclei \\citep[AGN; e.g.,][]{Bade1998, Schmidt1998, Akiyama2000, Bauer2000, Lehmann2001}. At the faintest currently achievable levels, however, a large population of ``normal'' galaxies emerge \\citetext{\\citealp[e.g.,][]{Brandt2001a}, hereafter Paper~IV; \\citealp{Hornschemeier2001}, hereafter Paper~II; \\citealp{Tozzi2001}}. These X-ray faint ``normal'' galaxies are often optically bright ($I\\la22$) and are predominantly identified as narrow emission-line galaxies with redshifts of $z=0.3$--$1.3$. The majority appear to have faint 15$\\mu$m {\\it ISOCAM} counterparts \\citep[][hereafter Paper XI]{Alexander2002}, suggesting a close association with the strongly evolving luminous infrared (IR) starburst galaxy population \\citep[e.g.,][]{Elbaz2002}. X-ray/IR correlations have also been found locally for late-type normal spiral and irregular galaxies \\citep[e.g.,][]{Fabbiano1988, David1992, Shapley2001}. Thus the study of optically bright, X-ray faint galaxies may provide important clues about the nature of star formation on cosmic timescales.  An equally good correspondence between X-ray faint galaxies and faint radio sources is perhaps to be expected, given (1) the X-ray/radio correlation found for local late-type normal spiral and irregular galaxies \\citep[e.g.,][]{Fabbiano1988, Shapley2001, Ranalli2002}, and (2) the tight relationship between radio and mid-IR/far-IR (MIR/FIR) emission found for star-forming galaxies both locally and at moderate redshifts \\citep[e.g.,][]{Helou1985, Condon1991, Price1992, Shapley2001, Elbaz2002, Garrett2002}. Such a correlation can be understood physically, since X-ray and radio emission are both tied to the evolution of massive stars in the form of mass-transfer in binaries and supernovae, respectively \\citep[e.g.,][]{Petre1993, Bressan2002}. Indeed, within the Hubble Deep Field North itself \\citep[\\hbox{HDF-N};][]{Williams1996}, the faint radio and X-ray populations share a large overlap ($\\sim$70\\%; Paper~IV).  In this paper, we investigate in detail the observational properties of the faint X-ray/radio matched sources in the vicinity of the \\hbox{HDF-N} using two of the deepest surveys ever performed at their respective wavelengths, the 1~Ms {\\it Chandra} Deep Field North X-ray \\citep[\\hbox{CDF-N};][hereafter Paper~V]{Brandt2001b} and the 40~$\\mu$Jy VLA 1.4~GHz radio \\citep[][hereafter R00]{Richards2000} surveys. This work extends the results of Paper~IV to a larger area and to deeper X-ray fluxes, and it examines in more detail the relationships between the X-ray, radio, and IR emission. To facilitate comparison with previous X-ray/IR matching results in Paper~XI, we adhere to much of the procedure outlined in that study. We describe the selection of our X-ray and radio samples in $\\S$\\ref{sample}, compare and contrast these samples in $\\S$\\ref{properties} and $\\S$\\ref{categories}, and discuss implications in $\\S$\\ref{discussion}. Throughout this paper, we adopt $H_{0}=65$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=\\onethird$, and $\\Omega_{\\Lambda}=\\twothirds$. The Galactic column~density toward the \\hbox{CDF-N} is $(1.6\\pm0.4)\\times10^{20}$ cm$^{-2}$ \\citep{Stark1992}. Coordinates are for the J2000 epoch. ",
        "conclusions": "Our main results are the following:  \\begin{itemize} \\item We find a large overlap between faint X-ray and radio sources detected within $3\\arcmin$ of the \\hbox{HDF-N} using the 1~Ms \\hbox{CDF-N} and 40~$\\mu$Jy VLA 1.4~GHz radio surveys. The matched sources exhibit a broad range of emission mechanisms, in agreement with the findings of Papers~II and IV.  \\item The highest matching fraction is found among X-ray- and radio-detected ELGs, which are comprised of apparently normal and starburst galaxies at redshifts of $z\\sim0.1$--1.3 thought to be undergoing recent star formation. The ELGs are distinct in terms of their X-ray, optical, and radio properties, and nearly all of them have MIR {\\it ISOCAM} detections.  \\item The radio-detected ELGs make up $\\approx$~35\\% of the X-ray source population at X-ray fluxes below $5\\times10^{-16}$ \\xflux\\ and further signal the emergence of a population of star-forming galaxies discovered in previous X-ray studies. The high matching fraction among the ELGs suggests that the deep 40~$\\mu$Jy VLA survey at 1.4~GHz is extremely well-matched to the 1~Ms {\\it Chandra} observation for the detection of galaxies.  \\item The X-ray-emitting AGN population (i.e., OBAGN $+$ OFXs) have significantly fewer radio matches than the ELGs. These sources exhibit a range of X-ray spectral slopes suggesting many are moderately obscured. Moreover, the radio properties and optical morphologies of the harder X-ray AGN indicate that their obscuration at X-ray energies may be related to active star formation (perhaps circumnuclear). The relative numbers of such starburst-obscuring AGN, however, suggest that they are not likely to be dominant contributors to the XRB.  \\item A good correlation exists between X-ray luminosity and radio luminosity density for the CDF-N ELGs and nearby late-type galaxies, and it suggests that the emission mechanisms from the evolution of massive stars (e.g., the creation of HMXBs) and their eventual destruction (e.g., supernova-accelerated electrons) are intimately linked, even at moderate redshifts. This link suggests that X-ray emission can likewise be used as an indicator of the SFR. The implied SFRs found for the X-ray-detected ELGs in the CDF-N are $\\sim$~0.2--450~$M_{\\odot} {\\rm yr}^{-1}$.  \\item Finally, there appear to be two statistically distinct classes of infrared star-forming galaxies detected with {\\it ISOCAM}, one of which shows correlated X-ray and radio emission, and one which does not. This latter group is likely to have significant stronger emission-line features than the those with X-ray and radio counterparts. Confirmation of this result must await the Guaranteed Time and Great Observatories Origins Deep Survey\\footnote{See http://www.stsci.edu/science/goods/.} observations of the \\hbox{CDF-N} with SIRTF, which will offer both deep spectroscopy and multi-band photometry at MIR wavelengths over a much larger area. We note that the X-ray exposure of the \\hbox{CDF-N} has recently increased to 2~Ms. Thus, the robust X-ray/radio matching fraction we find with the 1~Ms \\hbox{CDF-N} dataset also signals the need for deeper radio observations with the VLA to build upon this work.  \\end{itemize}"
    },
    "0207/astro-ph0207119_arXiv.txt": {
        "abstract": "Studies of absorption spectra of high-$z$ QSOs have revealed that the intergalactic medium at $z\\sim 2-3$ is enriched to $\\sim 10^{-3}-10^{-2}\\zsol$ for gas densities more than a few times the mean cosmic density, but have not yet produced an accurate metallicity estimate, nor constrained variations in the metallicity with density, redshift, or spatial location.  This paper discusses the ``pixel optical depth'' (POD) method of QSO spectrum analysis, using realistic simulated spectra from cosmological simulations.  In this method, absorption in $\\lya$ is compared to corresponding metal absorption on a pixel-by-pixel basis, yielding for each analyzed spectrum a single statistical correlation encoding metal enrichment information.  Our simulations allow testing and optimization of each step of the technique's implementation.  Tests show that previous studies have probably been limited by \\CIV\\ self-contamination and \\OVI\\ contamination by \\HI\\ lines; we have developed and tested an effective method of correcting for both contaminants.  We summarize these and other findings, and provide a useful recipe for the POD technique's application to observed spectra.  Our tests reveal that the POD technique applied to spectra of presently available quality is effective in recovering useful metallicity information even in underdense gas. We present an extension of the POD technique to directly recover the intergalactic metallicity as a function of gas density. For a given ionizing background, both the oxygen and carbon abundance can be measured with errors of at most a factor of a few over at least an order of magnitude in density, using a single high-quality spectrum. ",
        "introduction": "\\label{sec-intro}  Studies of absorption lines in QSO spectra have established that at high redshift ($z \\ga 2$) the intergalactic medium (IGM) is enriched by metals to metallicity $10^{-3} \\la Z/\\zsol \\la 10^{-2}$ at densities above a few times the cosmic mean (e.g., Tytler et al.\\ 1995; Cowie et al.\\ 1995; Songaila \\& Cowie 1996; Rauch, Haehnelt, \\& Steinmetz 1997; Dav\\'e et al.\\ 1998).  Since stars are thought to form in appreciable quantities only inside much denser mass condensations, the observed enrichment indicates that galaxies, which form from the IGM, also feed back some of their nucleosynthetic products in a process which is currently poorly understood but which is likely to be a crucial ingredient of galaxy formation and evolution.  Because galaxies form in relatively high-density regions, metals in low-density regions must have traveled a significant distance.  This puts strong constraints on the intergalactic (IG) enrichment mechanism, therefore determining the IG metallicity at the lowest possible densities is crucial for understanding the feedback of metals from galaxies to the IGM (see, e.g., Aguirre et al.\\ 2001; Madau, Ferrara \\& Rees 2001).  Previous studies, performed primarily with the Keck HIRES instrument, have used measured metal line column densities to show that, at $z\\approx 3$, most absorbers with $\\log N({\\rm H\\,I}) > 14.5$ (corresponding to gas overdensities $\\delta \\equiv \\rho/\\langle\\rho\\rangle \\ga 5-10$ [Schaye 2001]) have associated carbon lines (e.g., Tytler at al. 1995; Cowie et al.\\ 1995; Songaila \\& Cowie 1996; Ellison et al. 2000).  These studies have not, however, made clear how this metallicity changes with spatial location, redshift, or gas density.  Studies employing line-fitting could in principle treat the first two issues.  But pushing to lower gas densities where individual metal lines are undetectable requires a statistical analysis; at high redshift this is also required for metals absorption falling blue-wards of $\\lya$, where the metals lines are difficult to disentangle from \\HI\\ absorption.  Ellison et al.\\ (1999, 2000; see also Cowie \\& Songaila 1998) have compared two ways of doing this: first by stacking the spectra of the metal-line regions corresponding to a set of low column density $\\lya$ lines, or second by comparing the $\\lya$ optical depth in each pixel with the optical depth in the pixel at the corresponding metal-line wavelength.  The latter ``pixel optical depth'' (POD) method, which was devised and first used by Cowie \\& Songaila (1998), has several advantages: it is fast, it is objective, it preserves more of the spectrum's information, it can be applied to heavily contaminated regions, and it appears to be more sensitive to metals in low-density gas, even in uncontaminated regions.  Its shortcoming, however, has been that it is not immediately clear how to interpret the results of the POD method, and it is unclear how to convert the set of PODs into a metallicity $Z$ at a given overdensity $\\delta$, especially when contamination is important, as for \\OVI\\ (see Schaye et al.\\ 2000a).  The purpose of the present study is to present several important improvements to and extensions of the POD technique, to provide a complete and useful explanation of the method, and to test it using realistic spectra generated from hydrodynamic cosmological simulations. The method as described here will be applied to a sample of observed QSO spectra in a future paper. The present paper should also serve as a useful reference for prospective users of the POD technique.  The testing we perform allows us to estimate the accuracy with which \\HI\\ and metal line PODs can be recovered (and to refine the method for doing so), and to draw the connection between the optical depths and $Z(\\delta)$ for the gas. By comparing recovered quantities directly to the true (simulated) ones, we demonstrate that the POD method does work, and that we can interpret directly the meaning of features in the POD statistic.  Section~\\ref{sect-overview} contains the basics of the method, pointing forward to figures in the main text, for illustration. This section is meant for readers not interested in all the details, or as a general overview for readers unfamiliar with pixel statistics.  The full details of the method, and of the tests we performed, can be found in the following sections.  In \\S~\\ref{sec-spgen} we describe the simulations used in this study and how we generate spectra from them. In \\S~\\ref{sec-podmeth} we give a careful account of the recovery of metal and \\HI\\ PODs from QSO spectra, testing the individual steps by comparing recovered and ``true'' optical depths; this section can be used as a reference for the POD method of spectrum analysis, and is summarized in the Appendix. Section~\\ref{sec-interp} discusses and tests the interpretation of the recovered PODs: Sections~\\ref{sec-hiinterp} and~\\ref{sec-zinterp} relate the \\HI\\ and metal PODs to the density of the absorbing gas, and \\S~\\ref{sec-restest} shows the relation between metal PODs and the assumed gas metallicity, given various levels of noise and other uncertainties.  In \\S~\\ref{sec-invert} we formulate and test a procedure whereby $Z(\\delta)$ can be recovered directly and perform simple tests of this procedure. Finally, we discuss our results and conclude in \\S~\\ref{sec-conc}. ",
        "conclusions": "\\label{sec-conc}  Studies of the absorption spectra of QSOs at $2 \\la z \\la 4$ have revealed that the intergalactic medium is enriched with metals to the level of $\\sim 0.1-1\\%$ solar metallicity at overdensities $\\delta \\ga 5$. However, a more accurate estimate of the metallicity and its spatial variation, as well as the extent to which it varies with gas density and redshift, are currently unknown.  This paper is the first in a series that systematically analyzes a sample of QSO absorption spectra in an effort to glean as much information as possible about the enrichment of low-density gas at high redshift.  In the present paper, we have systematically described, tested, and discussed the ``pixel optical depth\" (POD) method of analyzing metal absorption in QSO spectra.  This method, while still less widely used than direct line-fitting techniques, is better able to extract information from low-density regions in which individual metal lines are difficult or impossible to detect.  The POD technique was developed and previously employed in QSO absorption studies of \\CIV\\ (Cowie \\& Songaila 1998; Ellison et al.\\ 2000) and \\OVI\\ (Schaye et al.\\ 2000a), but has suffered from difficulties in the interpretation of the results. Here we overcome this difficulty by testing the method on spectra generated from realistic cosmological hydrodynamical simulations that are able to reproduce the observed statistics of the Ly$\\alpha$ forest.\\footnote{It should be noted that this is {\\em all} that is required of the simulations, i.e., the results of this paper are not dependent upon the detailed accuracy of the cosmological simulations, but only on the simulations' proven ability to roughly match observed QSO absorption spectra.}  This has allowed us to both assess the effectiveness of the method, and to refine and calibrate it. The major improvement to the method developed and tested here is the removal of higher-order Lyman lines from the \\OVI\\ region, and the correction of self-contamination of \\CIV\\ by its own doublet.  These innovations significantly improve the accuracy of the \\OVI\\ and the \\CIV\\ recovery.  In the Appendix we gave a compact but complete recipe for implementing the method.  The most general result of our tests is that the POD technique is very effective at recovering information about the metallicity of low-density gas, even where (as for \\OVI) the metal lines are severely contaminated.  More specifically, our tests reveal the following:  \\begin{enumerate} \\item{Ly$\\alpha$ optical depths of up to several hundred can be reliably determined using higher order Lyman lines, and these optical depths in turn give an accurate estimate of the density of gas giving rise to the absorption.} \\item{With high-quality spectra (signal-to-noise $\\ga 50$) median \\CIV\\ optical depths binned in $\\lya$ optical depth can be accurately recovered down to $\\tau_{\\rm CIV} \\sim 10^{-3}$ for both $z_{\\rm qso}=3.5$ and $z_{\\rm qso}=2.5$.  Previous analyses were probably limited by self-contamination; with our correction for this effect, the recovery is limited by noise and by errors in the continuum fitting.} \\item{The recovery of median \\OVI\\ PODs is limited primarily by contamination due to \\HI\\ lines (and secondarily by continuum fitting errors). A significant fraction of this contamination can be removed, allowing a fairly accurate recovery of \\OVI\\ PODs with $\\tau_{\\rm OVI} \\ga 10^{-2}$ for $z_{\\rm qso}=2.5$, and a useful recovery of median PODs for $z_{\\rm qso}=3.5$.} \\item{The \\SiIV\\, (limited primarily by \\CIV\\ contamination) and \\NV\\, (limited primarily by $\\lya$ contamination) PODs can be usefully recovered for $\\lya$ optical depths ($\\tau_{\\lya} \\ga 1)$.} \\item{Using the POD method both \\CIV\\ and \\OVI\\ should be detectable in realistic spectra of $z\\sim 3$ QSOs even in gas of the mean cosmic density, if the metallicity is $\\sim 10^{-3}\\zsol$, and at even lower density if the metallicity is higher.} \\item{When applied to simulated spectra with different metallicities as a function of gas density, the method is able to distinguish models in which the metallicity is constant from those in which the metallicity declines significantly with gas density.  The results are {\\em not} very sensitive to spatial variations in metallicity at a given gas density.} \\item{Most important for accurate recovery of metallicity information are large wavelength coverage (so that a significant number of higher order Lyman lines can be used), accurate continuum fitting, and high signal-to-noise.  However very high S/N is only useful for \\CIV\\ recovery; obtaining a $S/N \\gg 50$ probably does not significantly improve the amount of information that can be recovered regarding \\OVI.} \\item{Due to different levels of thermal broadening in metals vs.\\ hydrogen, the inferred density/temperature of the gas is, somewhat different when weighted by metal line optical depth than when weighted by $\\lya$ optical depth, especially at relatively high gas densities.  However, this uncertainty does not appear to significantly affect the derived metallicities.} \\item{Using spectra generated from cosmological simulations, the POD method may be extended, as described in~\\S~\\ref{sec-invert}, to directly and accurately recover the metallicity of intergalactic gas vs. its density. The largest uncertainty in this ``inversion'' is the spectral shape (and possible spatial variations) of the ionizing background. For a given background, the abundances of C and O can be measured with errors of at most a factor of a few over at least an order of magnitude in density, using a single high-quality spectrum.}  \\end{enumerate}  The POD technique, as refined and calibrated in our study should, when applied to QSO spectra of currently-available quality, yield clear information about the metallicity of even underdense gas.  This information will place strong contraints on models of the enrichment of the IGM, and consequently on feedback and galaxy formation."
    },
    "0207/astro-ph0207605_arXiv.txt": {
        "abstract": "We present near infrared images in H$_2$ at 2.12\\mum~of the HH 7/11 outflow and its driving source SVS 13 taken with {\\it Hubble Space Telescope} NICMOS 2 camera, as well as archival \\Ha~and [\\ion{S}{2}] optical images obtained with the WFPC2 camera. The NICMOS high angular resolution observations confirm the nature of a small scale jet arising from SVS 13, and resolve a structure in the HH 7 working surface that could correspond to Mach disk H$_2$ emission. The H$_2$ jet has a length of 430 AU (at a distance of 350 pc), an aspect ratio of 2.2 and morphologically resembles the well known DG Tau optical micro-jet. The kinematical age of the jet ($\\sim 10$ yr) coincides with the time since the last outburst from SVS 13. If we interpret the observed H$_2$ flux density with molecular shock models of 20-30 \\kms, then the jet has a density as high as $10^5$ \\cc. The presence of this small jet warns that contamination by H$_2$ emission from an outflow in studies searching for H$_2$ in circumstellar disks is possible. At the working surface, the smooth H$_2$ morphology of the HH 7 bowshock indicates that the magnetic field is strong, playing a major role in stabilizing this structure. The H$_2$ flux density of the Mach disk, when compared with that of the bowshock, suggests that its emission is produced by molecular shocks of less than 20 \\kms. The WFPC2 optical images display several of the global features already inferred from groundbased observations, like the  filamentary structure in HH 8 and HH 10, which suggests a strong interaction of the outflow with its cavity. The H$_2$ jet is not detected in [\\ion{S}{2}] or \\Ha~, however, there is a small clump at $\\sim 5$ \\arcsec~NE of SVS 13 that could be depicting the presence either of a different outburst event or the north edge of the outflow cavity. ",
        "introduction": "\\label{intro}  Herbig-Haro (HH) objects trace optically the mass loss process from young stellar objects (YSOs) and their interaction with the surrounding medium (Reipurth  \\& Bally~\\citeyear{bo01}). Because of their morphology, energetics and size, HH objects are an integral part of the star formation process and its effect on molecular clouds. HH 7/11 is a chain of HH objects in Herbig's photographic plate catalogue (Herbig~\\citeyear{her74}), located in the very active star forming region NGC 1333 (Aspin et al.~\\citeyear{ASR94}; Bally, Devine \\& Reipurth~\\citeyear{betal96}) at a distance of 350pc (Herbig \\& Jones~\\citeyear{her83}). From the ground, at optical wavelengths, the HH 7/11 system is defined by an arch-shaped morphology (blue lobe) that spans $\\sim$ 2\\arcmin. The red-shifted counter-lobe is detected in the near infrared (NIR), e.~g. at 2.12\\mum~in the H$_2$ (1,0) S(1) line, and displays a more chaotic structure (Stapelfeldt et al.~\\citeyear{sta91}; Garden, Russell \\& Burton~\\citeyear{gar90}). Recent interferometric observations  (Bachiller et al.~\\citeyear{bach00}) have convincingly demonstrated that SVS 13 (Strom, Vrba \\& Strom~\\citeyear{str76}) is the driving source of the HH 7/11 outflow. SVS 13 has a Class 0/I spectral energy distribution (Bachiller et al.~\\citeyear{bach98}) and a luminosity of $\\sim 85$~\\lsol~ (Molinari, Liseau \\& Lorenzetti~\\citeyear{mol93}).  Detailed optical spectroscopic observations of HH 7/11 show a complex velocity field and a low excitation nature, consistent with shock velocities of 30-60 \\kms~(Solf \\& B\\\"ohm~\\citeyear{sb87}; B\\\"ohm \\& Solf~\\citeyear{khb90}). NIR spectroscopy displays a rich H$_2$ vibrational spectra and strong [Fe~II] 1.257 and 1.644\\mum~lines (Gredel~\\citeyear{gre96}; Everett~\\citeyear{eve97}), again consistent with collisional excitation by shocks. Far infrared spectroscopy, which includes the emission of molecular species (like H$_2$, H$_2$O and CO) and atomic fine structure lines (like [O~I] 63\\mum~and [Si~II] 34.8\\mum), requires both J-type and C-type shocks of 15-30 \\kms~to explain their ratios (Molinari et al.~\\citeyear{mol00}).  In this paper we present new high angular resolution (FWHM=0.1\\arcsec) {\\it Hubble Space Telescope} (HST) NICMOS H$_2$ images at 2.121\\mum~of HH 7/11 and its source SVS 13, as well as archive optical images in \\Ha~and [\\ion{S}{2}] taken with the WFPC2 camera at a similar epoch. ",
        "conclusions": "\\label{concl}  We have resolved two remarkable features in the molecular Hydrogen emission of the HH 7/11 outflow thanks to HST NICMOS images at 2.12\\mum: a jet with a length of 430 AU arising from the SVS 13 driving source, and the Mach disk in HH 7, leading working surface. These observations strongly support the presence of small-scale H$_2$ jets arising from Class I/O sources (Davis et al.~\\citeyear{chris02}), and open the possibility that a jet shock can be detected in H$_2$.  Using previously published infrared spectroscopic observations, coupled with molecular shock models, we have determined that: (i) the jet can have a density as high as $10^5$ \\cc, for shock velocities of 20-30 \\kms, (ii) the magnetic field plays a major role in stabilizing the HH 7 bowshock, (iii) the Mach disk H$_2$ emission is probably produced by shocks of less than 20 \\kms; and (iv) the complex distribution of the atomic and molecular gases in HH 7/11, depicted by the NICMOS and WFPC2 images, coupled with its kinematics, suggests a strong interaction of this outflow with its circumstellar medium.  \\begin{center} {\\it Acknowledgments} \\end{center} A.N-C. research is partially supported by NASA through a contract with Jet Propulsion Laboratory, California Institute of Technology, and ADP Grant NRA0001-ADP-096. We thank A. Moro-Mart{\\'{\\i}}n for critical and useful comments."
    },
    "0207/astro-ph0207449_arXiv.txt": {
        "abstract": "We develop a comprehensive quantitative description of the cross-section mechanism discovered several years ago by Lazarian. This is one of the processes that determine grain orientation in clouds of suprathermal cosmic dust. The cross-section mechanism manifests itself when an ensemble of suprathermal paramagnetic granules is placed in a magnetic field and is subject to ultrasonic gas bombardment. The mechanism yields dust alignment whose efficiency depends upon two factors: the geometric shape of the granules, and the angle $\\,\\Phi\\,$ between the magnetic line and the gas flow. We calculate the quantitative measure of this alignment, and study its dependence upon the said factors. It turns out that, irrelevant of the grain shape, the action of a flux does not lead to alignment if $\\,\\,\\Phi\\,=\\,\\arccos (1/\\sqrt{3})$. ",
        "introduction": "Polarisation of the starlight is a long-known effect. Due to the correlation of polarisation with reddening, this phenomenon is put down to the orientation of particles in the cosmic-dust nebulae (Hall 1949, Hiltner 1949). This orientation causes differential extinction of electromagnetic waves of different polarisations (effect known in optics as linear dichroism), and provides a remarkable example of order emerging in a seemingly chaotic system.  In a nutshell, the polarization is due to the fact that the grains are non-spherical (i.e., have different cross sections in the body frame) and these non-spherical cross sections are somehow aligned within the cloud. A remarkable feature of this  phenomenon is that, whatever orientational mechanisms show themselves in the dust-grain rotational dynamics, the orientation always takes place relative to the interstellar magnetic field.  Thus, whenever the word ``orientation'' is used, it always is an euphemism for ``orientation with respect to the magnetic line''. Another semantic issue is the conventional difference between the meaning of words ``orientation'' and ``alignment''. Historically, the word ``orientation'' has been used to denote existence of {\\it one} preferred direction in a physical setting. The analysis of cosmic-dust orientation has demonstrated that in all known cases this preferred direction is equivalent to its opposite. Simply speaking, were an instantaneous inversion of the magnetic field in the dust cloud possible, it would not alter the orientation of the dust ensamble. Because of such an invariance, the word ``orientation'' is often avoided and substituted by the word ``alignment'' which is thereby imparted with the desired broader meaning. (The adjective ``orientational'', though, remains in use.) We shall abide by this verbal code.  The rotational dynamics of an interstellar particle is determined by a whole bunch of accompanying physical processes whose combination produces a variety of orientational mechanisms. Which of these come into play in a particular physical setting, depends upon the suprathermality of the dust cloud. Suprathermal are, by definition, grains which spin so rapidly that their averaged (over the dust ensemble) rotational kinetic energy $\\,<E_{\\small{rot}}>\\,$ much exceeds the (multiplied by the Bolzmann constant $\\kappa$) temperature $T_{\\small{gas}}$ of the surrounding environment. The suprathermality degree is then introduced as the following ratio: \\be \\beta\\;=\\;<E_{\\small{rot}}>/{\\kappa}T_{\\small{gas}}\\;\\;\\;\\;.\\;\\; \\label{1.1} \\ee Dust ensembles with $\\,\\beta\\,$ not much different from unity are called thermal or Brownian. Clouds with $\\,\\beta\\,\\gg\\,1$ are called suprathermal. In the observable Universe, values of $\\,\\beta\\,$ of order $\\,10^2\\,$ are not unusual.  The leading reason for suprathermal rotation is formation of $\\bf{H_2}$ molecules at the defects on the granule surface: over such a defect (called active site), two atoms of $\\bf{H}$ couple to form a molecule, ejection whereof applies an uncompensated torque to the granule surface (Purcell 1979). These, so-called spin-up torques keep emerging at each active site until the site gets ``poisoned'' through the everlasting accretion. After that, some other active site will dominate the spin dynamics of the grain, by its $\\bf{H_2}$ ``rocket''. This change of spin state will, with some probability, go through a short-term decrease, to thermal values, of the grain's angular velocity. Such breaks are called ``cross-overs'' or ''flip-overs''.  Another pivotal circumstance to be mentioned here is the existing evidence of paramagnetic nature of a considerable portion of dust particles (Whittet 1992) which makes them subject to the Barnett effect. This phenomenon takes place in para-~~and ferromagnetics, due to interaction between the spins of unpaired electrons and the macroscopic rotation of crystal lattice (Stoner 1934). This coupling has its origin in the angle-dependent terms in the dipole-dipole interaction of neighbouring spins. It spontaneously endows a rotating para- or ferromagnetic body with a magnetic moment parallel to the angular velocity\\footnote{Another contribution to the magnetisation comes from the electric charge carried by the granule.} (Lazarian \\& Roberge 1997). Purcell offered the following illustrative explanation of the effect. If a rotating body contains equal amount of spin-up and spin-down unpaired electrons, its magnetisation is nil. Its kinetic energy would decrease, with the total angular momentum remaining unaltered, if some share of the entire angular momentum could be transferred to the spins by turning some of the unpaired spins over (and, thus, by dissipating some energy). This potential possibility is brought to pass through the said coupling.  An immediate outcome from granule magnetisation is the subsequent coupling of the magnetic moment $\\;\\bf M\\;$ with the interstellar magnetic field $\\;\\bf B\\;$: the magnetic moment precesses about the magnetic line. What is important, is that this precession goes at an intermediate rate. On the one hand, it is slower than the grain's spin about its instantaneous rotation axis. On the other hand, the precession period is much shorter than the typical time scale at which the relative-to-$\\bf B\\;$ alignment gets established\\footnote{ A more exact statement, needed in Section V below, is that the period of precession (about $\\;\\bf B\\;$) of the magnetic moment $\\;\\bf M\\;$ (and of $\\,Z\\,$ aligned therewith) is much shorter than the mean time between two sequent flip-overs of a spinning granule (Purcell 1979, Roberge et al 1993)}. The latter was proven by Dolginov \\& Mytrofanov (1976), for magnetisation resulting from the Barnett effect, and by Martin (1971), for magnetisation resulting from the grain's charge.  If we disembody the core idea of the Barnett effect from its particular implementation, we shall see that it is of quite a general sort: a free rotator, though conserves its angular momentum, tends to minimise its kinatic energy through some dissipation mechanism(s). This fact, neglected in the Euler-Jacobi theory of unsupported top, makes their theory inapplicable at time scales comparable to the typical time of dissipation (Efroimsky 2002). The needed generalisation of the insupported-top dynamics constitutes a mathematically involved area of study (Efroimsky 2000), which provides ramifications upon wobbling asteroids and comets (Efroimsky 2001), rotating spacecraft, and even precessing neutron stars (Tr\\\"{u}mper et al 1986, Alpar \\& \\\"{O}gelman 1987, Bisnovatyi-Kogan \\& Kahabka 1993, Stairs et al 2000). Fortunately, in the case of cosmic-dust physics, we need only some basics of this theory. A free rotator has its kinetic energy minimised (with its angular momentum being fixed) when the rotation axis coincides with the axis of major inertia. In this, so-called principal state, the major-inertia axis $\\,Z\\,$, the angular velocity $\\;\\bf \\Omega\\,$, and the angular-momentum vector $\\;\\bf J\\;$ are all aligned. In other, complex rotation states, both the maximal-inertia axis $\\,Z\\,$ and angular-velocity $\\;\\bf \\Omega \\,$ precess about the angular momentum $\\;\\bf J\\;$. This precession is also called ``wobble'' or ``tumbling'', in order to distinguish it from the precession of the magnetic moment $\\;\\bf M\\;$ about the magnetic line. Similarly, the wobble relaxation (i.e., gradual alignment of axis $\\,Z\\,$ and of $\\;\\bf \\Omega\\,$ toward $\\;\\bf J\\;$) should be distinguished from the granule alignment relative to the magnetic field $\\;\\bf B\\;$: the latter effect is the eventual target of our treatise, while the former is merely a trend in the tapestry. Still, the wobble relaxation is far more than a mere technicality: it is important to know if a typical time of the wobble relaxation is much less than the typical times of the external interactions (like, say, the period of precession of $\\;\\bf M\\;$ about $\\;\\bf B\\;$). In case the wobble relaxation is that swift, one may assume that the precession (about $\\;\\bf B\\;$) of the magnetic moment $\\;\\bf M\\;$ is the same as precession of the angular momentum $\\;\\bf J\\;$ about $\\;\\bf B\\;$: indeed, in this case, both $\\;\\bf J\\;$ and the major-inertia axis $\\,Z\\,$ will be aligned with $\\;\\bf \\Omega\\,$ which is parallel to $\\;\\bf M\\;$. This parallellism of all four vectors is often called not alignment but ``coupling'', to distinguish it from the alignment relative to $\\;\\bf B\\;$. This coupling is enforced by two different processes. One is an effect kin to that of Barnett: tumbling of $\\;\\bf \\Omega\\;$ relative to conserved\\footnote{At time scales shorter than a precession cycle about $\\;\\bf B\\;$.} $\\;\\bf J\\;$ (and, therefore, relative to an inertial observer) yields periodic remagnetisation of the material, which results in dissipation. The other effect is the anelastic dissipation: in a complex rotational state the points inside the body experience time-dependent accelaration that produces alternate stresses and strains. Anelastic phenomena entail inner friction (which may be understood also in terms of a time lag between the strain and stress). The contributions from the Barnett and anelastic effects to the coupling were compared by Purcell in his long-standing cornerstone work (Purcell 1979). Purcell came to an unexpected conclusion that the input from the Barnett effect much outweights that from anelasticity. Having set out the calculations, Purcell continued with phrase: ``It may seem surprising that an effect as feeble, by most standards of measure, as the Barnett effect could so decisively dominate the grain dynamics''. After that, Purcell tried to pile up some qualitative evidence, to buttress up the unusual result. Still, the afore quoted emotional passage reveals that, most probably, Purcell's tremendous physical intuition signalled him that something had been overlooked in his study. An accurate treatment (Lazarian \\& Efroimsky 1999) shows that the anelastic dissipation is several orders of magnitude more effective than presumed, and in many physical settings it dominates over the Barnett dissipation. The case of suprathermal dust is one such setting. Without going into redundant details, we would mention that combination of the two dissipation processes provides at least partial alignment of $\\;Z\\;$ and $\\;\\bf \\Omega\\;$ toward $\\;\\bf J\\;$ in Brownian clouds, and it provides a perfect alignment in suprathermal ones.  The presently known mechanisms of grain alignment can be classified into three cathegories: mechanical mechanisms, paramagnetic mechanisms, and via radiative torques. The latter mechanism was addressed in (Dolginov \\& Mytrofanov (1976); Lazarian 1995b; Draine and Weingartner 1996, 1997). It has not yet been well understood. The paramagnetic alignment is due to the Davis-Greenstein (1951) mechanism (initially suggested for Brownian dust particles), and due to the Purcell (1979) mechanism (which is a generalisation of the Davis-Greenstein mechanism, to the suprathermal case). The Davis-Greenstein and Purcell processes operate to bring the granule's rotation axis (which is, as explained above, fully or partially aligned with the granule's major-inertia axis), into parallelism with the magnetic line. This happens because precession of grain's spin axis about $\\;\\bf B\\;$ entails material remagnetisation\\footnote{It is assumed that the grain is either paramegnetic (Davis \\& Greenstein 1951) or ferromagnetic (Spitzer \\& Tukey 1951), (Jones \\& Spitzer 1967). The case of a diamagnetic granule has not had been addressed in the literature so far.} and, therefore, dissipation resulting in a slow removal of the rotation component orthogonal to $\\bf B$. The induced alternating magnetisation $\\bf M$ will lag behind rotating $\\bf B$, giving birth to a nonzero torque equal, in the body frame, to ${\\bf{M}}\\times{\\bf{B}}$. It can be shown (Davis \\& Greenstein 1951) that this torque will entail steady decrease of the orthogonal-to-$\\bf B$ component of the angular velocity\\footnote{A rigorous analysis of the Davis-Greenstein process should be carried out in the language of Fokker-Planck equation (Jones \\& Spitzer 1967).}. The so-called\\footnote{The word ``so-called'' is very much in order here, because the mechanical mechanisms, too, provide alignment relative to the magnetic line, and their very name, ``mechanical'' simply reflects the fact that these effects are not purely magnetic but involve the grains' mechanical interaction with the interstellar wind.} mechanical alignment comprises the Gold (1952) mechanism,  and those of Lazarian (1995 a,b,c,d,e). Lazarian suggested two mechanisms: the cross-over one and the cross-section one, and they show themselves in the case of suprathermal grains only.  The nature of the Gold mechanism is the following. Each collision of the dust particle with an atom or a molecule of the streaming gas adds to the particle's angular momentum a portion perpendicular to the relative velocity. As explained in one of the above footnotes, the major-inertia axis of the body tends to align with the angular momentum. One, hence, may say that the interstellar wind will spin-up the granule so that its maximal-inertia axis will ``prefer'' positions perpendicular to the wind. Since the said major inertia axis is, roughly, the shortest dimension of the rotator, one may deduce that, statistically, the particles tend to rotate with their shortest axes orthogonal to the gas flow. This picture is, though, complified by the precession of the magnetic moments (and of the angular momenta that tend, for the afore mentioned reason, to align with the magnetic moments) about the magnetic line. This mechanism works only for Brownian dust clouds, because it comes into being due to the elastic gas-grain collisions to which only thermal granules are sensitive. To be more exact, it is assumed here that the precession period is much shorter than a typical time during which the grain's angular momentum alters considerably.  The suprathermally-rotating dust particles ignore the random torques caused by the elastic gas-grain collisions, because the timescales for the random torques to alter the spin state are several orders of magnitude larger than the average time between subsequent crossovers (Purcell 1979). Still, the dust granules do become susceptible to the random torques during the brief cross-overs when the granule becomes, for a short time, thermal (i.e., slow spinning). This is the essence of the first Lazarian mechanism of alignment, introduced in Lazarian (1995 d) under the name of ``Cross-Over Mechanism''. Hence, the first Lazarian mechanism is the Gold alignment generalised for suprathermal grains; the generalisation being possible because even suprathermal granules become thermal for small time intervals.  The second Lazarian mechanism, termed by Lazarian (1995 d, e) ``Cross-Section Mechanism'', and studied in Lazarian \\& Efroimsky (1996) and Lazarian, Efroimsky \\& Ozik (1996), is not a generalisation of any previously known effect, but is a totally independent, very subtle phenomenon. Its essence can well be grasped on the intuitive level: a precessing (about the magnetic line) interstellar granule will ``prefer'' to spend more time in a rotational mode of the minimal effective cross section. In other words, the particle has to ``find'' the preferable mean value of its precession-cone's half-angle, value that will minimise the mean cross section. Here, the ``mean cross section'' is the averaged (over rotation, and then over precession) cross section of a granule as seen by an observer looking along the direction of interstellar drift. It is crucial that, though the alignment is due to gas-grain collisions, it establishes itself not relative to the wind direction but relative to the magnetic line about which the spinning grain is precessing.  Now, the goal is to understand how effective this mechanism is for the dust particles of various geometric shapes. Articles (Lazarian and Efroimsky 1996) and (Lazarian, Efroimsky \\& Ozik 1996) addressed the cross-section alignment of oblate and prolate symmetrical grains, correspondingly. In the current paper we intend to extend the study to ellipsoidal granules of arbitrary ratios between the semiaxes. ",
        "conclusions": "In the article thus far, we have investigated the cross-section mechanism of suparthermal-grain alignment in a supersonic interstellar gas stream. While the preceding efforts had been aimed at the cases of oblate and prolate ellipsoidal grains, in the current paper we studied the case of triaxial ellipsoid. We provided a comprehensive semianalytical-seminumerical treatment that reveals the dependence of the alignment measure (Rayleigh reduction factor $R$) upon the semiaxes' ratios and upon the angle $\\Phi$ between the magnetic line and gas drift. We provided a qualitative physical explanation of some of the obtained results.  However, the most intriguing result poses a puzzle and still lacks a simple physical explanation. This is the remarkable shape-independence of the critical value of $\\Phi$, at which $R$ vanishes and the cross-section mechanism fails. For all studied shapes (prolate, oblate, and triaxial with various ratios of semiaxes), this critical value is $\\;\\Phi_o\\;=\\;\\it{arccos}\\,(1/\\sqrt{3})\\;$. We hypothesise that this special nature of the said value of $\\Phi$ is shape-independent.  \\pagebreak    \\appendix"
    },
    "0207/astro-ph0207480_arXiv.txt": {
        "abstract": "We present \\textit{K}-band (1.9 -- 2.5 \\micron) spectra of the Type~Ic SN~2000ew observed with IRCS on the Subaru Telescope.  These data show the first detection of carbon monoxide (CO) emission in a Type~Ic supernova.  The detection of CO in SN~2000ew provides further evidence that molecule formation may be a common occurrence in core-collapse supernova ejecta. The spectrum also contains narrow emission lines of [\\ion{Fe}{2}] and \\ion{He}{1}\\ probably from dense clumps of hydrogen-poor circumstellar gas surrounding SN~2000ew.  Our spectrum of SN~2000ew shows no trace of an unidentified feature seen near 2.26~\\micron, just blueward of the CO emission, in the spectrum of SN~1987A and we discuss proposed detections of this feature in other Type~II supernovae. ",
        "introduction": "The detection of first-overtone carbon monoxide (CO) emission near 2.3~\\micron\\ in SN~1987A opened exciting new possibilities for the study of supernovae (SNe). Unfortunately, for nearly a decade, SN~1987A remained the sole supernova with detected molecular emission. However, with the maturing of near-infrared (NIR) spectrographs, observations of several SNe in the late-time nebular phase have been made.  As a result, CO emission has now been detected in other Type~II supernovae: SN~1995ad \\citep{spyromilio96}, SN~1998S \\citep{gerardy00,fassia01}, SN~1998dl, and SN~1999em \\citep{spyromilio01}.  These observations suggest that molecule formation may be a common occurrence in Type~II SNe.  The study of molecular emission in SNe can provide valuable information about the composition, explosion dynamics, and the late-time temperature evolution and energy balance in the ejecta.  For example, the detection of CO formation can place constraints on the mixing in the SN ejecta.  Because CO is quickly destroyed by the presence of ionized helium, the detection of CO emission in a Type~II or Ib supernova implies either that the CO is not microscopically mixed with helium or that the helium is not ionized \\citep{lepp90,gearhart99}.  In the case of a Type~Ic supernova there is likely little or no helium left in the outer envelope at the time of core collapse.  As a result, CO emission in a Type~Ic does not place strong constraints on the mixing between ejecta layers.  Due to its very large number of collisionally excitable energy levels, CO can be an important coolant of the SN ejecta.  Indeed, at temperatures of a few thousand K, CO emission may be the dominant cooling mechanism \\citep{HSZ89,liu92,spyromilio96,liu95,gearhart99}.  CO cooling may, in turn, play a key role in dust formation in supernova ejecta.  SN~1987A and SN~1998S, two supernovae with detected CO emission, also showed strong evidence of dust formation in the ejecta.  In SN~1987A, thermal emission from dust was detected in the mid and far-infrared while, at the same time, line emission from the ejecta shifted to the blue (\\cite{MC93}, and references therein).  Similarly, late-time spectra of SN~1998S exhibited H and He lines with a multi-peak line profile, the red side of which faded dramatically with time, while the blue side remained nearly constant. In both cases, the changes in line emission were attributed to the formation of dust in the ejecta, which obscured the far side of the supernova causing the redshifted emission to fade.  Thermal emission from dust was also seen in SN~1998S \\citep{gerardy00,fassia01}, but this was not due to dust formation in the ejecta.  In this case, the emission was likely from pre-existing dust in an extended cloud which was heated by X-rays and UV light from a strong interaction between the supernova and dense circumstellar gas \\citep{gerardy02}.  Analysis of near-infrared CO emission can also provide information about the CO-rich ejecta.  For example, the overall shape of the CO emission profile can be used as a temperature diagnostic for the CO emitting gas (e.g.~\\cite{spyromilio88, SH89,liu92}).  The fine structure in the near-infrared is formed by transitions between different vibrational levels, with higher level transitions emitting at longer wavelengths.  As the temperature decreases, the higher vibrational levels become de-populated and emission from the red end of the CO profile decreases.  For high S/N data, it is possible to determine the velocity of the CO emitting gas from the shape of the CO profile \\citep{gerardy00,fassia01}.  At low velocities, the band structure of the CO emission is well defined and the blue edge of the emission profile is quite sharp.  As the expansion velocity increases the bands blend together, smearing out the pattern of peaks and troughs.  Also, the short-wavelength end of the profile creeps farther toward the blue and the rise becomes progressively shallower.  High velocity CO can have important implications for the progenitor star.  To accelerate a significant amount of CO out to high velocity in a Type~II supernova, the progenitor must be massive ($M \\geq 20 M_\\odot$ for $V_{\\rm CO} \\geq 3000$ km~s$^{-1}$) and must have lost a large portion of its H and He rich mantle prior to core-collapse \\citep{gerardy00}.  The large progenitor mass is required to build up a large carbon/oxygen layer which will lie far enough out in the ejecta to accelerate to high velocity.  In a smaller progenitor, the C/O rich layer will be buried too deeply beneath the outer mantle.  On the other hand, for stripped-envelope supernovae, where the pre-collapse mass loss becomes extreme, the requirement for a massive progenitor may be relaxed. For instance, in a Type~Ic supernova there is likely little or no mass left above the C/O rich layer and quite large C/O velocities might be expected from a fairly low mass progenitor. However, if the C/O velocity becomes too high, then the density might drop too fast for a significant amount of CO to form.  Here we present the first detection of CO in a Type~Ic supernova, SN~2000ew. In \\S~2, we describe the observations and data reduction.  In \\S~3 we present the data and compare the CO emission observed in SN~2000ew to that seen in the NIR spectrum of SN~1987A.  We also discuss the conclusion of \\citet{spyromilio01} that both CO emission and an unidentified feature seen in the spectrum of SN~1987A near 2.26~\\micron\\ are ubiquitous features in the NIR spectra of Type~II SNe. ",
        "conclusions": "We have presented \\textit{K}-band spectra of the Type~Ic SN~2000ew which shows strong first-overtone emission of carbon monoxide.  This is the sixth core-collapse supernova for which a CO detection has been published and it is the first detection in a non-Type~II supernova.  The spectrum also shows narrow emission lines of [\\ion{Fe}{2}] and \\ion{He}{1} but is conspicuously lacking narrow \\ion{H}{1} (Br$\\gamma$) emission. We interpret this emission as coming from dense clumps of hydrogen-poor circumstellar gas.  \\citet{spyromilio01} have concluded that CO formation probably occurs in all Type~II supernovae.  We present a table listing ten core-collapse SNe for which \\textit{K}-band spectra were obtained, seven of which show CO emission.  CO emission was typically first observed between 100 and 200~d, and the three SNe without CO detections were not observed past 100~d.  The detection of CO in the Type~Ic SN~2000ew provides the first indication that CO formation may also be common in stripped-envelope core-collapse SNe (Types Ib, Ic, and ``IIb'').  The observed CO emission profile in SN~2000ew is quite similar to that of SN~1987A, and indicates an expansion velocity around 2000 km~s$^{-1}$, which is surprisingly low for a Type~Ic supernova.  The low observed CO velocity may be another indication that strongly non-spherical explosion models are needed for Type~Ic SNe.  To date, there have been no observations that provide strong evidence for high velocity CO emission.  An apparent blue wing in the CO spectrum of SN~1998S presented by GFHW might be due to high velocity CO. However a coeval spectrum taken by \\citet{fassia01} doesn't show this blue extension to the CO profile and they derive a much lower CO velocity from their observations.  Thus the CO velocity in SN~1998S is uncertain. In all of the other CO detections, the observed velocity is around 2000 km~s$^{-1}$ or lower.  It could be that a low C/O velocity is a requirement for CO formation. CO formation is a highly density sensitive process, and it is possible that in higher velocity gas the density decreases too fast to allow much CO to form. However, for most type~II supernovae, the C/O rich layers are buried too deep in the ejecta to expect CO velocities much higher than 2000 km~s$^{-1}$.  Looking for CO emission in type~Ib/c supernovae will provide a better test as much higher C/O velocities are expected in these objects.  We find no evidence in SN~2000ew for the presence of a 2.26~\\micron\\ feature seen in SN~1987A and, contrary to \\citet{spyromilio01}, we suspect it is either faint or absent in other Type~II and Type~Ib/c objects.  SN~1987A remains the only clear detection of this feature, and its identification remains a mystery.  ~\\\\ We would like to thank the Subaru observatory staff for their excellent support,especially Dr.~Hiroshi Terada.  C.~L.~G. and R.~A.~F.'s research is supported by NSF grant 98-76703.  K.~N. has been supported in part by the Grant-in-Aid forScientific Research (07CE2002, 14047206, 14540223) of the Ministry of Education,Culture, Sports, and Technology in Japan.  Research of JCW is supported by NSF Grant 0098644. ~\\\\"
    },
    "0207/astro-ph0207163_arXiv.txt": {
        "abstract": "From equivalent widths of the S\\,{\\sc i} lines at 8694\\,\\AA , Israelian \\& Rebolo (2001) and Takada-Hidai et al. (2002) have derived a surprisingly high sulphur-to-iron ratio ([S/Fe] $\\simeq$ 0.5 to 0.7) in six halo stars with [Fe/H] $\\simeq -2.0$ suggesting perhaps that hypernovae made a significant contribution to the formation of elements in the early Galaxy. To investigate this problem we have used high-resolution spectra obtained with the ESO VLT/UVES spectrograph to determine the S/Fe ratio in 19 main-sequence and subgiant stars ranging in [Fe/H] from $-3.2$ to $-0.7$. The sulphur abundances are determined from S\\,{\\sc i} lines at 8694\\,\\AA\\  $\\em and$ 9212 - 9237\\,\\AA , and the iron abundances from about 20 Fe\\,{\\sc ii} lines. S/Fe ratios as derived from 1D model atmospheres are presented and possible 3D effects are discussed. The initial results from our survey do not confirm the high values of [S/Fe] quoted above; instead we find that the ratio [S/Fe] remains constant at about 0.35 dex for metallicities $-3 < \\feh < -1$. ",
        "introduction": "Sulphur is generally regarded as an $\\alpha$-capture element. The work on Galactic stars by Fran\\c{c}ois (1987, 1988) supported this view by showing that [S/Fe] increases from zero at solar metallicities to a plateau level of about +0.5 dex in the metallicity range $-1.8 < \\feh < -0.8$, an analogous behaviour to that of other $\\alpha$-elements Mg, Si, and Ca, see Ryan et al. (1996). The standard interpretation is that this trend arises from the time delay in the production of 2/3 of the iron by SN of Type Ia relative to the near-instantaneous release of the $\\alpha$-elements by Type II SN.  Recent observations of sulphur in metal-poor stars by Israelian \\& Rebolo (2001) have, however, challenged this view. Their data suggest that \\sfe\\ increases linearly with decreasing \\feh\\ to a level as high as $\\sfe \\sim +0.7$ at $\\feh = -2.0$. The study of Takada-Hidai et al. (2002) based on Keck HIRES observations supports a quasi linear dependence of \\sfe\\ on \\feh\\ although in their case \\sfe\\ reaches only 0.5 dex at $\\feh = -2.0$.  As a possible explanation of the high value of \\sfe\\ in metal-poor stars, Israelian \\& Rebolo propose that massive supernovae with exploding He-cores and a high explosion energy make a significant contribution to the early chemical evolution of galaxies. According to Nakamura et al. (2001) these hypernovae overproduce S with respect to O, Mg and Fe. With this intriguing possibility in mind a more thorough investigation of sulphur abundances in halo stars seems worthwhile. A clarification of the trend of S abundances is also much needed in deciphering the chemical enrichment of damped Ly$\\alpha$ systems (DLAs), widely regarded as the progenitors of present-day galaxies at high redshift. Its importance stems from the fact that, unlike most other heavy elements, S is not depleted onto dust. Consequently, observations of the relatively weak \\SII\\ triplet resonance lines at $\\lambda\\lambda 1250, 1253, 1259$ yield a direct measurement of the abundance of S in DLAs. The only other element for which this is the case is Zn, and indeed most of our current knowledge of the chemical evolution of the universe at high redshift is based on surveys of [Zn/H] in DLAs (e.g. Pettini et al. 1999, Prochaska \\& Wolfe 2002). {\\it If} S is an $\\alpha$-capture element, then its abundance relative to Zn (an iron-peak element) could be used as `a chemical clock' to date the star-formation process at high $z$. Specifically, if the major star formation episodes in the DLAs observed occurred within the last $\\approx 0.5$\\,Gyr, we would expect to measure enhanced [S/Zn] ratios, and {\\it vice versa}. Data on the abundance of S in DLAs have been relatively scarce until recently, but are now becoming available at a progressively faster rate thanks largely to the high sensitivity of the VLT/UVES high resolution spectrograph at blue and near-UV wavelengths. The picture is still confused (Centuri\\'{o}n et al. 2000, Prochaska \\& Wolfe 2002), but one thing is clear: without a secure knowledge of the behaviour of S in metal-poor Galactic stars we stand no hope of interpreting the situation at high $z$.  In this paper we report on a large scale survey of sulphur abundances in metal-poor stars carried out with VLT/UVES. Effects on the derived S/Fe ratio from the modelling of stellar atmospheres are discussed, and preliminary results for about half of the 35 stars observed are presented. ",
        "conclusions": ""
    },
    "0207/astro-ph0207355_arXiv.txt": {
        "abstract": "{PKS~0537--441, a bright $\\gamma$-ray emitting blazar, was observed at radio, optical, UV and X-ray frequencies during various EGRET pointings, often quasi-simultaneously. In 1995 the object was found in an intense emission state at all wavelengths. BeppoSAX observations made in 1998, non-simultaneously with exposures at other frequencies, allow us to characterize precisely the spectral shape of the high energy blazar component, which we attribute to inverse Compton scattering. The optical-to-$\\gamma$-ray spectral energy distributions at the different epochs show that the $\\gamma$-ray luminosity dominates the bolometric output. This, together with the presence of optical and UV line emission, suggests that, besides the synchrotron self-Compton mechanism, the Compton upscattering of photons external to the jet (e.g., in the broad line region) may have a significant role for high energy radiation.  The multiwavelength variability can be reproduced by changes of the plasma bulk Lorentz factor. The spectrum secured by IUE in 1995 appears to be partially absorbed shortward of $\\sim$1700 \\AA.  However, this signature is not detected in the HST spectrum taken during a lower state of the source.  The presence of intervening absorbers is not supported by optical imaging and spectroscopy of the field. ",
        "introduction": "Radiation in the MeV-GeV energy range has been firmly detected by EGRET in 65 active galactic nuclei, all of blazar type (von Montigny et al. 1995;  Hartman 1999; Hartman et al. 1999).  The $\\gamma$-ray emission, coupled to the small emitting volumes inferred by variability time scales, implies beaming in a relativistic jet (which makes the emitting regions transparent to $\\gamma$-ray photons, e.g., McBreen 1979; Maraschi et al.  1992), a distinctive characteristic of blazars.  Among EGRET blazars, many exhibit variability on a range of timescales from years to days (Wehrle et al.  1998;  Mattox et al. 1997; Bloom et al. 1997; Hartman 1996; Hartman 1999).  There is a general consensus that the emission of blazars at energies higher than 10 MeV should be attributed to inverse Compton scattering of relativistic electrons off soft photons produced in or in the vicinity of the jet.  However, the exact role played in the scattering process by photons created in the jet (synchrotron photons) or outside (broad emission line region or accretion disk photons) is not known, nor is the cause of the huge, dramatic $\\gamma$-ray flares observed. Depending on which physical parameters are varying (electron density, magnetic field density, bulk Lorentz factor), variations of different amplitude are expected in the synchrotron and inverse Compton components (see Hartman et al. 1996;  Wehrle et al. 1998; Ghisellini \\& Madau 1996; Ghisellini \\& Maraschi 1996). The most effective way to discriminate among the different scenarios consists in observing these sources simultaneously at several wavelengths spanning a broad range.  PKS~0537--441 ($z = 0.896$) is one of the most luminous and variable blazars at all frequencies, and has been the target of monitoring from radio to X-rays at many epochs (Cruz-Gonzalez \\& Huchra 1984; Maraschi et al. 1985; Tanzi et al. 1986; O'Brien et al. 1988; Bersanelli et al. 1992; Edelson et al. 1992; Falomo et al. 1993a; Treves et al. 1993; Falomo et al. 1994; Romero et al. 1994; Heidt \\& Wagner 1996; Tingay et al. 1996; Sefako et al. 2000; Tornikoski et al. 2001).  The source was detected by EGRET for the first time in 1991 (Michelson et al. 1992; Thompson et al. 1993), and then re-observed at many successive epochs. It is bright and variable in $\\gamma$-rays, and its luminosity in the highest state is comparable to the average luminosity of the strongest and best studied EGRET blazar, 3C~279. Multiwavelength modeling, based on non-simultaneous $\\gamma$-ray and lower frequencies data, has been proposed by Maraschi et al. (1994a) within the synchrotron self-Compton scheme.  On the basis of its radio variability characteristics (Romero et al. 1995) and an off-centered surrounding nebulosity (Stickel et al. 1988), it was proposed that PKS~0537--441 is microlensed by stars in a foreground galaxy. Search for extended optical emission around PKS~0537--441 has a long and controversial history (e.g., Falomo et al. 1992; Lewis \\& Ibata 2000; Scarpa et al. 2000). Clarifying the nature of this emission is of importance both for the study of the properties of the galaxy and environment hosting this very active nucleus and for the alleged possibility of microlensing effects.  In this paper we present multiwavelength observations of PKS~0537--441 at various epochs during the EGRET lifetime, and particularly focus on the $\\gamma$-ray flare of 1995 (Sect. 2.1). We also report on 1998 BeppoSAX observations in the 0.2-50 keV band (Sect. 2.2) and consider archival HST spectra (Sect. 2.3). In Sect. 2.4.2 we compare and discuss the results of all imaging studies of the field. In Sect. 3 we discuss the overall energy distribution and its implications for the nuclear emission mechanisms.   \\begin{table*}[t!] \\caption[]{Multiwavelength observations of PKS~0537--441} \\begin{center} \\begin{tabular}{lccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Date    & Instrument &  $\\alpha_\\nu^a$  & $F_\\nu^b$   & $\\nu^c$    & Ref.$^d$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1991 Feb 11      & ESO 1.5m+CCD            & 1.35$\\pm$0.05 & 2.19$\\pm$0.04 mJy     & $5.45 \\times 10^{14}$ & 1 \\\\ 1991 Apr 16      & ROSAT+PSPC         & 1.1$\\pm$0.4   & 0.79$\\pm$0.05 $\\mu$Jy & $2.41 \\times 10^{17}$  & 2 \\\\ 1991 Jul 26-Aug 08 & CGRO+EGRET            & 1.50$\\pm$0.32 & 39$\\pm$8 pJy         & $9.66 \\times 10^{22}$ & 3,4 \\\\ 1992 May 14-Jun 04 & CGRO+EGRET            & ...           & $< 48^e$ pJy          & $9.66 \\times 10^{22}$  & 5 \\\\ 1992 May 21.41     & IUE+LWP         & ...           & 0.66$\\pm$0.05$^f$ mJy & $1.15 \\times 10^{15}$ & 4 \\\\ 1992 May 21.55     & IUE+LWP         & ...           & 0.78$\\pm$0.03$^f$ mJy & $1.15 \\times 10^{15}$ & 4 \\\\ 1993 Jul 12        & HST+FOS+G130H     & ...           & 0.09$\\pm$0.01$^f$ mJy & $2.14 \\times 10^{15}$ & 4 \\\\ 1993 Sep 16        & HST+FOS+G270H   & $1.92 \\pm 0.09$ & 0.371$\\pm$0.007$^f$ mJy & $1.15 \\times 10^{15}$ & 4 \\\\ 1995 Jan 10-24     & CGRO+EGRET            & 0.96$\\pm$0.18 & 137$\\pm$21 pJy        & $9.66 \\times 10^{22}$ & 6 \\\\ 1995 Jan 30.37   & IUE+SWP        & 1.2$\\pm$0.1$^g$ & 0.90$\\pm$0.02$^f$ mJy & $1.67 \\times 10^{15}$ & 4 \\\\ 1995 Jan 31.44   & IUE+LWP        & 1.2$\\pm$0.1$^g$ & 1.45$\\pm$0.08$^f$ mJy & $1.15 \\times 10^{15}$ & 4 \\\\ 1995 Feb 01.17   & ESO SEST          & ...           & 5.16$\\pm$0.21 Jy      & $90 \\times 10^9$ & 4 \\\\ 1995 Feb 03      & ESO 1.5m+CCD      & 1.29$\\pm$0.04 & 3.03$\\pm$0.15 mJy & $5.45 \\times 10^{14}$ & 7 \\\\ 1995 Feb 05      & ESO 1.5m+CCD       & 1.27$\\pm$0.05 & 3.64$\\pm$0.20 mJy & $5.45 \\times 10^{14}$ & 7 \\\\ 1995 Feb 07      & ESO 1.5m+CCD       & 1.09$\\pm$0.04 & 5.26$\\pm$0.25 mJy & $5.45 \\times 10^{14}$ & 7 \\\\ 1995 Feb 27.14   & ESO SEST                & ...           & 2.88$\\pm$0.23 Jy      & $230 \\times 10^9$ & 4 \\\\ 1998 Nov 28-30   & BeppoSAX      & 0.80$\\pm$0.13 & 0.46$\\pm$0.09 $\\mu$Jy & $2.41 \\times 10^{17}$ & 4 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{6}{l}{$^a$ Spectral index ($F_\\nu \\propto \\nu^{-\\alpha_\\nu}$). Uncertainties, both for $\\alpha$ and fluxes, are at 90\\% confidence level}\\\\ \\multicolumn{6}{l}{ ~~ for the X- and $\\gamma$-ray measurements, and at 68\\% for UV, optical and millimetric. No spectral fit has been}\\\\ \\multicolumn{6}{l}{ ~~ tried when the data signal-to-noise ratio was too low or the spectral range too limited.}\\\\ \\multicolumn{6}{l}{$^b$ Flux density. EGRET flux conversion follows Thompson et al. (1996). Optical-to-X-ray data are corrected}\\\\ \\multicolumn{6}{l}{ ~~ for Galactic extinction.}\\\\ \\multicolumn{6}{l}{$^c$ Frequency to which the flux density refers, in Hz.}\\\\ \\multicolumn{6}{l}{$^d$ References. {\\bf 1:} Falomo et al. 1994; {\\bf 2:} Treves et al. 1993; {\\bf 3:} Thompson et al. 1993;}\\\\ \\multicolumn{6}{l}{ ~~ {\\bf 4:} This paper; {\\bf 5:} Hartman et al. 1999; {\\bf 6:} Mukherjee et al. 1997; {\\bf 7:} Scarpa \\& Falomo 1997.}\\\\ \\multicolumn{6}{l}{$^e$ 2-$\\sigma$ upper limit.}\\\\ \\multicolumn{6}{l}{$^f$ The uncertainty is only statistical. For IUE data, this was evaluated following Falomo et al. (1993b).}\\\\ \\multicolumn{6}{l}{$^g$ From a spectral fit over the band 1700-5500 \\AA.}\\\\ \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{}  In Fig. 2 we have reported the broad-band spectral energy distributions of PKS~0537--441 constructed with quasi-simultaneous data from the millimetric to the $\\gamma$-ray frequencies (see Table 1).  In addition, we have reported the unpublished BeppoSAX spectrum.  The multiwavelength state of PKS~0537--441 in January-February 1995 was one of the brightest recorded for this source during the lifetime of EGRET (cf. Tornikoski et al.  1996; Falomo et al. 1994; Hartman et al. 1999). The comparison of $\\gamma$-ray detection and optical flux of 1991 and the EGRET upper limit and UV flux of 1992 is suggestive of little or no multiwavelength variability between the two epochs.  The increase of the optical-to-UV flux from the 1991-1992 to the 1995 level corresponds to a variation of similar amplitude (a factor of $\\sim$2) in the $\\gamma$-rays, or only slightly larger. The flat BeppoSAX spectrum suggests that a single emission component dominates in the energy band 0.1-30 keV.  PKS~0537--441 exhibits the typical double-humped multiwavelength spectral shape of ``Low-frequency-peaked\" blazars (Padovani \\& Giommi 1995; Sambruna et al. 1996; Fossati et al. 1998): the first component peaks at wavelengths longer than the optical (likely in the far-infrared, as suggested by the IRAS data, taken at a much earlier epoch, Impey \\& Neugebauer 1988) and is due to synchrotron radiation.  The second component, which peaks around $\\sim 10^{22}-10^{24}$ Hz (Fig. 2) and is a factor $\\sim$6 more powerful than the synchrotron maximum, is probably produced via inverse Compton scattering between relativistic electrons and synchrotron photons (Maraschi et al. 1992; Maraschi et al. 1994b; Bloom \\& Marscher 1993) or external photons (broad emission line region or accretion disk, Dermer \\& Schlickeiser 1993; Sikora et al. 1994).  The overall energy distribution of PKS~0537--441, the inverse-Compton dominance, and the presence of optical and UV emission lines suggest that external Compton upscattering may be significant with respect to the synchrotron self-Compton process in producing the high energy emission (Ghisellini et al. 1998; Ghisellini 2001). To model the broad-band energy distribution of PKS~0537--441 we have assumed that the emission is produced in a region filled by relativistic particles which radiate at low energies via synchrotron, and upscatter both synchrotron photons and accretion disk photons reprocessed in the broad emission line region. We have not included the direct contribution of the accretion disk to the population of seed photons to be upscattered by the jet electrons, because those would be highly redshifted in the frame of the blob, which is emitting at a distance of $\\sim$ 0.3 pc (see caption to Fig. 2) from the jet base (Sikora et al. 1994). The size of the emitting region has been constrained with the variability time scale in $\\gamma$-rays ($\\sim$2 days); the electron energy distribution is modeled with a double power-law (see Tavecchio et al. 2000 and Ballo et al. 2002 for more details of the model). The low state model has been constrained with the BeppoSAX spectrum and with the 1991-1992 optical and $\\gamma$-ray data, while the high state model reproduces the 1995 multiwavelength data.  Note that the ROSAT spectral measurement of April 1991, which is affected by a large uncertainty, is consistent both with the low and high state model curves. The IRAS data are marginally consistent with the high state model.  The difference between the two model curves (reported in Fig. 2; see caption for the model parameters) is solely determined by a change in the Lorentz factor of the relativistic plasma bulk motion (increasing by 10\\% from the low to the high state) and in the index of the upper branch of the electron energy distribution ($n_2$ in the notation of Ballo et al. 2002, slightly flatter in the brighter state, see caption to Fig. 2).  Under our assumption that the high energy component results from the contribution of both synchrotron-self Compton and external Compton mechanisms, the variation of the bulk Lorentz factor would produce a correlation between the variability amplitude of the synchrotron and inverse Compton components which is intermediate between linear and half-cubic (Ghisellini \\& Maraschi 1996). This is consistent with the observations, although the difficulty of exactly locating the peaks of the emission components, the non-strict simultaneity of the optical-to-UV spectra and the rapid $\\gamma$-ray flare seen at the end of the EGRET viewing period make the fit results only indicative.  From our measured UV line intensities reported in Table~2 and from dereddened MgII emission line intensity (Sect. 2.4.1), assuming $H_0$ = 65 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, $\\Omega_\\Lambda = 0.7$, we derive a total line luminosity $L_{\\rm BLR} \\sim 4.7 \\times 10^{44}$ erg s$^{-1}$ (the other observed optical emission lines are sufficiently weak that their contribution to the line luminosity is not significant, see Lewis \\& Ibata 2000). Using this and the external photon density assumed in our model, $U_{\\rm ext} = 6 \\times 10^{-3}$ erg cm$^{-3}$, we can evaluate the size of the broad line region, $R_{\\rm BLR} \\simeq 7.9 \\times 10^{17}$ cm.  As an independent check, we have considered the empirical relationship determined by Kaspi et al. (2000) between the size of the broad line region and the luminosity of the thermal continuum at 5100 \\AA\\ in quasars.  Assuming a covering factor of the broad line clouds of 10\\%, we can estimate the disk luminosity to be $L_{\\rm disk} \\sim 4.7 \\times 10^{45}$ erg s$^{-1}$.  This value is consistent with the approximate upper limit which can be derived from the broad band spectrum, $\\sim 2 \\times 10^{46}$ erg s$^{-1}$.  After accounting for a factor of $\\sim$3 difference between a bolometric and a monochromatic disk output, we find that our disk luminosity estimate implies, according to Kaspi et al.'s formula, $R_{\\rm BLR} \\simeq 5.8 \\times 10^{17}$ cm, well consistent, given the uncertainties, with the size we have determined based on the observed $L_{\\rm BLR}$ and on the assumed photon density. We finally note that our observed $L_{\\rm BLR}$ is about a factor 4 lower than that estimated for this source by Celotti et al. (1997), based on an old measurement of the MgII emission only.  The validity and origin of the spectral dip seen in the IUE spectrum of January 1995 at wavelengths shorter than 1700 \\AA\\ remain to be established. Assuming it to be real, it could be qualitatively described by an absorption edge or broad trough, and may be identified with neutral hydrogen ionization discontinuity (Lyman limit)  at the redshift of the source or with Ly$\\alpha$ forest blanketing up to a maximum redshift of $\\sim$0.4.  Continuum decrements shortward of Ly$\\alpha$ emission, Lyman limit systems and damped Ly$\\alpha$ absorption features have been reported in few blazars at redshifts $0.5 \\simlt z \\simlt 1.5$ (e.g., PKS~0637--752, Cristiani et al. 1993; PKS~0735+178, PKS~2223-052 Courvoisier \\& Paltani 1992, Lanzetta et al. 1995).  Absorption features have recently been reported also in X-ray blazar spectra (Tavecchio et al. 2000). The host galaxy or the halos of the galaxies located in the vicinity of PKS~0537--441 and at similar redshift (projected distances of $\\sim$30 kpc, Heidt et al. 2002) may be responsible for absorption in the bluer part of the UV nuclear spectrum. Using the photoelectric cross-section computed by Rumph et al. (1994) to model the interstellar opacity at extreme UV wavelengths, we estimated that in the case of Lyman continuum absorption the equivalent intrinsic $N_{\\rm HI}$ would be about 3 orders of magnitude lower than the Galactic $N_{\\rm HI}$ in the direction of the blazar, and therefore its effect on the X-ray spectrum would be undetectable. Alternatively, the FOS G130H spectrum may suggest resemblance with those of Broad Absorption Line quasars (e.g., Hamann \\& Ferland 1999; Arav et al. 2001), so that, in the lower resolution IUE SWP spectrum, many broad features would appear blended in a unique edge.  Higher signal-to-noise ratio data (ideally HST STIS spectra simultaneously covering the wavelength range 1100-2700 \\AA) would be necessary to confirm the presence of the feature and accurately model it. The nuclear emission and environmental characteristics of the blazar PKS~0537--441 make it a case study for future $\\gamma$-ray missions like AGILE and GLAST and for the NGST, respectively."
    },
    "0207/gr-qc0207062_arXiv.txt": {
        "abstract": "We consider the problem of searching for gravitational waves emitted during the inspiral phase of binary systems when the orbital plane precesses due to relativistic spin-orbit coupling. Such effect takes place when the spins of the binary members are misaligned with respect to the orbital angular momentum. As a first step we assess the importance of precession specifically for the first-generation of LIGO detectors. We investigate the extent of the signal-to-noise ratio reduction and, hence, detection rate that occurs when precession effects are not accounted for in the template waveforms. We restrict our analysis to binary systems that undergo the so-called simple precession and have a total mass $\\lesssim 10\\,M_{\\odot}$.  We find that for binary systems with rather high mass ratios (e.g., a 1.4\\,M$_\\odot$ neutron star and a 10 M$_{\\odot}$ black hole) the detection rate can decrease by almost an order of magnitude. Current astrophysical estimates of the rate of binary inspiral events suggest that LIGO could detect at most a few events per year, and therefore the reduction of the detection rate even by a factor of a few is critical. In the second part of our analysis, we examine whether the effect of precession could be included in the templates by capturing the main features of the phase modulation through a small number of extra parameters. Specifically we examine and tested for the first time the 3-parameter family suggested in \\cite{Apost96}. We find that, even though these ``mimic'' templates improve the detection rate, they are still inadequate in recovering the signal-to-noise ratio at the desired level. We conclude that a more complex template family is needed in the near future, still maintaining the number of additional parameters as small as possible in order to reduce the computational costs. ",
        "introduction": "With the international network of ground-based gravitational wave (GW) detectors LIGO \\cite{Abram92}, VIRGO \\cite{Caron97}, GEO600 \\cite{Danzm95}, TAMA \\cite{Tagos01} coming online, the need for accurate source modeling, which can guide the construction of optimal data analysis strategies and therefore maximizes the detection efficiency is becoming increasingly pressing. The expected signals will be so weak compared to the intrinsic noise of the detectors, that one relies on a number of different data-processing techniques for signal extraction and detection, followed by search for coincidences in two or more instruments. Matched-filtering provides the optimal linear search technique \\cite{Helst68,OwenS99}, particularly suitable to search for signals that are characterized by a large number of wave cycles within the interferometer observational window ($\\simeq 40\\,{\\rm Hz} - 2$\\,kHz), and whose waveforms can be accurately modeled from a theoretical point of view.  A well-known example of such a signal is the binary inspiral of compact objects, either neutron stars (NS) or stellar-mass black holes (BH).  Such inspiral events are the most promising candidates for ground-based interferometers, and, for the instruments of first generation, especially those with high total mass (BH/NS and/or BH/BH binary systems), see \\cite{CutleT02} and references therein.  A crucial component in the successful implementation of matched-filtering is the construction of a reliable family of gravitational waveforms to be used as ``templates''. As one might expect intuitively, the closer the templates are to the ``true'' signal, the higher the signal-to-noise ratio (SNR) at which a detection can be achieved. For this reason in recent years a tremendous effort has been devoted to the computation of the waveforms that characterize the inspiral phase, using different approaches, such as post-Newtonian expansions (see \\cite{Blanchet02} and references therein for an extensive review) and Pad\\'e approximants \\cite{DamouIS98,DamouIS00}.  The simplest family of templates corresponds to GW signals from two point-masses orbiting each other.  However, it is possible that the true inspiral is characterized by other more complicated effects, such as relativistic spin-orbit and spin-spin couplings due to the presence of high spins misaligned one with respect to the other and to the orbital angular momentum. One of the most dramatic effects induced by such coupling is the precession of the binary orbital plane around a fixed direction, with a significant number of precession cycles within the GW frequency band of interest. As a result of this precession the polarization of the waves impinging on the detector changes during the observation, which in turn produces an amplitude and phase modulation of the signal at the detector output \\cite{ApostCST94}. It is well known that if such modulation is not accounted for in the search templates the detection efficiency could decrease significantly, especially for binaries containing BHs and/or characterized by large mass ratios \\cite{Apost95,Apost96}. This would lead to a reduction of the volume of the Universe accessible by any given search and hence of the rate of detection. Current predictions of the detection rates for binary inspiral with the first-generation of laser interferometers are relatively low, in the most optimistic cases up to just a few events per year even for massive binaries with two black holes \\cite{KalogB01,BelczKB02}. Such estimates are obtained by assuming perfect match between the signal and, at least, one of the templates of the filter bank. If modulations induced by spin-orbit precession are not included into the family of templates, the actual signal-to-noise ratio that can be achieved in any given search is lower. If such a loss of SNR occurs for a significant portion of the parameter space that characterize binary systems, the chance of detections by LIGO and other ground-based interferometers would be severely compromised. This would in turn require to construct a new family of template that takes precession effects into account. In the case of binaries with two neutron stars, both observations of pulsar spin periods \\cite{TayloW82,Wolsz91} and our current theoretical understanding of neutron star tidal evolution \\cite{BildsC92,Kocha92} suggest that the spins are essentially negligible in neutron star binary systems. Therefore, we concentrate on binaries with either a black hole and a neutron star or two low-mass black holes.  The PN equations of motion, up to second post-Newtonian order, including spin-orbit and spin-spin corrections, have been derived in \\cite{Kidde95}. However, in this paper, instead of integrating the complete set of equations, we use an analytical approximation of the 1.5PN order equations derived by Apostolatos {\\em et al.} \\cite{ApostCST94,Apost95}. This approximation is valid in the regime of {\\em simple precession}, which is relevant in the following two cases~:  (i) two spinning objects of equal mass $\\l(m_1=m_2\\r)$ or (ii) unequal-mass objects but with only the most massive one spinning $\\l(S_2=0\\r)$. Here we restrict our computations to systems for which only the most massive object is spinning. In the case of binaries with a black hole and a neutron star, such an assumption is physically well justified based on the same arguments relevant for binary neutron stars. In the case of binary black holes, the most massive of the two is also expected to be spinning more rapidly, since it is the one that in most cases is formed first, and therefore may have been spun up through accretion from its non-degenerate companion, the progenitor of the second black hole \\cite{BelczKB02}. We restrict ourselves to systems with relatively moderate total mass (i.e.  $m_1+m_2 \\lesssim 10 M_\\odot$) so that one expects the PN-expansions to be fairly accurate \\cite{BradyCT98}. The detectability of two, high mass, non-spinning black holes, for which the PN-expansions break down, has been recently assessed in \\cite{BuonaCV02}. Then, we consider whether the inspiral detection efficiency could be improved by using a new family of ``mimic'' templates, initially suggested by Apostolatos in \\cite{Apost96} as a possible solution for the poor fitting factor achieved by non precessing templates. This new class of templates depends only on a small number of additional parameters -- three in this case -- but attempt to ``mimic'' the effect of precession. Apostolatos presented heuristic arguments supporting the choice of waveform: here we quantitatively explore for the first time whether they indeed recover most of the signal to noise ratio which is lost due the modulation effects. The result of this investigation is actually that this new family of templates is still inadequate to search for precessing binaries over essentially the entire parameter space.  This paper is organized as follows: Sec. \\ref{s:signal} presents briefly the simple precession regime and reviews the various types of templates used throughout this paper, (post)-Newtonian and the ``mimic'' templates which include a correction term to the phase. The concept of the fitting factor is also introduced. Results are shown in Sec. \\ref{s:results}. We draw conclusions and discuss lines for future work in Sec.  \\ref{s:conclu}. ",
        "conclusions": "\\label{s:conclu} In this paper we present a systematic and quantitative study of the effect of precession on the detection of binary inspiral signals.  Unlike earlier work \\cite{Apost96}, we focus on (i) implications for the initial LIGO observations, (ii) the signal-to-noise ratio reduction due to the effects of precession alone -- by analyzing signals and templates computed at the same PN order for the non-precessing portion of the phase -- and (iii) the quantitative study of the fitting factor of a family of mimic templates that have been suggested for recovering high SNR. We have examined results for 4 pairs of masses for the full range of spin properties and directions, and discussed in more detail the results of one of these pairs ($\\l(10;1.4\\r)$ solar masses) that exhibits the strongest precession effects and is appropriate for binaries with a black hole and a neutron star.  We first addressed the question of how important is precession and for what binary properties. We found that it can seriously affect detection, if mass ratios are in excess of about 2, spin magnitudes in excess of about 30\\% of maximum, and spin tilts in excess of about 35 degrees. These results help in the identification of the parameter space where ways of improving the detection efficiency must be found. We found that the detection rate can decrease by almost an order of magnitude if searches are performed with templates that do not include precession effects. Such a loss of events can be very concerning, given the current low-estimates for the expected detection rates \\cite{BelczKB02}.  Precession waveforms depend on a large number of parameters and their use as a template family is not feasible. Therefore, it is important to introduce a family of templates that can ``mimic'' precession effects well enough, but they depend on a small number of parameters. Here we tested one such family suggested in \\cite{Apost96}, and found that, although they do increase the detection rate, this increase is not significant enough to raise the detection rate to desirable levels (leading to improvement of more than a factor of a few for the case of strong precession). We conclude that other forms of ``mimic'' templates must be explored in the near future. Such an exploration is beyond the scope of the present paper, but we are undertaking it as the next step in this project.  One can further restrict the physical parameter space over which precession is crucial for detection, by convolving our results from the first part of the paper with astrophysically relevant distributions of binary and spin parameters of double compact objects, derived based on our current understanding of the formation of compact binaries with black holes. We are currently working on such a convolution with detailed population calculations for many formation models \\cite{BelczKB02} of BH binaries, with the goal of obtaining an astrophysically-motivated picture of how important it will be to include precession in the template families for the search of GW inspiral signals in the next few years \\cite{IhmKGB02}."
    },
    "0207/astro-ph0207025_arXiv.txt": {
        "abstract": "{ We have carried out a survey of the Andromeda galaxy for unresolved microlensing (pixel lensing). We present a subset of four short timescale, high signal-to-noise microlensing candidates found by imposing severe selection criteria: the source flux variation exceeds the flux of an $R=21$ magnitude star and the full width at half maximum timescale is less than 25 days. Remarkably, in three out of four cases, we have been able to measure or strongly constrain the Einstein crossing time of the event. One event, which lies projected on the M31 bulge, is almost certainly due to a stellar lens in the bulge of M31. The other three candidates can be explained either by stars in M31 and M32 or by MACHOs.} ",
        "introduction": "}  The galactic dark matter may be partly composed of compact objects (e.g., faint stars, brown dwarfs, Jupiters) that reside in halos and are popularly called MACHOs (``MAssive Compact Halo Objects''). Microlensing surveys towards M31 \\cite{crotts92,baillon93} have the potential to resolve the puzzling question raised by searches toward the Magellanic Clouds: the optical depth $\\tau\\sim 10^{-7}$ measured by MACHO \\cite{macho} is too large by a factor 5 to be accounted for by known populations of stars and too small by the same factor to account for the dark matter, while the mass scale inferred for the lenses \\mbox{$M\\sim 0.4\\,M_\\odot$} is in the mid-range of normal stars. EROS \\cite{eros} obtained upper limits that are consistent with the MACHO results.  Since M31 is 15 times more distant than the Magellanic Clouds, the stars are about 200 times fainter and more densely packed on the sky. Even with new techniques that are required to monitor flux changes of unresolved stars in the face of seeing variations \\cite{cro96,ans97,ans99}, the low signal-to-noise ratio (S/N) engenders a whole range of problems. First, the detection efficiency is reduced. Second, there is a degeneracy between the Einstein crossing time, the impact parameter and the source flux \\cite{gould96}. Third, some variable stars can not be easily distinguished from microlensing events and so will contaminate the signal. We elaborate on each of these points as follows:\\\\ i) The loss of detection efficiency is severe because a microlensing event can be rejected by the selection procedure if the source star or neighbouring blended stars are variable. Indeed, if it is to be detected as microlensing, an event must rise above the photon noise due to all the blended neighbouring stars. For a fixed impact parameter, the brighter the source star, the easier it is to detect the event. So, bright sources are the most likely microlensing candidates. Unfortunately, the Hipparcos catalogue shows that most of the bright sources with $M_V <0$ are prone to intrinsic variability \\cite{perry}.\\\\ ii) The degeneracy between parameters of the lightcurves occurs mainly around the time of maximum magnification and becomes more severe as the impact parameter becomes smaller. It can be partly broken for events with good S/N and good sampling on the wings -- as for three of the four events presented later.\\\\ iii) To distinguish between any MACHO population and variable stars, we intend to exploit the fact that M31 is highly inclined ($i \\sim 77^\\circ$) to our line of sight. Therefore, if MACHOS are distributed in a roughly  spherical halo, the density of MACHOs along the line of sight is larger on the far side of the M31 disk than on the near side. This implies a larger optical depth and an excess of microlensing events on the far side \\cite{crotts92,ker01}.  The POINT-AGAPE collaboration is carrying out a pixel-lensing survey of M31 using the Wide Field Camera (WFC) on the \\mbox{$2.5\\,$m} Isaac Newton Telescope (INT).  We monitor two fields, each of \\mbox{$\\sim 0.3\\,$deg$^2$}, located North and South of the M31 centre.  After a brief description of the observations and data analysis in \\mbox{Section \\ref{sec:obsdata}}, we present four events with high S/N and short durations in \\mbox{Section \\ref{sec:candidates}}, for which microlensing is by far the most plausible interpretation. ",
        "conclusions": "\\label{sec:discuss} We have reported first results extracted from two years of data in our pixel lensing survey of the Andromeda galaxy.  By imposing the stringent requirements $R(\\Delta F)<21$ and $t_{1/2}<25\\,$days, we have selected four very convincing microlensing candidates with high signal-to-noise ratio and short timescales. For three of our four events we have been able to make reliable determinations of the Einstein crossing time, which provides additional clues as to the probable origin of these events. In the case of PA-00-S3, the event is most likely caused by a stellar lens in the M31 bulge. In the three other cases, MACHOs and stellar lensing are plausible.  \\bigskip \\noindent {\\bf Acknowledgments}: YLD was supported by a PPARC postdoctoral fellowship and SJS by a PPARC advanced fellowship. NWE acknowledges help from the Royal Society.  Work by AG was supported in part by a grant from the Centre National de la Recherche Scientifique and in part by grant AST 02-01266 from the NSF.  \\begin{figure} \\resizebox{\\hsize}{!}{\\epsfig{file=INTevts.eps,scale=0.5,clip=}} \\caption{Positions of the four microlensing candidates projected on M31 (\\textsf{http://aladin.u-strasbg.fr}, POSSII). The dotted lines show the boundaries of observed field and the white cross indicates the position of the centre of M31. Note that S4 lies just next M32.} \\label{fig:positions} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\epsfig{file=n1db.eps,scale=0.4,clip=} } \\caption{Top panel: $r'$-lightcurve of the PA-99-N1 microlensing candidate between August 1999 and January 2001. Encircled variations show three bumps, the first one being the microlensing candidate. As in Figure \\ref{fig:lightcurves}, the solid line shows the best-fit Paczy\\'nski curve (data points for the secondary bumps being masked for this fit). Bottom panels: image differencing (at left around the maximum magnification of the microlensing event, at right on data points belonging to secondary bumps) showing that the microlensing event and the secondary bumps are separated by $\\sim 3\\,$pixels ($\\sim 1''$). On the bottom right panel, a black square shows the superpixel centred on the microlensing candidate.} \\label{fig:n1doublebump} \\end{center} \\end{figure}"
    },
    "0207/astro-ph0207213_arXiv.txt": {
        "abstract": "We use a compilation of cosmic microwave anisotropy data (including the recent VSA, CBI and Archeops results), supplemented with an additional constraint on the expansion rate, to directly constrain the parameters of slow-roll inflation models. We find good agreement with other papers concerning the cosmological parameters, and display constraints on the power spectrum amplitude from inflation and the first two slow-roll parameters, finding in particular that $\\epsilon_1 < 0.057$. The technique we use for parametrizing inflationary spectra may become essential once the data quality improves significantly. ",
        "introduction": "Recent measurements of the cosmic microwave background (CMB) show a flat portion at low multipole number $\\ell$ and a sharp peak around $\\ell\\sim 200$, as well as tentative evidence for a peak structure beyond $\\ell = 200$. This represents a tremendous success for the simplest models of the universe described by a flat Friedmann--Robertson--Walker metric with adiabatic perturbations, which are in excellent qualitative agreement with these observations. The power of the CMB is that it can be used to constrain cosmological parameters, as well as allowing us to test our assumptions about the form of the initial irregularities in qualitative and now quantitative ways. The most popular assumption concerning the initial irregularities is that they originated during a period of cosmological inflation (see Liddle \\& Lyth 2000 for an extensive account of inflationary cosmology).  There have now been several papers which have searched for possible effects in this data from quite complicated inflationary models. One example is the inclusion of extra isocurvature degrees of freedom in the primordial power spectrum in addition to a dominant adiabatic component (Trotta, Riazuelo \\& Durrer 2001; Amendola et al.~2002), with the conclusion that the current data set is consistent with a sub-dominant isocurvature component (or even a dominant one on large scales in the case where the isocurvature perturbations are correlated with the adiabatic ones) and that the allowed values and ranges of the cosmological parameters are sensitive to the type of perturbations under consideration. Another example is attempts to fit inflation-motivated `features in the power spectrum' to the data (Griffiths, Silk \\& Zaroubi 2001; Barriga et al.~2001; Adams, Cresswell \\& Easther 2001). These scalar power spectra have the intrinsic property of introducing extra degrees of freedom --- the shape parameters associated with the feature --- which can be used alter the peak heights at will. However it is necessary to choose the features to coincide with characteristic scales in the CMB power spectrum, such as the first or second peaks [early work in this direction (Adams, Ross \\& Sarkar 1997) was also motivated by possible features in the matter power spectrum]. The CMB spectrum alone offers no evidence for any extra features and so smooth power spectra are currently best motivated.  It is surprising that relatively little attention has been paid to the CMB spectrum resulting from slow-roll inflation, even though these models have been the most intensively studied since its conception. Slow-roll inflation has acted as a guiding principle for inflation model builders, and is the simplest assumption and capable of giving excellent agreement with observations. The reason why specific studies of slow-roll inflation have been lacking is the usual assumption that inflation predicts a nearly power-law shaped spectrum, and hence that any information about inflation can be extracted as some linear combination of the constraints on the two key parameters, the scalar spectral index $n_{{\\rm S}}$, and the tensor fraction $R$, an approach used by Kinney, Melchiorri \\& Riotto (2001) and by Hannestad et al.~(2002) to discuss constraints on inflation. Hansen \\& Kunz (2001) also included the running of the spectral index, translating constraints on these parameters to place bounds on derivatives of the inflaton potential. Other parameter analyses (Wang, Tegmark \\& Zaldarriaga 2002; Percival et al.~2002) have tended to focus on results for other cosmological parameters such as the densities of the various matter components, and have been content to use this simple parametrization.  However, this approach ignores the fact that current data are only weakly constraining, and the current data set permits parameter regions where a significant deviation from a Harrison--Zel'dovich spectrum is allowed, and where the use of the full slow-roll power spectra is required to obtain robust results. The principal aim of this paper is to make the first direct estimation of slow-roll inflation parameters from CMB data. While the full predictions are relevant only in extreme regions of parameter space given current data, as the global cosmological data set improves (and in particular as significant observational weight develops on the high $\\ell$ part of the CMB spectrum) it may well be that these types of corrections take on increasing importance, depending which (if any) inflation models prove capable of fitting the data.  We do not expect any dramatic new results from this analysis given the quality of present data, though it is a useful test of the robustness of results under more general forms of the initial power spectra. However it is an important test of principle to bring these methods to bear on cosmological parameter estimation, as the high-quality data of coming years may well require the high-accuracy description of the power spectrum that this approach allows. ",
        "conclusions": "We have implemented the detailed second-order predictions for the inflationary power spectra, given by equations (\\ref{bs0}) to (\\ref{eqn:psnorm}), into a CMB parameter search method, using the logarithmic expansion of ${\\mathcal P}(k)$, Eq.~(\\ref{eqn:chaotic}), for the first time. Although the present data set cannot hope to actually measure the weak running of the spectral index induced by a significant tensor component, we derived sensible limits on the power spectrum amplitude ${\\mathcal P}_{{\\mathcal R}}(k=0.05\\,{\\rm Mpc}^{-1})$ and the first two parameters, $\\epsilon_1$ and $\\epsilon_2$, by assuming the running to be weak, which was achieved by considering models with $\\epsilon_1<0.07$, $-0.4<\\epsilon_2<0.3$ (and fixing $\\epsilon_3=0$). We also derived a sensible limit for $\\omega_{{\\rm b}}$ which acts as a useful consistency check on the assumption of adiabatic perturbations with an approximately power-law form.  While the results of the present paper do not add much to existing studies parametrizing the spectra as power-laws, our paper represents an important point of principle in implementing precise slow-roll inflation predictions for the first time. As the global dataset improves, including MAP and then {\\it Planck} data, it is quite likely that these techniques are required to ensure robust estimation even of cosmological parameters such as the densities of the various components. Further, these techniques will be essential to squeeze the maximum possible amount of information out of the data regarding inflation, should slow-roll inflation continue to give the simplest viable interpretation of observational data. As the data-set improves, it will be interesting to open up the $\\epsilon_3$ direction as well as exploring the possibility of a negligible tensor prior, in order to differentiate between inflationary models as effectively as possible. An interesting goal would be to determine the signs of both $\\epsilon_2$ and $\\epsilon_3$, which would allow us to immediately rule out three quarters of all single-field slow-roll inflation models."
    },
    "0207/astro-ph0207207_arXiv.txt": {
        "abstract": "We present a compilation of radio, infrared, optical and hard X-ray (2-10kev) data for a sample of 90 Seyfert 2 galaxies(Sy2s) with spectropolarimetric observations (41 Sy2s with detection of polarized broad lines  (PBL) and 49 without PBL). Compared to Sy2s without PBL, Sy2s with PBL tend to be earlier-type spirals, and show warmer mid-infrared color and significant excess of emissions (including the hard X-ray(2-10kev), [O {\\sc iii}]$\\lambda$5007, infrared (25 $\\mu$m) and radio). Our analyses indicate that the majority of Sy2s without PBL are those sources having less powerful AGN activities, most likely caused by low accretion rate. It implies that the detectability of the polarized broad emission lines in Sy2s may depend on their central AGN activities in most cases. Based on the available data, we find no compelling evidence for the presence of two types of Sy2s, one type of them has been proposed to be intrinsically different from Sy2s claimed in Unification Model. ",
        "introduction": "In the scheme of the standard unification model, Seyfert 1 and 2 galaxies (Sy1s and Sy2s hereafter) are intrinsically the same objects and the absence of broad emission lines in Sy2s is ascribed to the obscuration by a pc-scale dusty torus oriented along the line of sight (see the reviews by  Antonucci 1993). The observational evidence for this model includes the detection of polarized broad emission lines in some Seyfert 2 galaxies (Antonucci \\& Miller 1985; Tran 1995 and 2001; Young et al. 1996; Heisler et al. 1997; and Moran et al. 2000), the detection of broad lines in the infrared spectra of some Sy2s (Rix et al. 1990; Ruiz et al. 1994; and Veilleux et al. 1997) and the detection of a prominent photoelectric cutoff in the X-ray spectra of Sy2s indicating the presence of large columns of gas along the line of sight (Koyama et al. 1989; Awaki et al. 1991; Maiolino et al. 1998; Risaliti et al. 1999). However, recent investigations suggest that this strictest version of unification model needs modifications. Among them, we may find, for example, the outflowing wind model (Elvis 2000), and the existence of two intrinsically different populations of Sy2s $-$ the hidden Sy1 and the \"real\" Sy2 with a weak or absent Sy1 nucleus $-$ ( Tran 2001) based on a spectropolarimetric survey of the CfA  and 12$\\mu$m samples of Seyfert 2 galaxies. Nevertheless, Antonucci (2001) strongly argued that the evidence claimed by Tran is quite uncertain.  With the improvement of the techniques and instruments for the spectropolarimetry, now people have observed a large sample of Sy2s with less bias (e.g.  Heisler et al. 1997; Moran et al. 2000). For the present,  polarized broad lines (PBL) have been detected in several dozens of  Sy2s, while not detected in other several dozens. Although such surveys are probably biased inherently since pre-selection was done according to the broad-band polarization, they still provide a largest sample for more meaningful analysis than before so that we may, or may not, find some compelling evidence for the proposed modifications, especially the presence of two types of Sy2s. That is what we would present in this paper.  In this work, we provide the multiwavelength data for a sample of 90 Seyfert 2 galaxies collected from recent literatures in \\S2, and compare the properties of host galaxies, infrared, radio and hard X-ray continua and [O {\\sc iii}] emission in \\S3. The implications and discussions on the results are given in \\S4, and conclusions in \\S5 . ",
        "conclusions": "In this paper, we collect radio, infrared, optical and hard X-ray data for a sample of 90 Seyfert 2 galaxies with spectropolarimetric observations. Out of these 90 objects, 41 show polarized broad lines (most likely ascribed to scattering of the broad line region) and 49 do not. Compared to Sy2s without PBL, Sy2s with PBL tend to be earlier-type spirals, and show warmer mid-infrared color and significant excess of emissions (including the hard X-ray(2-10kev), [O {\\sc iii}]$\\lambda$5007, infrared (25 $\\mu$m) and radio), while their distributions of blue luminosity and absorbing column density are similar. Our analyses suggest that the majority of Sy2s without PBL are those sources having less powerful AGN activities, most likely caused by low accretion rate. It implies that the detectability of the polarized broad emission lines in Sy2s may depend on their central AGN activities in most cases. Based on the available data, we find no compelling evidence for the presence of two types of Sy2s, one type of them has been proposed to be intrinsically different from Sy2s claimed in Unification Model."
    },
    "0207/astro-ph0207031_arXiv.txt": {
        "abstract": "{ High resolution spectra of seven early B-type giant/supergiant stars in the SMC cluster NGC330 are analysed to obtain their chemical compositions relative to SMC field and Galactic B-type stars. It is found that all seven stars are nitrogen rich with an abundance approximately 1.3 dex higher than an SMC main-sequence field B-type star, AV304. They also display evidence for deficiencies in carbon, but other metals have abundances typical of the SMC. Given the number of B-type stars with low rotational projected velocities  in NGC330 (all our targets have $v$sin$i < 50$\\,km/s), we suggest that it is unlikely that the stars in our sample are seen almost pole-on, but rather that they are intrinsically slow rotators. Furthermore, none of our objects displays any evidence of significant Balmer emission excluding the possibility that these are Be stars observed pole-on. Comparing these results with the predictions of stellar evolution models including the effects of rotationally induced mixing, we conclude that while the abundance patterns may indeed be reproduced by these models, serious discrepancies exist. Most importantly, models including the effects of initially large rotational velocities do not reproduce the observed range of effective temperatures of our sample, nor the currently observed rotational velocities. Binary models may be able to produce stars in the observed temperature range but again may be incapable of producing suitable analogues with low rotational velocities. We also discuss the clear need for stellar evolution calculations employing the correct chemical mix of carbon, nitrogen and oxygen for the SMC. ",
        "introduction": "NGC330 is one of the brightest and most populous young clusters in the Small Magellanic Cloud (SMC). The photometric surveys of Arp (\\cite{Arp}), Robertson (\\cite{Rob74}) and Carney et al. (\\cite{Car}) illustrate the key features of the cluster's colour-magnitude diagram, namely the presence of two groups of blue and red supergiants well separated from the cluster's supposed main-sequence blue plume.  These two groups of stars have been widely interpreted as core helium burning stars and the cluster is therefore considered as a key test of stellar evolution theory and physics for stars of intermediate mass in a low metallicity regime.  Essentially the ratio of blue (B) to red (R) supergiants is an indicator of the relative times a massive star spends in the these phases, and these quantities are extremely sensitive to assumptions made concerning convection and mixing. In fact the B/R ratio in NGC330 is generally assumed to be representative of the SMC as a whole and is used as a calibrator for stellar evolution calculations at low metallicity (Stothers \\& Chin \\cite{SCa}, \\cite{SCb}; Keller et al. \\cite{Kel00}; Chiosi et al. \\cite{Chi95}). The specific problem of the B/R ratio as a function of metallicity has been discussed by Langer \\& Maeder (\\cite{Lan95}), where a more detailed discussion of the various treatments of convection and overshooting may be found.  The interpretation of the cluster's HR diagram is complicated by the surprise finding that many main sequence B-type stars in the cluster have H$\\alpha$ emission implying a very high incidence of Be stars (Feast \\cite{Fst72}). A subsequent intermediate band and H$\\alpha$ photometric study indicated that at least 60\\% of all main-sequence B-type stars are of Be-type (Grebel et al. \\cite{Gre96}), this high fraction being confirmed independently by the spectroscopic observations of Lennon et al. (\\cite{Len94}), Mazzali et al (\\cite{Maz96}) and Keller \\& Bessell (\\cite{Kel98}). As with the ratio of B/R supergiants, the ratio Be/B-type stars in NGC330 is often taken as being representive for the SMC metallicity (Maeder et al. \\cite{Mae99}). A second complication arises concerning uncertainty over the metallicity of stars in NGC330; some estimates of the metallicity based upon spectroscopy of the brightest K and F-type supergiants (Spite et al. \\cite{Spi91}) and one B-type giant (Reitermann et al. \\cite{Rei90}) imply that these objects are metal poor even with respect to SMC field stars, while Str\\\"omgren photometric observations of supergiants by Grebel \\& Richtler (\\cite{Gre92}) have been interpreted as evidence for a metal deficiency of 0.5 dex with respect to field stars. However more recent analyses of K-type supergiants have tended to suggest that this difference in metallicity is much smaller, or indeed not significant (Hill \\cite{Hil99}), confirming the results obtained from the analysis of two B-type stars in the cluster by Lennon et al. (\\cite{Len96}, hereafter Paper I).  The spectroscopic work of Lennon et al. (\\cite{Len94}) also found that the bright non-Be and weak Be-type stars occupied that region of the HR-diagram known as the post main sequence gap, or blue Hertzsprung gap (BHG). That is, they are giant/supergiant stars lying red-wards of the main sequence band, but blue-wards of the A/F-type supergiant regime. Caloi et al. (\\cite{Cal93}) and Grebel et al. (\\cite{Gre96}) have also commented on this fact, the latter suggesting that these stars are most likely a mixture of rapidly rotating B/Be-type stars of varying orientation and blue stragglers formed by interaction in binary stars.  Keller et al. (\\cite{Kel00}) also attempted to address this problem using far-UV photometry (from the F160BW filter on WFPC2 of HST) to constrain B-type stellar effective temperatures and find significantly fewer stars in the blue Hertzsprung gap (BHG). However they assumed that the logarithmic surface gravities were 4.0 in their work (Keller, private communication), which may result in spuriously high effective temperatures for stars near the turn-off since they have much lower surface gravities (Lennon et al. \\cite{Len94}). Note that Caloi et al. (\\cite{Cal93}) also adopted lower values for the surface gravities.  Clearly a detailed spectroscopic analysis of the BHG stars leading to estimates of both stellar parameters and atmospheric abundances is needed for comparison with the predictions of various stellar evolution calculations.  In Paper I, we derived metallicities of two such B-type stars in NGC330 and while we found them in general to be compatible with that of the SMC field both stars had a significant nitrogen overabundance. The magnitude of the nitrogen enrichment was uncertain due to the small number of NGC330 targets analysed and also the difficulty in estimating the low nitrogen abundance of the SMC field, at least, from B-type stars.  Also in Paper I we found that the carbon abundance was not significantly depleted, contrary to what one expects if the nitrogen were produced in the CN cycle. The carbon abundance was uncertain however, and coupled with the uncertainty of the magnitude of the nitrogen enrichment, made interpretation difficult. An additional puzzling aspect was that both stars are narrow-lined and therefore if the nitrogen enhancements are the result of high rotation we must be observing them almost pole-on, which seemed unlikely give that they were drawn from the sample of about 20 stars observed by Lennon et al. (\\cite{Len94}).  In the present paper we analyse seven targets in NGC330 (including the two discussed in Paper I) belonging to both the blue supergiant group and the tip of the blue main sequence plume discussed above. Hence all stars lie in, or close to, the BHG. We estimate stellar parameters, radial and projected rotational velocities, as well as both absolute abundances and differential abundances relative to a galactic target in the h and $\\chi$ Persei cluster (Vrancken et al. \\cite{Vra00}) and to an SMC field star, AV304. For the latter we utilise the results from a recent analysis (Rolleston et al \\cite{Rol02}) based on high quality VLT data, which now give us a reliable estimate of the pristine nitrogen abundance in unevolved B-type stars in the SMC. We also attempt to provide improved estimates for carbon abundances and compare our results with stellar evolution calculations, including the recent models of Maeder \\& Meynet (\\cite{Mae01}) which include the effects of rotationally induced mixing. ",
        "conclusions": "\\subsection{The chemical composition of our B-type sample in NGC330} The principle conclusion from the differential analysis is that the cluster targets have a much higher nitrogen and possibly lower oxygen abundance than AV304. Relative nitrogen to oxygen abundances, $[\\frac{N}{O}]$~of $-0.40\\pm$0.24 dex and $-$1.6 dex are deduced for the NGC330 targets and AV304 respectively. For the former the estimates from the individual stars have been weighted by the number of lines observed (although if the stars are uniformly weighted the ratio is only changed by 0.1 dex), whilst for the latter the uncertainty in the individual element abundances imply an error in the ratio of approximately 0.2 dex. For the NGC330 targets, if the errors in the ratios are randomly distributed, the error on the mean would be reduced to approximately $\\pm$0.1 dex.  As discussed in Paper I, the theoretical \\ion{N}{ii} and \\ion{O}{ii} line strengths have a similar dependence on the adopted atmospheric parameters and hence the estimated nitrogen to oxygen abundance ratios are unlikely to be affected by uncertainties in these quantities. Hence we conclude that the $[\\frac{N}{O}]$~ratio is enhanced by 1.2 dex with respect to AV304 and this estimate is unlikely to be significantly affected by uncertainties in atmospheric parameters or the relatively simple LTE analysis adopted.  The simplest explanation for this nitrogen enhancement is that it represents the products of hydrogen burning by the CNO bi-cycle. In such circumstances, it might be expected that there were a corresponding enhancement in helium and underabundances in carbon and possibly oxygen, with the sum of CNO nuclei in the NGC330 stars comparable to that in AV304.  It is therefore important to try to map the abundances for the NGC330 targets as derived from our differential analysis onto an absolute scale. There are a number of options available to us, which we now discuss.  If we assume that the composition of AV304 is the baseline initial composition of NGC330, then we can use the difference in the LTE abundances of the NGC330 targets relative to AV304. However we must also address the probable impact of non-LTE (NLTE) effects on these abundances. We have used NLTE calculations similar to those discussed in McErlean et al. (\\cite{McE00}) to estimate the difference in NLTE and LTE CNO abundance estimates for AV304. The corrections are approximately +0.07, -0.11 and -0.07 dex respectively for C, N and O.  In addition we note that our carbon abundance in AV304 relies heavily on the 4267\\AA\\ line which is well known to give spuriously low abundances. Following the discussion of Vrancken et al. (\\cite{Vra00}) and comparing their results with those of Gies \\& Lambert (\\cite{Gie92}) we further correct the carbon abundance by +0.34 dex.  Our final CNO NLTE abundances estimates for AV304 are then 7.41, 6.55 and 8.16 dex in good agreement with the H\\,{\\sc ii} region results of 7.4, 6.6 and 8.1 dex as summarized in the discussion of baseline SMC abundances for A-type supergiants in the SMC by Venn (\\cite{Venn99}).  We can now use these modified NLTE abundance estimates for AV304 and the difference in LTE abundances listed in Tables \\ref{abs_analysis} and \\ref{gal_stand} to estimate absolute CNO abundance for our NGC330 targets; these are summarized in Table \\ref{CNO}. In turn we can then estimate sum of the CNO and CN abundances for AV304, which are are 8.24 and 7.46 dex respectively.  These may be compared with the mean NGC330 totals of 8.17 and 7.71 dex respectively. There is some slight evidence for an increase in CN but the difference of +0.25 should be compared with uncertainties in the mean carbon and nitrogen abundances of 0.15 and 0.18 dex (the total being dominated by the more abundant species).  The sum of CNO is in good agreement but again the total is dominated by the oxygen abundance for which the uncertainty is approximately 0.13 dex. Clearly it is unproductive to compare summations of abundances when one species is substantially more abundant than all other species, and the uncertainty in that abundance is similar to, or larger than, that of the less abundant species.  We arrive at a similar picture if we instead consider the differential abundances of the NGC330 stars relative to BD+56\\,576 using the NLTE abundances published by Vrancken et al. (\\cite{Vra00}) to put them on an absolute scale.  This results in CNO abundances for the NGC330 targets of 7.28, 7.52 and 7.98 dex respectively which are in good agreement with the values obtained using AV304 as the standard (see Table \\ref{CNO}).  This is possibly fortuitous, and given the difficulty in modeling the 4267\\AA\\ line, unexpected for carbon. Nevertheless it reinforces the previous discussion of the absolute abundances.  We conclude that the chemical peculiarities of the NGC330 targets may be understood in terms of nuclear processing by the CNO bi-cycle, perhaps with some weak evidence that ON processing has occurred, but it is not necessary to invoke primary nitrogen production (although this cannot be precluded).  We can also search for correlations between element abundances within the NGC330 targets. Linear least squares fits show a positive correlation between the helium and nitrogen abundances and negative correlations between the carbon and nitrogen and between the oxygen and nitrogen abundances. Interestingly all these trends are consistent with the transformation of hydrogen into helium using the CNO bicycle. Unfortunately however, none of the correlations are convincing and the coefficients are not significantly different from zero at even the 1$\\sigma$\\ level.  \\subsection{Other stellar abundances in NGC330 and the SMC}  Korn et al. (\\cite{Korn00}) have recently published C, O, Mg and Si NLTE abundances for another B-type giant in NGC330, the star B30,  and their abundances are in good agreement with ours given the magnitude of the NLTE corrections and the uncertainties in both studies.  Unfortunately they do not give a nitrogen abundance but two other similar SMC giants in their sample have NLTE nitrogen abundances of 7.3 and 7.2 dex. We will return to these stars in the discussion of their evolutionary status below.  Furthermore an LTE nitrogen abundance for star B30 was published by Reitermann et al. (\\cite{Rei90}) who obtained a value of 7.4 dex for a microtubulence of 5 km/s.  There have also been many studies of cool giants and supergiants in NGC330. These results are summarized in Table \\ref{CNO}.  Of interest here are CNO abundances and there are two recent estimates for NGC330 stars by Hill (\\cite{Hil99}, H99) and Gonzalez \\& Wallerstein (\\cite{Gon99}, GW). One obvious difference between these is that the mean nitrogen abundance of GW is systematically lower than that of H99. However we note that the nitrogen abundance is derived from molecular CN features, and depends on the adopted carbon abundance (which is derived from C$_{2}$).  While H99 independently derive their carbon abundances, GW adopted mean values from the literature and there is a small systemic overestimation relative to H99 (approximately 0.2 dex).  Such a small change in the carbon abundance is the most likely reason for their low nitrogen abundances and given that the carbon abundance of H99 may be more reliable and that their results agree better with other samples of evolved B, A and F-type stars in the SMC, we prefer their results for carbon and nitrogen.  Comparing with other SMC samples we note that our mean CNO abundances are in good agreement with the results of Venn (\\cite{Venn99}) and Hill et al. (\\cite{Hil97}) although the NGC330 stars may be mildly metal poor. There have been previous suggestions that NGC330 may be relatively metal poor with respect to the SMC (Grebel \\& Richtler \\cite{Gre92}) but our results confirm recent work in that any metal deficiency must be small ($< 0.2$ dex). We therefore conclude that the pattern of the mean CNO abundances found in the NGC330 B-type giants is very similar to that found for samples of other evolved A, F and K-type giants and supergiants. We also note that Dufton et al. (\\cite{Duf00}) investigated a large sample of B-type supergiants in the SMC and the nitrogen abundances found in their less luminous stars, although uncertain, are comparable to those found here.  \\subsection{Evolutionary status}  It is useful to preface our discussion of the evolutionary status of these stars by considering their positions in the HR-diagram. We follow the procedure described in VDLL, adopting a distance modulus for the SMC of 18.9 and extinction estimates as discussed in section 3.1.  As in VDLL we estimate spectroscopic masses (M$_{spec}$) from our derived stellar radii and surface gravities. Figure \\ref{hrdiag} illustrates the positions of the stars compared with the non-rotating stellar evolution tracks of Charbonnel et al. (\\cite{Cha93}) which are computed with metallicity of $Z = 0.004$ (SMC-like).  Again following VDLL we estimate evolutionary masses (M$_{evol}$) by interpolation in this diagram and these may be compared with M$_{spec}$ in Table \\ref{masses}. where all derived quantities are summarized. We note that Maeder \\& Meynet (\\cite{Mae01}) have produced a grid of calculations including the effects of rotation for a metallicity of $ Z = 0.004$. Figure \\ref{hrdiag} also shows the locus of the end of H-burning main sequence phase for an assumed initial rotational velocity of 300 km/s (taken from their Figure 6). Note that this locus also corresponds to the approximate position of the end of the main sequence for the non-rotating models. This is a consequence of the fact that the main sequence widening in Charbonnel et al. comes from their inclusion of convective overshooting, while a similar effect is obtained by Maeder \\& Meynet with rotationally induced mixing alone.  The fact that overshooting and rotation both have similar results in this respect has been discussed previously, see for example Talon et al. (\\cite{Tal97}).  We will return to this point later in the discussion, after first considering the observed abundance pattern.  \\begin{figure*} \\epsfig{file=fig1.eps} \\caption[]{HR-diagram for stars in NGC330 showing the positions of the stars relative to the stellar evolution tracks (solid lines) of Charbonnel et al. (\\cite{Cha93}) which are computed for a metallicity of $Z = 0.004$. Tracks are labeled with their initial masses. The dashed line represents the approximate position of the end of the core H-burning main sequence for the rotating models of Maeder \\& Meynet (\\cite{Mae01}) for an assumed initial rotational velocity of 300\\,km/s.  Blue stars (analysed here) are labeled with their identifications. For comparison we also show the positions of the red supergiants in NGC330 analysed by Hill (\\cite{Hil99}) and Gonzalez \\& Wallerstein (\\cite{Gon99}), whose chemical compositions we also discuss.  The error bars represent typical uncertainties discussed in the text, for example 15\\% in effective temperature. } \\label{hrdiag} \\end{figure*}  \\begin{table*} \\caption{Derived quantities for B-type stars in NGC330: Stellar radii (R/R$_{\\odot}$), absolute visual magnitudes (M$_v$), bolometric corrections (B.C.), bolometric magnitudes (M$_{bol}$), luminosities (log(L/L$_{\\odot}$) and estimates of spectroscopic (M$_{spec}$) and evolutionary (M$_{evol}$) masses in units of solar mass. The distance modulus to the SMC was assumed to be 18.9 in the derivation of these quantities.  The last two columns compare the difference in spectroscopic and evolutionary masses with the uncertainties in log\\,g ($\\Delta$logg). } \\label{masses} \\begin{center} \\begin{tabular}{lrllllrrrc} \\hline \\hline Star & R/R$_{\\odot}$ & M$_v$ & B.C. & M$_{bol}$ & log(L/L$_{\\odot}$) & M$_{spec}$ & M$_{evol}$ & log(M$_{evol}$/M$_{spec}$) & $\\Delta$logg \\\\ \\hline A01 & 12.6 & -4.48 & -2.50 & -6.98 & 4.67 & 29 & 15 & $-0.29$ & 0.3 \\\\ A02 & 40.4 & -6.28 & -1.49 & -7.77 & 4.98 & 12 & 18 & 0.18 & 0.1 \\\\ A04 & 17.2 & -4.64 & -1.77 & -6.41 & 4.44 &  7 & 12 & 0.23 & 0.2 \\\\ B04 &  8.5 & -3.58 & -2.34 & -5.92 & 4.25 & 10 & 11 & 0.04 & 0.2 \\\\ B22 & 17.6 & -4.89 & -2.03 & -6.92 & 4.65 & 11 & 14 & 0.10 & 0.2 \\\\ B32 & 11.5 & -4.17 & -2.23 & -6.40 & 4.45 &  5 & 12 & 0.38 & 0.3 \\\\ B37 & 31.9 & -5.99 & -1.76 & -7.75 & 4.98 &  9 & 18 & 0.30 & 0.1 \\\\ B30 & 21.6 & -5.08 & -1.64 & -6.72 & 4.54 & 11 & 13 & 0.07 & 0.2 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   Nitrogen enrichment is clearly a sign of contamination of the surface by the products on the nuclear burning, in this case from the CNO bi-cycle. What is especially interesting about the current sample of stars (to which we can add the star B30 mentioned above) is that they represent a coeval group of stars with rather homogeneous properties, whose N/C surface abundance ratio is enhanced by typically a factor of 10.  Stellar models which include the effects of rotationally induced mixing have been proposed as a means of producing the observed nitrogen enhancements in main sequence and evolved massive stars.  Maeder \\& Meynet (\\cite{Mae01}) have followed the evolution of the surface abundances for a range of initial rotational velocities.  Comparing our results to the models with high initial rotational velocity, $v_{\\rm ini} = 300$ km/s, and therefore relatively large nitrogen enhancements, we find that the best agreement is with their red supergiant or blue loop stars.  At the end of the main sequence these models predict an increase in [N/H] by a factor of only 3, or about 0.5 dex. Note however that their initial abundance ratios are assumed to be solar and therefore they overestimate the initial nitrogen abundance significantly, adopting one fifth solar, which is approximately 7.3 dex in our notation.  We can perform a simple recalibration of their results by assuming that a calculation with an initial lower nitrogen abundance of 6.6 dex (but similar carbon and oxygen) produces the same excess of nitrogen in absolute terms. In this case the nitrogen abundance at the end of the main sequence would be approximately 7.7 dex rather than 7.8 dex, a consequence of the fact that the initial nitrogen abundance is negligible compared with that which is produced by the star.  These models therefore could conceivably reproduce the observed nitrogen enhancements.  In fact if the initial N/O or N/(O+C) abundance ratios are small and can be neglected, which in the case of the SMC is a good approximation, then it is easy to show that one only needs a relatively small fraction of core material in ON equilibrium to produce a big change in the observed nitrogen abundance.  Nevertheless, as Figure \\ref{hrdiag} shows, there are other more serious discrepancies. For example, our objects tend to lie red-wards of the main sequence despite the widening provided by the rotating models (and models with convective overshooting). In addition all the B-type giants/supergiants in our sample are slow rotators with values of $v$sin$i$ less than 50\\,km/s.  By contrast the rotating models require very high initial rotation to produce the enhanced nitrogen but do not predict significant slow down by the end of the main sequence. As discussed in Paper I, while there may be a selection effect in our observer sample (we can only analyse the slow rotators) it is highly unlikely that there are so many fast rotators in NGC330 oriented pole on. For example, both Mazzali et al. (\\cite{Maz95}) and Keller \\& Bessell (\\cite{Kel98}) give $v$sin$i$ values for 22 of the brightest B and Be-type stars in NGC330.  They find that maximum values lie in the range 300 -- 400\\,km/s and if we assume that our sample have similar rotation rates ($v$) this implies that sin$i$ is less than 10 degrees.  In other words if the distribution of $i$ is random we should expect our sample to be drawn from about 1.5\\% of all the B-type stars in NGC330.  This is clearly incompatible with the number of B-type stars in NGC330 in the relevant magnitude range. One is left with the conclusion that our objects are intrinsically slow rotators at the present time.  Either they were fast rotators in the past, and have somehow slowed down, or some process other than rotationally induced mixing is responsible for the observed abundance pattern.  Given the similarity between the carbon and nitrogen abundances in the blue and red stars in NGC330 it is tempting to invoke blue loops as a means of explaining our abundances. This seems unlikely given that no stellar evolution calculations predict loops which progress hotter than effective temperatures corresponding to late-B spectral types.  While Venn (\\cite{Venn99}) invoked blue loops for the A-type supergiants this does not seem a viable option for our early B-type stars, despite similarities in CNO abundances. Mass-transfer in binaries may also be invoked to explain enhanced nitrogen abundances and Wellstein et al. (\\cite{Wel01}) have recently produced models which produce nitrogen enriched blue stars which can reside in the post main-sequence gap. Such stars may also appear to be under-massive for their luminosities and is therefore tempting to ascribe the discrepancies between spectroscopic and evolutionary masses in Table \\ref{masses} to binarity. One should be cautious however because our spectroscopic masses were estimated from the derived surface gravities and in some cases the uncertainties in this quantity are quite substantial.  The final two columns of Table \\ref{masses} compares the differences between spectroscopic and evolutionary masses with the uncertainties in gravity and in all but three cases (A02, B32 and B37) the mass differences are easily accounted for.  There does appear to be a suggestion of a correlation between nitrogen enhancement and luminosity in our sample.   Stars A02 and B37 are the most luminous and most N-rich of our sample,  however their positions are perhaps consistent with being core-helium burning stars on their way red-wards (this phase represented by the slight kink in the 20 solar mass track for the non-rotating models). The problem for the rotating models is that this phase becomes progressively cooler and shorter lived as initial rotational velocity is increased, with only the slow rotators spending any significant time in this part of HR-diagram.  However, such stars should not be significantly N-enriched, in contradiction to the observations. B32 would appear to be the best candidate for the binary evolution hypothesis, the real problem being that much better constraints are needed on the gravity of this object, and indeed all other stars in our sample, before definitive statements can be made about possible mass discrepancies. Finally, while binarity  appears to be an attractive scenario, it has a significant problem in that it is expected that the products of mass accretion will be fast rotators having been spun up by the accreted material. In addition radial velocities of all our stars are typical of the NGC330 cluster (Table 1) although it must be noted that expected radial velocity amplitudes of the kind of systems predicted by the models of Wellstein \\& Langer (\\cite{Wel01}) are typically the order of 10--20 km/s.  Finally we note that VDLL carried out a similar study to the present one but for the solar metallicity cluster h+$\\chi$ Per.  They did not find evidence for significant nitrogen overabundances in any of the evolved B-type stars in this cluster.  However if we take the nitrogen enrichments found here, in absolute terms, apply these to their galactic counterparts it is clear that this leads to enhancements the order of a factor of 2--3.  For some objects in h+$\\chi$ Per this magnitude of a nitrogen enhancement is consistent with the observations.  It is simply more obvious in the SMC stars given their initial very low nitrogen abundance.  \\begin{table*} \\caption{Comparison of our abundance estimates corrected for NLTE effects with results for other stars in NGC330, viz. the B-type star B30 analysed by Korn et al. (NLTE C and O) and Reitermann et al. (LTE N) (B30) and cool supergiants from Hill (H99) and Gonzalez \\& Wallerstein (GW). Also tabulated are other SMC abundance estimates, viz. the main sequence star AV304 (Rolleston et al. \\cite{Rol02} corrected for NLTE effects), A-type supergiants from Venn (A-stars), B-type giants and supergiants from Korn et al. (B-stars) and cool supergiants from Hill et al. (K-stars).} \\begin{center} \\label{CNO} \\begin{tabular}{llllllllll} \\hline \\hline & & \\multicolumn{4}{c}{NGC330}& & \\multicolumn{3}{c}{SMC supergiants} \\\\ \\cline{3-6} \\cline{8-10} Element & AV304 & This paper & B30 & H99 & GW & &A-stars & B-stars & K-stars \\\\ \\hline C & 7.47 & 7.26 & 7.3 & 7.40 & 7.55 & & -- & 7.40 & 7.55 \\\\ N & 6.55 & 7.52 & 7.4 & 7.21 & 6.96 & & 7.33 & 7.25 & 7.51 \\\\ O & 8.16 & 7.98 & 8.25 & 7.71 & 7.92 & & 8.14 & 8.15 & 8.01 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "0207/astro-ph0207388_arXiv.txt": {
        "abstract": "Balmer lines are an important diagnostic of stellar atmospheric structure, since they are formed at a wide range of depths within the atmosphere. The different Balmer lines are formed at slightly different depths making them useful atmospheric diagnostics. The low sensitivity to surface gravity for stars cooler than $\\sim$8000~K makes them excellent diagnostics in the treatment of atmospheric convection. For hotter stars Balmer profiles are sensitive to both effective temperature and surface gravity. Provided we know the surface gravity of these stars from some other method (e.g. from eclipsing binary systems), we can use them to determine effective temperature.  In previous work, we have found no significant systematic problems with using $uvby$ photometry to determine atmospheric parameters of fundamental (and standard) stars. In fact, $uvby$ was found to be very good for obtaining both $T_{\\rm eff}$ and $\\log g$. Using H$\\alpha$ and H$\\beta$ profiles, we have found that both the Canuto \\& Mazzitteli and standard Kurucz mixing-length theory without approximate overshooting are both in agreement to within the uncertainties of the fundamental stars. Overshooting models were always clearly discrepant. Some evidence was found for significant disagreement between {\\em all\\/} treatments of convection and fundamental values around 8000$\\sim$9000~K, but these results were for fundamental stars {\\em without\\/} fundamental surface gravities. We have used stars with fundamental values of both $T_{\\rm eff}$ and $\\log g$ to explore this region in more detail. ",
        "introduction": "Balmer lines are an important diagnostic of stellar atmospheric structure since they are formed at a wide range of depths within the atmosphere. The depth of formation of H$\\alpha$ is higher than that of H$\\beta$, thus observations of these profiles provide useful diagnostics (e.g. Gardiner 2000). Balmer profiles are relatively insensitive to surface gravity for stars cooler than $\\sim$8000~K (Gray 1992, see also Heiter et al. 2002), whilst sensitive to the treatment of atmospheric convection (e.g. van't Veer \\& M\\'{e}gessier 1996, Castelli et al. 1997, Gardiner 2000, Heiter et al. 2002). For stars hotter than $\\sim$8000~K the profiles are sensitive to both effective temperature and surface gravity. However, provided we know surface gravity from some other means (e.g. from eclipsing binary systems), we can use them to determine effective temperature.  In previous work, Smalley \\& Kupka (1997) found no significant systematic problems with $uvby$ and fundamental (and standard) stars. In fact, $uvby$ was found to be very good for obtaining $T_{\\mathrm eff}$ and $\\log g$. Using H$\\alpha$ and H$\\beta$ profiles, Gardiner et al. (1999) found that both the Canuto \\& Mazzitteli (1991, 1992) and standard Kurucz (1993) mixing-length theory without overshooting (see Castelli et al. 1997) are both in agreement to within the uncertainties of the fundamental stars. Overshooting models were always clearly discrepant. However, Gardiner et al. (1999) found some evidence for significant disagreement between {\\em all\\/} treatments of convection and fundamental values around 8000$\\sim$9000~K. In this region the effects of $\\log g$ cannot be ignored, and the majority of the $T_{\\mathrm eff}$ stars did not have fundamental values of $\\log g$. We have used binary systems with fundamental values of $\\log g$, determined fundamental values of $T_{\\mathrm eff}$ and compared the results with those from fitting models to Balmer-line profiles. ",
        "conclusions": "Balmer line profiles have been fitted to the fundamental binary systems. To within the errors of the fundamental $T_{\\mathrm eff}$ values, neither the H$\\alpha$ or H$\\beta$ profiles exhibit any significant discrepancies for the CM and MLT without approximate overshooting models. As in previous work, the MLT with overshooting models are found to be discrepant. Moreover, there are no systematic trends, such as offsets, between results from H$\\alpha$ and H$\\beta$ as long as $\\alpha$ in MLT models is chosen small enough (e.g. 0.5). The discrepancies exhibited by the fundamental $T_{\\mathrm eff}$ stars in Gardiner et al. (1999) can be explained by rapid rotation in two cases and by the fact that the Balmer profiles become sensitive to $\\log g$ and less sensitive to $T_{\\mathrm eff}$ in the other two cases. However, for the time being the lack of any stars with fundamental values of both $T_{\\mathrm eff}$ and $\\log g$ in this region precludes the conclusion that there is not a problem with the models in the $T_{\\mathrm eff}$ range 8000 $\\sim$ 9000~K.  Full details of this work are given in Smalley et al. (2002)."
    },
    "0207/astro-ph0207341_arXiv.txt": {
        "abstract": "We present here the results of an investigation of the pulse averaged and pulse phase resolved energy spectra of two high luminosity accretion powered X-ray pulsars SMC~X-1 and LMC~X-4 made with ASCA. The phase averaged energy spectra definitely show the presence of a soft excess in both the sources. If the soft excess is modeled as a separate black-body or thermal bremsstrahlung type component, pulse phase resolved spectroscopy of SMC~X-1 shows that the soft component also has a pulsating nature. Same may be true for LMC~X-4, though a very small pulse fraction limits the statistical significance. The pulsating soft component is found to have a nearly sinusoidal profile, dissimilar to the complex profile seen at higher energies, which can be an effect of smearing. Due to very high luminosity of these sources, the size of the emission zone required for the soft component is large (radius $\\sim$300--400 km). We show that the pulsating nature of the soft component is difficult to explain if a thermal origin is assumed for it. We further investigated with alternate models, like inversely broken power-law or two different power-law components and found that these models can also be used to explain the excess at low energy. A soft power-law component may be a common feature of the accreting X-ray pulsars, which is difficult to detect because most of the HMXB pulsars are in the Galactic plane and experience large interstellar absorption. In LMC X-4, we have also measured two additonal mid-eclipse times, which confirm the known orbital decay. ",
        "introduction": "The X-ray continuum spectra of accreting pulsars are often described as a broken power-law or a power-law with exponential cutoff. The break in the spectrum is in the range of 10--20 keV and power-law photon index below the break energy is in the range of 0--1 (White, Nagase \\& Parmar 1995). Some binary X-ray pulsars which are away from the Galactic plane and therefore experience less interstellar absorption, show the presence of a soft component in the spectrum which is often modeled as a black-body and/or thermal bremsstrahlung emission (SMC~X-1: Marshall, White, \\& Becker 1983; Woo et al. 1995; Wojdowski et al. 1998; LMC~X-4: Dennerl 1989; Woo et al. 1996; RX~J0059.2--7138: Kohno, Yokogawa, \\& Koyama 2000; 4U~1626--67: Orlandini et al. 1998; Her~X-l: McCray et al. 1982, Dal Fiume et al. 1998; Oosterbroek et al. 1997, 2000; Endo, Nagase \\& Mihara 2000) or an inversely broken power-law (XTE~J0111.2--7317: Yokogawa et al. 2000b). The soft component in 4U~1626--67, when modeled as a black-body emission requires the size of emission region to be comparable to that of the neutron star because the intrinsic luminosity of this source is of the order of 10$^{35}$ erg s$^{-1}$ (Orlandini et al. 1998). On the other hand, the soft component in Her~X-1 can be modeled as a black-body which is reprocessed emission from the innermost part of the accretion disk (Endo et al. 2000). However, for the bright pulsars in the Magellanic Clouds for which the distance is of the order of 50--60 kpc and the luminosity is close to the Eddington limit, the luminosity of the soft excess is about an order of magnitude larger compared to the same of Her~X-1. Therefore, a pulsating nature of the soft component that has been observed in some high luminosity X-ray pulsars (LMC~X-4 : Woo et al. 1996 and XTE~J0111.2--7317 : Yokogawa et al. 2000b) needs to be probed with greater detail.  To investigate the pulse phase dependence of the soft component of accreting X-ray pulsars in detail, we have chosen two luminous X-ray pulsars SMC~X-1 and LMC~X-4, which are in the Magellanic Clouds and suffer less interstellar absorption. SMC~X-1 and LMC~X-4 are two bright, eclipsing, accreting, binary X-ray pulsars with spin periods of $\\sim$0.7 and $\\sim$13.5 s and binary periods of $\\sim$3.9 and $\\sim$1.4 day respectively. The companion of SMC~X-1 is a B0 supergiant while the companion star of LMC~X-4 is of type O7III-V. The binary orbits of both the systems are nearly circular. The orbital period of the two binaries are found to decay with time scale of 3 $\\times$ 10$^5$ yr in SMC~X-1 (Wojdowski et al. 1998), and 10$^6$ yr in LMC~X-4 (Levine, Rappaport \\& Zojcheski 2000). Another striking similarity between these two sources is a long-period of 50--60 day and 30.5 day respectively, that is known to be quasi-stable in SMC~X-1 (Wojdowski et al.  1998) and stable in LMC~X-4 (Lang et al. 1981). The long-period is believed to be a result of (quasi) periodic obscuration of the neutron star by a precessing accretion disk, similar to that in Her~X-1. Broad band energy-spectra (0.2--37 keV) of these two sources were studied by performing a combined fit to the observations made with ROSAT and GINGA (Woo et al. 1995; 1996). In addition to a cutoff power-law type component in SMC~X-1 (photon index of $\\sim$0.93, $E_{\\rm C}$ = 5.6 keV, $E_{\\rm F}$ = 15 keV) and power-law type component in LMC~X-4 (photon index of $\\sim$0.67) broad iron emission lines and soft excess were detected. The soft component was modeled as a single black-body component ($kT_{\\rm BB}$ = 0.16 keV) in SMC~X-1 and as a sum of a low temperature black-body emission ($kT_{\\rm BB}$ = 0.03 keV) and a thermal bremsstrahlung emission ($kT_{\\rm TB}$ = 0.35 keV) in LMC~X-4. Beppo-SAX observations also showed presence of similar soft component (La Barbera et al. 2001). From HEAO-1 observations, the soft X-rays from SMC X-1 were found to be nonpulsating (Marshal et al. 1983). However, the ROSAT and ASCA observations detected clear pulsations with a pulse profile different from the hard component (Wojdowski et al. 1998). Pulse phase resolved spectroscopy of combined ROSAT and GINGA data of LMC~X-4 revealed modulation of the thermal bremsstrahlung and iron line emission components with pulse phase (Woo et al. 1996).  We note that the combined fit of the ROSAT and GINGA spectra of these two sources were performed on non-simultaneous observations and the intensity and spectral shape of these sources are known to be variable. We have carried out pulse phase averaged and pulse phase resolved spectral studies in the 0.5--10.0 keV band. In this paper we present the results of our investigation to the nature of the soft excess through pulse profiles in different energy bands and variations of the different spectral components with pulse phase. ",
        "conclusions": "In the present work we have found that these two high luminosity pulsars show pulsations in the entire energy band of 0.5--10.0 keV, with strong energy dependence in the pulse shape. The pulse phase averaged energy spectra definitely show the presence of soft excess in both the sources. If the soft excess is modeled as a separate black-body or thermal bremsstrahlung type component, pulse phase resolved spectroscopy shows that the soft component has a pulsating nature in SMC X-1, which may also be true for LMC X-4. We have found that the pulse profile of the soft component is nearly sinusoidal, significantly different from the sharp, complex profile of the hard power-law component. Due to very high luminosity of these sources, the size of the emission zone required for the soft component is large (radius $\\sim$300--400 km) and we find that it is difficult to explain the pulsations detected at low energies. We have found that alternate models like inversely broken power-law or two different power-law components can also be used to describe the spectra."
    },
    "0207/astro-ph0207177_arXiv.txt": {
        "abstract": "In this contribution we present a few selected examples of how the latest generation of space-based instrumentation -- NASA's {\\it Chandra} X-ray Observatory and the Far-Ultraviolet Spectroscopic Explorer ({\\it FUSE}) -- are finally answering old questions about the influence of massive star feedback on the warm and hot phases of the ISM and IGM. In particular, we discuss the physical origin of the soft thermal X-ray emission in the halos of star-forming and starburst galaxies, its relationship to extra-planar H$\\alpha$ emission, and plasma diagnostics using {\\it FUSE} observations of O {\\sc vi} absorption and emission. ",
        "introduction": "Massive stars exercise a profound influence over the baryonic component of the Universe, through their return of ionizing radiation, and via supernovae (SNe), kinetic energy and metal-enriched gas, back into the ISM from which they form --- usually called ``feedback''. Feedback influences gas-phase conditions in the immediate environment of the clusters within which the massive stars form, on galactic-scales the phase structure and energetics of the ISM, and on multi-Mpc scales the thermodynamics and enrichment of the inter-galactic medium (IGM).  The vast range of spatial scales involved is only one of the difficulties encountered in attempting to study feedback. Another is the broad range of complicated gas-phase physics -- (magneto)hydrodynamic effects such as shocks and turbulence, thermal conduction, and non-ionization equilibrium emission processes. A final complication is that much of the energy and metal-enriched material involved is in the hard-to-observe coronal ($T \\ga 10^{5}$ K) and hot ($T \\ga 10^{6}$ K) gas phases.   \\begin{figure}[!t] \\plotone{strickland_f1.eps} \\caption[Sample galaxies]{A few examples from our larger survey of edge-on star-forming galaxies (Strickland et al 2002b). The top panels plot {\\it Chandra} ACIS-S 0.3 -- 2.0 keV X-ray surface brightness. All galaxies are shown on the same logarithmic intensity scale, with contours spaced by 0.5 dex. Point sources have been left in these images, but are removed for scientific analysis. The lower panels are continuum-subtracted H$\\alpha$+[N {\\sc ii}] emission, with the IRAS $f_{60}/f_{100}$ ratio (a measure of the SF intensity) given at the bottom of the image. Each image shows a $20 \\times 20$ kpc region. Tick-marks represent $1\\arcmin$ --- {\\em Chandra}'s resolution is sub-arcsecond. } \\label{fig:dks:sample_xray_ha} \\end{figure} ",
        "conclusions": "In the three years since their launch in mid-1999, the un-matched capabilities of both {\\it Chandra} and {\\it FUSE}, have allowed subtantial progress to be made in understanding the physics of the feedback in star-forming galaxies. Old questions have been answered, and new ones have arisen. With several, no-doubt productive, years left for each instrument, it is likely that progress will continue to be rapid in understanding the relationship between massive stars and the warm and hot phases of the ISM and IGM."
    },
    "0207/astro-ph0207494_arXiv.txt": {
        "abstract": "{ Generally, chemical peculiarity found for stars on the upper main sequence excludes $\\delta$ Scuti type pulsation (e.g. Ap and Am stars), but for the group of \\LB stars it is just the opposite. This makes them very interesting for asteroseismological investigations. The group of \\LB type stars comprises late B- to early F-type, Population\\,I objects which are basically metal weak, in particular the Fe group elements, but with the clear exception of C, N, O and S. The present work is a continuation of the studies by Paunzen et al. (1997, 1998), who presented first results on the pulsational characteristics of the \\LB stars. Since then, we have observed 22 additional objects; we found eight new pulsators and confirmed another one. Furthermore, new spectroscopic data (Paunzen 2001) allowed us to sort out misidentified candidates and to add true members to the group. From 67 members of this group, only two are not photometrically investigated yet which makes our analysis highly representative. We have compared our results on the pulsational behaviour of the \\LB stars with those of a sample of $\\delta$ Scuti type objects. We find that at least 70\\% of all \\LB type stars inside the classical instability strip pulsate, and they do so with high overtone modes ($Q$\\,$<$\\,0.020\\,d). Only a few stars, if any, pulsate in the fundamental mode. Our photometric results are in excellent agreement with the spectroscopic work on high-degree nonradial pulsations by Bohlender et al. (1999). Compared to the $\\delta$ Scuti stars, the cool and hot borders of the instability strip of the \\LB stars are shifted by about 25\\,mmag, towards smaller $(b-y)_0$. Using published abundances and the metallicity sensitive indices of the Geneva 7-colour and Str\\\"omgren $uvby\\beta$ systems, we have derived [Z] values which describe the surface abundance of the heavier elements for the group members. We find that the Period-Luminosity-Colour relation for the group of \\LB stars is within the errors identical with that of the normal $\\delta$ Scuti stars. No clear evidence for a statistically significant metallicity term was detected. ",
        "introduction": "In this paper we present an extensive survey to analyse the pulsational characteristics of the \\LB stars. This small group comprises late B- to early F-type, Population\\,I stars which are metal weak (particularly the Fe group elements), but with the clear exception of C, N, O and S. Only a maximum of about 2\\% of all objects in the relevant spectral domain are believed to be \\LB type stars.  Several theories were developed to explain the peculiar abundance pattern for members of this group. The most acknowledged models include diffusion as main mechanism together either with mass-loss (Michaud \\& Charland 1986, Charbonneau 1993) or with accretion of circumstellar material (Venn \\& Lambert 1990, Turcotte \\& Charbonneau 1993). Another two theories deal with the influence of binarity on this phenomenon (Andrievsky 1997, Faraggiana \\& Bonifacio 1999). Heiter (2002) and Heiter et. al (2002) also tried to explain the abundance pattern in the context of the proposed theories.  In general, chemical peculiarity inhibits $\\delta$ Scuti type pulsation (e.g. for Ap and Am stars, see Kurtz 2000 for a recent discussion) but for the group of \\LB stars it is just the opposite. In two previous studies (Paunzen et al. 1997, 1998), we presented non-variable as well as pulsating \\LB stars. Since then, we have observed 22 additional objects and found eight new pulsators and confirmed another. Furthermore, new spectroscopic data (Paunzen 2001) has allowed us to sort out misidentified candidates and to add true members of the group.  Turcotte et al. (2000) investigated the effect of diffusion (probably the main cause of the \\LB phenomenon) on the pulsation of stars at the upper main sequence. Although these authors mainly investigated the theoretical behaviour of apparently metal-rich objects, their conclusions also have an impact for the \\LB group: little direct pulsational excitation from Fe-peak elements was found, but effects due to settling of helium along with the enhancement of hydrogen are important. Turcotte et al. (2000) find that, as their models of peculiar stars evolve, they become generally pulsational unstable near the red edge of the instability strip, whereas the behaviour at the blue edge is mainly sensitive to the surface metal abundance. Although the proposed models are still simplified (e.g. treatment of convection) these preliminary results already point towards the most important effects on the theoretical pulsational instability and behaviour of chemically peculiar stars.  The aim of the present paper is to analyse the pulsational characteristics of the group of \\LB stars and to test for the presence of a possible Period-Luminosity-Colour-Metallicity relation. The latter is especially interesting in the light of the models by Turcotte et al. (2000). The pulsational characteristics of the \\LB group (e.g. ratio of variable to non-variable objects and distribution of pulsational constants) may help to put tighter constraints on these models.  \\begin{table} \\caption[]{Sites, dates and telescopes used for our survey} \\label{sites} \\begin{center} \\begin{tabular}{lcccc} \\hline Site & Date & Telescope & Stars & Ref. \\\\ \\hline APT (Fairborn) & 05.2001 & 0.75 & 3 & 1 \\\\ SAAO & 04.2001 & 0.50 & 8 & 2 \\\\ & 07.2001 \\\\ & 08.2001 \\\\ & 09.2001 \\\\ & 10.2001 \\\\ SAAO & 12.2000 & 0.75 & 9 & 3 \\\\ & 01.2001 \\\\ SAAO & 08.2001 & 1.00 & 1 & 4 \\\\ Siding Spring & 01.2002 & 0.60 & 1 & 5 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table*} \\caption[]{Observing log of eight newly discovered and one confirmed (HD~75654) pulsating \\LB stars. Some information on the comparison stars is also given. The differential light curves are shown in Fig. \\ref{lightcurve}.} \\label{log1} \\begin{center} \\begin{tabular}{rccccccc} \\hline HD & JD & hrs & $m_{\\rm V}$ & Spec. & Freq. & Amp. & Ref. \\\\ & & & [mag] & & [d$^{-1}$] & [mag] \\\\ \\hline 13755 & 2451899 & 3.2 & 7.84 & $\\lambda$\\,Boo & 12.50 & 0.015 & 3 \\\\ & 2451903 & 2.5 &\t\t & \t   & 16.85 & 0.007 \\\\ & 2451905 & 3.1 \\\\ & 2451909 & 3.1 \\\\ 13602 &\t\t\t& \t  & 8.52 & F6 \\\\ 13710 &\t\t\t&\t  & 8.32 & K5 \\\\ \\hline 35242 & 2451900 & 2.1 & 6.35 & $\\lambda$\\,Boo & 38.61 & 0.005 & 3 \\\\ & 2451902 & 5.4 &\t\t & \t   & 34.16 & 0.003 \\\\ & 2451908 & 3.3 &\t\t & \t   & 41.33 & 0.003 \\\\ 35134 &\t\t\t& \t  &\t6.74 & A0 \\\\ 34888 &\t\t\t&\t  & 6.78 & A5 \\\\ \\hline 42503 & 2452291 & 4.2 & 7.45 & $\\lambda$\\,Boo & 7.00 & 0.015 & 5 \\\\ & 2452292 & 1.9 \\\\ 42058 & \t\t& \t  & 6.99 & A0 \\\\ 43452 &\t\t    &\t  & 7.71 & F5 \\\\ \\hline 75654 & 2451898 & 3.0 & 6.38 & $\\lambda$\\,Boo & 14.80 & 0.005 & 3 \\\\ & 2451902 & 1.8 &\t\t & \t   & 15.99 & 0.002 \\\\ & 2451905 & 3.1 \\\\ & 2451906 & 1.2 \\\\ & 2451907 & 3.8 \\\\ & 2451909 & 3.6 \\\\ 74978 &\t\t\t& \t  & 6.87 & A1 \\\\ 75272 &\t\t\t&\t  & 6.98 & B9.5 \\\\ \\hline 111604 & 2452061 & 4.1 & 5.89 & $\\lambda$\\,Boo & 8.77 & 0.020 & 1 \\\\ 112412 & \t\t & \t   & 5.61 & F1 \\\\ 110375 &\t\t &\t   & 8.33 & F5 \\\\ \\hline 120896 & 2452097 & 3.9 & 8.50 & $\\lambda$\\,Boo & 17.79 & 0.010 & 2 \\\\ 121372 & \t\t & \t   & 8.67 & G5 \\\\ \\hline 148638 & 2452097 & 4.6 & 7.90 & $\\lambda$\\,Boo & 16.32 & 0.016 & 2 \\\\ & 2452123 & 5.0 \\\\ 148596 & \t\t & \t   & 8.60 & F2 \\\\ 148573 &\t\t &\t   & 8.63 & B9 \\\\ \\hline 213669 & 2451823 & 6.5 & 7.42 & $\\lambda$\\,Boo & 15.01 & 0.023 & 2 \\\\ & 2451826 & 1.6 \\\\ & 2451827 & 1.1 \\\\ 211878 & \t\t & \t   & 7.70 & F5 \\\\ 214390 & \t\t & \t   & 7.90 & F3 \\\\ \\hline 290799 & 2451904 & 3.0 & 10.63 & $\\lambda$\\,Boo & 23.53 & 0.006 & 3 \\\\ & 2451906 & 4.8 \\\\ 37652  & \t\t & \t   & 7.35 & F5 \\\\ 290798 &\t\t &\t   & 10.40 & A2 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We have investigated the pulsational characteristics of a group of \\LB stars and compared it to a sample of $\\delta$ Scuti pulsators. The latter was chosen such that it matches our program stars within the global astrophysical parameters. The following properties of the \\LB stars are different from those of the $\\delta$ Scuti pulsators: \\begin{itemize} \\item At least 70\\% of all \\LB types stars inside the classical instability strip pulsate \\item Only a maximum of two stars pulsate in the fundamental mode but there is a high percentage with $Q$\\,$<$\\,0.020\\,d (high overtone modes) \\item The instability strip of the \\LB stars at the ZAMS is 25\\,mmag bluer in $(b-y)_0$ than that of the $\\delta$ Scuti stars. \\end{itemize} We find no clear evidence for a significant term for a [Z] correlation with the period, luminosity and colour but the PLC relation is within the errors identical with that of the $\\delta$ Scuti type stars. We note that for all but one of the investigated pulsators, high-degree nonradial modes were detected spectroscopically (Bohlender et al. 1999), which represents excellent agreement with our work. The spectral variability of the \\LB stars is very similar to that seen in rapidly rotating $\\delta$ Scuti stars (Kennelly et al. 1992).  \\begin{figure}[t] \\begin{center} \\epsfxsize = 70mm \\epsffile{ms2573f8.eps} \\caption{Correlation of [Z] for the \\LB (filled circles) and selected $\\delta$ Scuti type (open triangles) stars.} \\label{z_all} \\end{center} \\end{figure}"
    },
    "0207/astro-ph0207047.txt": {
        "abstract": "We present a method for measuring the cosmic matter budget without assumptions about speculative Early Universe physics, and for measuring the primordial power spectrum $\\Pstar(k)$ non-parametrically, either by combining CMB and LSS information or by using CMB polarization. Our method complements currently fashionable ``black box'' cosmological parameter analysis, constraining cosmological models in a more physically intuitive fashion by mapping measurements of CMB, weak lensing and cluster abundance into $k$-space, where they can be directly compared with each other and with galaxy and Ly$\\alpha$ forest clustering. Including the new CBI results, we find that CMB measurements of $P(k)$ overlap with those from 2dF galaxy clustering by over an order of magnitude in scale, and even overlap with weak lensing measurements. We describe how our approach can be used to raise the ambition level beyond cosmological parameter fitting as data improves, testing rather than assuming the underlying physics. ",
        "introduction": "% WHERE DO WE GO FROM HERE? What next? An avalanche of measurements have now lent support to a cosmological ``concordance model'' whose free parameters have been approximately measured, tentatively answering many of the key questions posed in past papers. Yet the data avalanche is showing no sign of abating, with spectacular new measurements of the cosmic microwave background (CMB), galaxy clustering, Lyman $\\alpha$ forest (Ly$\\alpha$F) clustering and weak lensing expected in coming years. It is evident that many scientists, despite putting on a brave face, wonder why they should care about all this new data if they already know the basic answer. The awesome statistical power of this new data can be used in two ways: \\begin{enumerate} \\item To measure the cosmological parameters of the concordance model (or a replacement model if it fails) to additional decimal places \\item To test rather than assume the underlying physics \\end{enumerate} % TESTING ASSUMPTIONS IS MORE FUN THAN DECIMALS. This paper is focused on the second approach, which has received less attention than the first in recent years. As we all know, cosmology is littered with ``precision'' measurements that came and went. David Schramm used to hail Bishop Ussher's calculation that the Universe was created 4003 b.c.e. as a fine example --- small statistical errors but potentially large systematic errors. A striking conclusion from comparing recent parameter estimation papers (say \\cite{9par,10par,concordance,consistent} by the authors for methodologically uniform sample) is that the quoted error bars have not really become smaller, merely more believable. For instance, a confidence interval for the dark energy density that would be quoted three years ago by assuming that four disparate data sets were all correct \\cite{9par} can now be derived from CMB + LSS power spectra alone \\cite{consistent,Efstathiou02,Melchiorri02,Lewis02} and independently from CMB + SN 1a as a cross-check.   \\begin{figure}[tb] %\\vskip-1.2cm \\centerline{\\epsfxsize=9.0cm\\epsffile{pplot.ps}} \\vskip-0.8cm \\smallskip \\caption{\\label{pplotFig}\\footnotesize% Measurements of the linear matter power spectrum $P(k)$ computed as described in the text, using the concordance model of \\protect\\cite{Efstathiou02} (solid curve) to compute window functions. The locations of the CMB points depend on the matter budget and scales with the reionization optical depth as $e^{2\\tau}$ for $k\\simgt 0.002$. Correcting for bias shifts the 2dF galaxy points \\protect\\cite{2df} vertically ($b=1.3$ assumed here) and should perhaps blue-tilt them slightly. The cluster point scales vertically as $(\\Omega_m/0.3)^{-1.2}$, and its error bars reflects the spread in the literature. The lensing points are based on \\protect\\cite{Hoekstra02}. The Ly$\\alpha$F points are from a reanalysis \\protect\\cite{GnedinHamilton01} of \\protect\\cite{Croft00} and have an overall calibration uncertainty around 17\\%. } \\end{figure}  % MOST COSMOLOGICAL PARAMETERS CAN BE REPLACED BY JUST TWO FUNCTIONS This paper aims to extend this trend, showing how measurements can be combined to raise the ambition level beyond simple parameter fitting, testing rather than assuming the underlying physics. Many of the dozen or so currently fashionable cosmological parameters merely parametrize two cosmological functions \\cite{gravity,spacetime}: the cosmic expansion history $a(t)$ and the cosmic clustering history $P(k,z)$, the observables corresponding to 0th and 1st order cosmic perturbation theory, respectively. %For example, it is often assumed that $(\\d lna/dz)^{-3}$ is a quartic polynomial %whose coefficients are the densities of various This means that non-parametric measurements of these cosmological functions allows testing whether the assumptions associated with the cosmological parameters are in fact valid. Moreover, if there are discrepancies, comparing measurements of these functions from different data sets reveals whether the blame lies with theory, data or both.  % AND WE'RE GOING AFTER P: We will limit our treatment to the 1st order function, $P(k,z)$, since the 0th order function has been extensively discussed previously \\cite{gravity,spacetime,WangGarnavich01}. One of the key ideas of this paper is summarized in \\fig{pplotFig}, showing how CMB, LSS, clusters, weak lensing and Ly$\\alpha$F all constrain $P(k,z)$ at $z=0$. The first plots that we are aware of showing CMB in k-space go back a decade \\cite{ScottWhiteSilk95}, when CMB merely probed scales much larger than accessible to large-scale structure measurements. Since then, CMB has gradually pushed to smaller scales with improved angular resolution while LSS has pushed to larger scales with deeper galaxy surveys. What is particularly exciting now, and makes this paper timely, is that the two have met and overlapped, especially with the CBI experiment \\cite{CBI} and the 2dF \\cite{Colless01} and SDSS \\cite{York00} redshift surveys. \\Fig{pplotFig} shows that CMB now overlaps also with the scales probed by cluster abundance and even, partly, with weak lensing.  % SEPARATING THE EARLY UNIVERSE FROM THE LATE UNIVERSE $P(k,z)$ can be factored as the product of a primordial power spectrum $\\Pstar(k)$ and a transfer function, corresponding to the physics of the Early Universe and the Late Universe, respectively\\footnote{We will assume that the primordial fluctuations are adiabatic, discussing the most general case in \\sec{DiscSec}.}. The two involve completely separate physical processes and assumptions that need to be tested, and the purpose of our method is to measure these two factors separately using observational data. Given a handful of cosmological parameters specifying the cosmic matter budget and the reionization epoch, the transfer function can be computed from first principles using well-tested physics (linearized gravity and plasma physics at temperatures similar to those at the solar surface). The primordial power spectrum is on shakier ground, generally believed to have been created in the Early Universe at an energy scale never observed and involving speculative new physical entities. Most work has parametrized this function as a power law $\\Pstar(k)\\propto k^n$ or a logarithmic parabola $\\Pstar(k)\\propto k^{n+\\alpha \\ln k}$, inspired by the slow-roll approximation in inflationary models \\cite{LiddleBook}, usually with $\\alpha=0$. More general parametrizations have included broken power laws \\cite{Amendola95,Kates95,Atrio97,Einasto99} a piecewise constant function \\cite{Wang99} and other forms \\cite{Kinney01,Matsumiya02}. It has also been shown \\cite{Wang99} that the MAP CMB data \\cite{MAP} in combination with SDSS power spectrum measurements should be able to constrain the shape of $\\Pstar(k)$ in considerable detail. The key challenge is breaking the degeneracy between the two factors, $\\Pstar(k)$ and the transfer function. Although a future brute-force likelihood analysis parametrizing $\\Pstar(k)$ with, say, 20 parameters would be interesting and perfectly valid, it would obscure the simplicity of the underlying physics. Such a ``black box'' approach would entail computing many different curves for each point in parameter space (such as $C_\\l$ for CMB, $P(k)$ for galaxies, the aperture mass function $\\Mapp(\\theta)$ for lensing and the cluster mass function), and mapping out the 20-dimensional likelihood function numerically by marginalizing over other cosmological parameters like those of the matter budget. This would be overkill, since (modulo nonlinearity complications treated below) all measurements shown in \\fig{pplotFig} can be recast directly as weighted averages of $\\Pstar(k)$. % For lensing and clusters, certainly better to fit to data directly. % Map theory to data space rather than vice versa - often better in practice.  The rest of this paper is organized as follows. In \\sec{Psec}, we describe the construction of \\fig{pplotFig}, explaining how CMB, weak lensing and cluster abundance measurements can be mapped into (linear) $k$-space. In \\sec{DegenSec}, we turn to the degeneracy between $\\Pstar(k)$ and cosmological parameters such as the various matter densities, and present our method for breaking it. We show how this allows measuring the cosmic matter budget without assuming anything about $\\Pstar(k)$ and obtaining a non-parametric measurement of $\\Pstar(k)$.    \\def\\pk{P(k)}  % If I put this directly in the title, k gets capitalized... ",
        "conclusions": "\\label{DiscSec}  We have presented a method which complements the traditional ``black box'' likelihood approach to cosmological parameter estimation in two ways: by testing underlying physical assumptions and by improving physical intuition for where the constraints come from. We described how CMB, galaxy, lensing, cluster and Ly$\\alpha$F could be compared directly in (linear) $k$-space in \\sec{Psec}, then showed how a graphical chi-by-eye test could be transformed into a statistically rigorous method in \\sec{DegenSec}, providing independent measurements of Early Universe parameters (the primordial power spectrum $\\Pstar(k)$) and Late Universe parameters ($\\tau$ and the matter budget). We found that requiring consistency between unpolarized CMB measurements and either polarized CMB or large-scale structure data  is quite promising in this regard.  Separating Early and Late Universe physics is particularly timely given the excess in the small-scale CMB power spectrum recently reported by the CBI  team \\cite{Mason02}. We have seen that the angular scales where this excess is seen correspond to spatial scales around $k\\sim 0.2h/\\Mpc$ where the power spectrum is already constrained by galaxy, lensing and cluster observations. This makes it difficult to blame the excess on the Early Universe, say  by tilting or adding a feature in $\\Pstar(k)$. The alternative explanation in terms of contamination from discrete SZ-sources has been shown to be extremely sensitive to the power spectrum normalization on the cluster scale, tentatively requiring $\\sigma_8\\approx 1$ \\cite{Bond02}. With our convention $(\\star)$, $\\sigma_8^2$ probes the range $k\\approx 0.17^{+0.09}_{-0.08}$, which according to \\fig{keffFig} corresponds approximately to the multipole range $\\l=2000\\pm 800$ if the concordance model \\cite{Efstathiou02} is not too far from the truth. % k-range  for sigma8: 0.09351159632......0.1677356809...... 0.2620461881 % l-range  for sigma8: 1262.940674......1989.069092...... 2865.904053 The cluster normalization $\\sigma_R^2$ with $R\\approx 15h^{-1}\\Mpc$ corresponds to $\\l\\approx 1100^{+600}_{-500}$, scales already probed by the shallower (mosaic) CBI observations \\cite{CBI}. % l-range for clusters: 710.5515747......1051.329224...... 1879.100586 The concordance model \\cite{Efstathiou02} normalized to the CMB data gives $\\sigma_8=0.81$ assuming $\\tau=0.05$, and it appears likely that quite accurate and robust $\\sigma_8$-measurements based on the CMB normalization will be available down the road. % My chi-by-eye for the Efstathiou et al model (shown in fig 3 - but need to rerun with tau) gives % sigma8=0.81 (for tau=0.05).  This paper is not intended to the final word on testing physical assumptions in cosmology, merely a small step to be followed by many more. In this spirit, let us close by summarizing some of the most important things that we have not done and some promising directions for future work.  An obvious first step is implementing our method to independently measure the Early and Late Universe parameters. % Although this is way beyond the scope of the present paper, Next year will be an appropriate time to do this, taking advantage of the revolutionary precision that will be offered by MAP. Since this will involve working with a large grid of CMB transfer functions, the approximations described in \\cite{DASH} will be useful for this.  We have made one important assumption throughout this paper that it would desirable to test: that the primordial fluctuations are adiabatic. The adiabatic assumption means that the process generating the fluctuations in the Early Universe created density fluctuations without altering the density ratios of different matter components (photons, baryons, neutrinos, dark matter, \\etc). By tinkering with these relative densities, it is possible to generate a variety of so-called isocurvature fluctuations. In addition to the familiar baryonic and CDM isocurvature modes, there are obscure ones, for instance a neutrino isocurvature mode where the ratio of neutrinos to photons varies spatially but the net density perturbation vanishes, and it can be shown that the most general case corresponds to a function $\\Pstar(k)$ that is not a scalar but a $5\\times 5$ symmetric matrix \\cite{Bucher00}. Although it has been shown that CMB polarization will help constrain isocurvature modes, real progress in this endeavor is likely to be some way off, requiring the sensitivity of the Planck satellite \\cite{Enqvist00,Bucher01,Amendola02}. On the bright side, these complications matter only when studying CMB data, so it is valid to compare and combine the power spectra measured with galaxies, Ly$\\alpha$F, lensing and clusters at low ($z\\simlt 10$, say) redshift as described above assuming purely adiabatic fluctuations.  Staying on the topic of still more general tests, there exist an elegant ``generalized dark matter'' formalism for describing the gravitational effects of the most general matter budget \\cite{HuGDM}, and this is likely to place robust constraints on dark matter and dark energy as power spectrum measurements continue to accumulate over a range of redshifts.  Needless to say, CMB and LSS data will keep improving dramatically in coming years, providing greater sensitivity, $z$-range, $k$-range and $k$-resolution. In addition to the obvious advantages of sensitivity and range, galaxy window functions will become narrower as the SDSS sky coverage improves, improving the ability to detect sharp features in $\\Pstar(k)$. For lensing, the shear power spectrum contribution from a given $z$ has delta function windows on $P(k)$ in the small-angle approximation, so the smearing in $k$-space comes from projection effects and can probably be substantially reduced using photometric redshift and topography techniques. For the Ly$\\alpha$F, the $\\sim 10^5$ SDSS quasar redshifts should greatly improve the measurement errors on larger spatial scales, where one may hope that uncertainties associated with nonlinear physics are smaller. In addition to the CMB and LSS probes we have utilized in this paper, many more appear promising. For instance, the recently claimed detection of galaxy halo substructure \\cite{Dalal02} falls right on our concordance curve in \\fig{pplotFig} but two orders of magnitude further right than the other data points, at $k\\sim 100\\Mpc/h$. Searches for phase space clumpiness in the Milky Way halo with tidal streamers and future space-based astrometry may provide further constraints on dark matter clustering on such small scales, and additional clever observational ideas are undoubtedly waiting to be thought of. The SZ power spectrum is emerging as another promising cosmological observable, sensitive the $k$-range near the cluster scale \\cite{Seljak02}. %  weak+strong lensing map reconstruction can also probe smallish scale. % Mention constraints from reionization? Primordial black holes?  This avalanche of precision data offers an exciting challenge to theorists in the community: to raise the ambition level to making precision cosmology mean more than merely more apocryphal decimal  places, placing our understanding of the Universe on a solid foundation where the underlying physics has been tested rather than assumed.  The authors wish to thank Ang\\'elica de Oliveira-Costa, Arthur Kosowsky, Uros Seljak and David Spergel for helpful comments. This work was supported by NSF grants AST-0071213, AST-0134999, AST-0098606 and PHY-0116590, NASA grants NAG5-9194 \\& NAG5-11099, and two Fellowships from the David and Lucile Packard Foundation. MT is a Cottrell Scholar of Research Corporation.  %%%%%%%%%%%%%%%%%%%%%% REFERENCES: %%%%%%%%%%%%%%%%%%%%%%%%%  %\\clearpage %\\end{multicols}  %\\vskip-1.0cm"
    },
    "0207/astro-ph0207427_arXiv.txt": {
        "abstract": "We present a dynamical analysis of the flow in the jets of the low-luminosity radio galaxy 3C\\,31 based on our earlier geometrical and kinematic model \\citep{LB02} and on estimates of the external pressure and density distributions from {\\em Chandra} observations \\citep{Hard_3C31}. We apply conservation of particles, energy and momentum to derive the variations of pressure and density along the jets and show that there are self-consistent solutions for deceleration by injection of thermal matter.  We initially take the jets to be in pressure equilibrium with the external medium at large distances from the nucleus and the momentum flux to be $\\Pi = \\Phi / c$, where $\\Phi$ is the energy flux; we then progressively relax these constraints. With our initial assumptions, the energy flux is well determined: $\\Phi \\approx$ 9 -- 14 $\\times 10^{36}$\\,W.  We infer that the jets are over-pressured compared with the external medium at the flaring point (1.1\\,kpc from the nucleus) where they start to expand rapidly.  Local minima in the density and pressure and maxima in the mass injection rate and Mach number occur at $\\approx$ 3\\,kpc.  Further out, the jets decelerate smoothly with a Mach number $\\approx$ 1.  The mass injection rate we infer is comparable with that expected from stellar mass loss throughout the cross-section of the jet close to the flaring point, but significantly exceeds it at large distances.  We conclude that entrainment from the galactic atmosphere across the turbulent boundary layer of the jet is the dominant mass input process far from the nucleus, but that stellar mass loss may also contribute near the flaring point.  The occurrence of a significant over-pressure at the flaring point leads us to suggest that it is the site of a stationary shock system, perhaps caused by reconfinement of an initially free jet.  Our results are compatible with a jet consisting of $e^-e^+$ plasma on parsec scales which picks up thermal matter from stellar mass loss to reach the inferred density and mass flux at the flaring point, but we cannot rule out an $e^-p^+$ composition with a low-energy cut-off. ",
        "introduction": "\\label{Introduction}  The measurement of basic flow variables such as velocity, pressure and density in extragalactic radio jets has proved to be an intractable problem, most of the estimates in the literature being highly model-dependent (see \\citealt{Leahy91} for a review).   We have recently shown that the total and polarized emission of the inner jets in the nearby radio galaxy 3C\\,31 can be modelled accurately on the assumption that they are symmetrical, axisymmetric, relativistic, decelerating flows \\citep{LB02}, and we derived a kinematic model for the jet flow that combined longitudinal deceleration and a transverse velocity gradient.   In order to make further progress in understanding jet dynamics, we need a physical model for the deceleration process. Mass loading must occur, but without disruption of the flow. As \\citet{Beg82} first pointed out, a jet can decelerate without being completely decollimated, but only in the presence of an external galactic pressure gradient, which effectively transforms heat back into kinetic energy.  It is not straightforward to estimate the mass input from observations: synchrotron emission gives no direct evidence for the jet composition on kpc scales, and constraints from Faraday rotation are weak. Two principal mechanisms have been proposed for mass loading: \\begin{enumerate} \\item injection from stellar winds within the volume that is traversed by the jet \\citep*{Phi83,Kom94,BLK96}, and \\item entrainment from the galactic atmosphere across an unstable boundary layer, and subsequent communication with the rest of the jet through ingestion of the thermal material and viscous interactions \\citep{Baa80,Beg82,Bic84,Bic86,DeY96,RHCJ99,RH00}. \\end{enumerate} In the remainder of this paper, we will refer to these processes as {\\em internal} and {\\em external} entrainment, respectively.  The majority of theoretical work in the literature concerns non-relativistic jets, but there have been two approaches to the quantitative study of relativistic jet deceleration: through analytical models and simulations \\citep{Kom94,BLK96} and through conservation law analysis \\citep{Bic94}. \\citet{Kom94} considered analytically the case of an electron-positron jet decelerating as a result of internal entrainment and \\citet{BLK96} made numerical simulations of decelerating electron-proton jets. Both of these references assumed that the jet dynamics were dominated by thermal particles (with energies too low to be seen via synchrotron radiation), although some of the cases they considered were hot enough to have a relativistic equation of state. These calculations were not designed to be compared directly with observations of individual objects and are restricted to internal entrainment.  \\citet{Bic94} used the laws of conservation of mass, momentum and energy in a quasi-one-dimensional approximation to demonstrate the feasibility of deceleration from pc to kpc scales for relativistic jets, considering two specific sources: NGC\\,315 and NGC\\,6251. In contrast to \\citet{Kom94} and \\citet{BLK96}, Bicknell assumed that relativistic particles are energetically dominant, and therefore that an ultra-relativistic equation of state is appropriate throughout. His formulation is general enough to cover both internal and external entrainment.  Our model for the jets in 3C\\,31 provides one essential ingredient for a dynamical analysis -- the velocity field -- but we also need to estimate how much mass participates in the flow.  Our solution for the jet kinematics can be used to constrain the mass flux using the conservation-law formalism of \\citet{Bic94}, but only if we also have an accurate prescription for the external pressure and density.  Such a prescription has recently been derived from {\\em Chandra} observations by \\citet{Hard_3C31}, and the present paper describes the resulting dynamical analysis of jet deceleration in 3C\\,31.  The conservation-law approach is described in Section~\\ref{Model}.  The results are presented in Section~\\ref{results} and are discussed in the context of theoretical models in Section~\\ref{discussion}. Our conclusions are summarized in Section~\\ref{conclusions}.  Throughout this paper, we adopt a Hubble constant $H_0$ = 70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  We take the redshift of NGC\\,383 (the parent galaxy of 3C\\,31) to be 0.0169; this is the mean of values from \\citet{Smith2000}, \\citet{HVG} and \\citet{RC3}.  The resulting conversion factor between angular and linear size is 0.34\\,kpc /arcsec.  We refer to two quantities that are conventionally notated as $\\beta$.  We use $\\beta$ alone for the normalized velocity $v/c$, and $\\beta_{\\rm atm}$ for the form parameter in models of hot galactic atmospheres. ",
        "conclusions": "\\label{discussion}  \\subsection{General}  Our analysis shows that the hypothesis that the jets decelerate by entrainment and are recollimated by the external pressure gradient is quantitatively consistent with our model velocity field for external gas parameters derived from {\\em Chandra} measurements.  The uniqueness of our solution depends primarily on the assumptions that the jets are in pressure equilibrium with the external medium in the outer region and that the momentum flux $\\Pi = \\Phi/c$.  We have argued that both assumptions are likely to be correct, but have also demonstrated the effects of relaxing them.  In the remainder of this Section, we assume that they hold precisely.  \\subsection{Comparison with numerical simulations} \\label{numsym}  Several groups have made numerical hydrodynamic or magnetohydrodynamic simulations of the effects of internal or external entrainment on jets. \\begin{enumerate} \\item \\citet{DeY96} modelled the development of turbulent eddies and subsequent entrainment, described as ``ingestion'' followed by ``digestion''; \\item \\citet{BLK96} studied the effects of mass input from stars (considered as a continuous mass source) on a two-dimensional, relativistic jet; \\item \\citet{Lok96} and \\citet{Lok97} modelled external entrainment into a non-relativistic jet using 3D hydrodynamical simulations and derived a mass entrainment rate. \\item \\citet{RHCJ99} and \\citet{RH00} investigated the effects of magnetic fields via 3D, MHD simulations of non-relativistic jets. These showed the non-linear development of Kelvin-Helmholtz (KH) instabilities and their role in external entrainment. \\end{enumerate} Although all of these simulations give important insights into the physics of jet deceleration, none can be compared directly with our results. Except for the calculations of \\citet{BLK96}, which deal specifically with mass input from stars and exclude external entrainment, all are non-relativistic.  Without exception, they assume that the jets are initially in pressure equilibrium with the external medium.  This is inconsistent with our inference of a significant over-pressure in the flaring region.  Only \\citet{BLK96} include a realistic galactic atmosphere.  Finally, the jets in the highest-resolution three-dimensional simulations \\citep{RH00} are much denser than we infer ($\\eta = 0.25$, compared with $\\approx$10$^{-5}$).  Nevertheless, comparison of Fig.~\\ref{lindens-ref} with fig.\\,2 of \\citet{RHCJ99} shows some  similarities: the linear density initially grows slowly, then increases rapidly through the flaring region and levels off at the beginning of the outer region.  The three phases are interpreted as the linear, non-linear and saturated stages of the Kelvin-Helmholtz instabilities.  There are obvious differences, however: the simulations do not show an abrupt increase in emissivity at the flaring point, nor do they predict the further rapid increase in entrainment rate in the outer region.  \\subsection{The onset of deceleration}  A common feature of all our acceptable solutions is that the jet becomes significantly over-pressured at the start of the flaring region.  Such a localised region can persist in a steady-state jet \\citep{Leahy91}: it is apparently ``unconfined'', but the fluid passing through it is expanding, and by the time it has expanded, it is further down the jet, and close to pressure equilibrium with the surroundings.  In fact, the rapid expansion in this region causes the pressure to drop abruptly and the jet becomes over-expanded, starts to recollimate and attains pressure equilibrium over roughly a sound crossing distance.  We have also established that the flaring point in 3C\\,31 is a discontinuity at which the jet collimation, emissivity (and perhaps the velocity) change abruptly and we have argued elsewhere that this is a general property of FR\\,I jets \\citep{LPdRF}.  What causes this sudden transition? It has frequently been suggested that it represents the onset of turbulence \\citep{Bic84} or (almost equivalently) the point at which Kelvin-Helmholtz instabilities start to grow non-linearly \\citep{RHCJ99,RH00}.  Our requirement for a significant over-pressure at the flaring point leads us instead to consider the possibility that the flaring point is associated with a stationary shock system. The boundary position is roughly consistent with the expected location of the reconfinement shock formed when the internal pressure of a freely-expanding supersonic jet falls below that of the external medium \\citep{Sand83}.  Our estimates of $p_{\\rm sync}$ for the inner jet (Fig.~\\ref{ext-fig}) are indeed consistent with an over-pressure for the first 0.5\\,kpc (but note that the external pressure might be underestimated; \\citealt{Hard_3C31}).  The shock is expected to occur at a distance \\[ z_{\\rm shock} \\approx \\left ( \\frac{2 \\Phi}{3 \\pi p_{\\rm ext} c} \\right )^{1/2} \\approx \\mbox{0.5 kpc} \\] from the nucleus for a relativistic jet \\citep{Kom94}, in fortuitously good agreement with our results. The  flaring point cannot be the initial reconfinement shock (by definition, the jet recollimates at that point), but the actual shock structure is likely to be more complicated. In the calculations of \\citet{Sand83}, a conical incident shock forms where the jet has become significantly under-pressured. After this shock, the internal pressure is still slightly below the ambient value.  The incident shock is reflected off the jet axis to form a second conical shock, after which the flow is over-pressured and expanding (Fig.~\\ref{reconf}). It is possible that the {\\rm reflected} shock may represent the visible start of the flaring region.  More detailed simulations will be required to ascertain whether the over-pressure is consistent with the values inferred earlier ($p/p_{\\rm ext} \\approx 8$), although \\citet{Falle87} suggests that values of $p / p_{\\rm ext}$ as high as 12.5 are possible in a non-relativistic jet if the oblique shocks are strong.  \\begin{figure} \\epsfxsize=8.5cm \\epsffile{fig21.eps} \\caption{The shock structure for a reconfining jet, after \\citet{Sand83}. The top panel shows a sketch with the edge of the jet (full line), shocks (bold, full lines) and a representative streamline (dotted line) marked.  Note that the vertical scale is expanded for clarity. The bottom panel shows a sketch of internal (full) and external (dotted) pressures against distance from the nucleus for the streamline in the upper panel. The flaring point is set at 1.1\\,kpc from the nucleus, as for 3C\\,31, the external pressure is as used in our models, as is the internal pressure in the flaring region. All other quantities are notional. \\label{reconf}} \\end{figure}   \\subsection{Internal versus external entrainment}  Whilst the occurrence of a reconfinement shock provides a plausible explanation for the over-pressure, it does not by itself explain the rapid increase in mass input.  However: \\begin{enumerate} \\item the increase in expansion rate will naturally lead to a larger mass injection from stars, which will in turn expand the jet still further in a runaway process and \\item the jet is expected to entrain the external medium more efficiently when it becomes transonic. \\end{enumerate} We estimate that internal entrainment from stars is within a factor of two of that required to slow the jet over the first kiloparsec of the flaring region (but note that the assumptions used to estimate the stellar mass input are extremely crude, since they assume that the loss rates inside and outside a jet are identical and the extent to which mass lost from stars mixes with the jet is also poorly known).  A number of lines of evidence suggest, however, that external entrainment becomes dominant further out. Firstly, our observations and kinematic model \\citep{LB02} show directly that there is an appreciable reduction in velocity at the edges of the jet, as expected in external entrainment models.  The {\\it shape} of the transverse velocity profile in our best fit model changes relatively little down the jet as it decelerates, so the profile could just be set close to the nucleus.  An error analysis shows, however, that an evolution from a top-hat velocity profile at the flaring point to a centrally-peaked profile at larger distances would also be consistent with the data. Secondly, the very sharp peak in the entrainment rate at a distance of 3.5\\,kpc from the nucleus is difficult to reproduce with stellar mass input alone: a maximum is indeed expected, but it should be much broader (e.g.\\ Fig~\\ref{ent-ref}).  Thirdly, the appearance of locally low polarization at the edges of the jets in the flaring region \\citep{LB02} requires the addition of radial magnetic field components at the edges of this region, so that the field becomes almost isotropic there. The appearance of such a field component suggests the onset of local radial motions in this region, consistent with the inflow of ambient material into the jet.  This effect occurs at exactly the position of the local maximum of the entrainment rate (e.g.\\ Fig.~\\ref{ent-ref}). Finally, the monotonic increase of entrainment rate at large distances is clearly inconsistent with the fall-off in stellar density.  A general feature of the models of \\citet{Kom94} and \\citet{BLK96}, where only mass-loading from stars is considered, is that the jets are re-accelerated at distances $>$1\\,kpc, becoming significantly supersonic. The reason is that the entrainment rate, which is proportional to the stellar density, decreases rapidly.  Our results, which indicate a continuous deceleration, therefore require additional entrainment, almost independent of the details of the conservation-law analysis.  We conclude that external entrainment across the jet boundary from the galactic atmosphere must be important in decelerating the jets, but that internal entrainment from stars within the jet may play a significant role in the initial phases.  Indeed, once the area of the jet starts to increase, the mass input from stars may slow the flow to the point where entrainment of external material becomes efficient \\citep{Bic94}.  \\subsection{External pressure and density}  The presence of a component of hot gas with a small core radius is essential for the jets to decelerate without disruption. We have confirmed, for example, that no solutions are possible with the large-scale component associated with the NGC\\,383 group alone. The core radius ($r_c =$ 1.2\\,kpc) and the distance of the flaring point from the nucleus (1.1\\,kpc) are almost exactly equal, so that the external pressure gradient is steepest in the flaring region (Fig.~\\ref{ext-fig}). We would expect a significant external pressure gradient to drive the recollimation of any flaring jet.  In one- or two-component beta-models of the type that we have fit, this inevitably requires that the core radius (or one of the core radii) be close to the flaring distance.  \\subsection{Jet composition}  \\subsubsection{Composition at the flaring point}  By comparing our estimate of jet density with the number of radiating particles required to generate the observed synchrotron emissivity, we can constrain the composition of the jet.  With the assumptions used earlier to calculate the synchrotron minimum pressure, the number density of radiating particles is: \\begin{equation} n_{\\rm rad} \\approx 60 \\gamma_{\\rm min}^{-2\\alpha} \\end{equation} with $\\alpha = 0.55$ if the power-law spectrum inferred from synchrotron emission observed between 1.4 and 8.4\\,GHz continues to lower energy. This estimate uses the on-axis emissivity inferred for our best-fitting kinematic model, which is close to the mean of the transverse profile.  The range of densities at the flaring point for pressure-matched models is $\\rho \\approx$ 1.5 -- 3.5 $\\times 10^{-27}$\\,kg\\,m$^{-3}$.  This corresponds to $\\gamma_{\\rm min} \\approx$ 50 -- 20 if every radiating electron is associated with a proton.  This is a rough estimate whose uncertainties include: \\begin{enumerate} \\item the derivation of $n_{\\rm rad}$ from the emissivity, which assumes a minimum-pressure condition; \\item the form of the spectrum at low energies, where we cannot observe synchrotron radiation directly; \\item the range of densities derived from different models. \\end{enumerate} If we drop the assumption of pressure balance in the outer region, the constraints on $\\gamma_{\\rm min}$ are relaxed, but the only circumstance in which we can avoid a low-energy cut-off entirely is if the high-momentum-flux solutions are valid.  If, in contrast, the jet consists only of electrons and positrons at the flaring point, then there would have to be a significant excess of low-energy particles above the power-law extrapolation.  We conclude that, although the pressure-matched jets are very light, we cannot exclude any of the following possibilities for their composition at the flaring point: \\begin{enumerate} \\item relativistic electrons with a power-law spectrum with energy index $2\\alpha + 1 = 2.1$ and minimum Lorentz factor $\\gamma_{\\rm min} \\approx$ 20 -- 50, each accompanied by a proton; \\item an electron-positron plasma with some admixture of thermal matter, the latter dominating the density; \\item a pure electron-positron plasma with an excess of particles over the power-law prediction at low energies. \\end{enumerate} Other intermediate compositions are possible.  \\subsubsection{Entrainment in the inner jet}  Given that stellar mass input must occur in the inner jet, it is of interest to estimate the mass flux at the flaring point due to this effect alone.  This depends on knowledge of the luminosity density of the galaxy at small radii, which is not available directly for NGC\\,383 (HST optical images show heavy dust obscuration, and infrared observations at sufficiently high resolution are not yet available; \\citealt{Vk}).  Given the stellar luminosity of the galaxy, it is likely that the light profile is of the ``core'' type, in which the surface-brightness profile shows a break from a steep power-law at large radii ($\\Sigma(r) \\propto r^{-1.65}$ for NGC\\,383; \\citealt{Owe89}) to a shallow one ($\\Sigma(r) \\propto r^{-a}$, with 0 \\la $a$ \\la 0.3) at small radii \\citep{Lauer95}.  This transition occurs around a break radius $r_b$ which is correlated with absolute magnitude and is likely to be 100\\,pc \\la $r_b$ \\la 1\\,kpc for NGC\\,383 \\citep{Faber97}.  Given the uncertainties, we have chosen to estimate two extreme limiting cases for the mass input into the inner jet. In the first, we extrapolate the $r^{-1.65}$ surface brightness profile seen at large radii inwards from the flaring point. In the second, we assume a flat profile over the whole of the inner jet, normalized at the flaring point (i.e.\\ $r_b \\approx$ 1\\,kpc and $a = 0$).  The number of particles injected per unit volume per unit time is a Lorentz invariant, so we can derive the mass flux at the flaring point by integrating the mass input rate (equation~\\ref{mdot-eq}) over the volume of the inner jet.  The results are: $\\Psi \\approx 9.8 \\times 10^{19}$\\,kg\\,s$^{-1}$ for $\\Sigma(r) \\propto r^{-1.65}$ and $\\Psi \\approx 1.7 \\times 10^{19}$\\,kg\\,s$^{-1}$ for a constant surface brightness.  The predicted mass flux is at least commensurate with that estimated at the flaring point (2.8 -- 3.4 $\\times$ 10$^{19}$\\,kg\\,s$^{-1}$ for the pressure-matched models). It is therefore possible that essentially {\\em all} of the mass of the jet comes from stars within $\\approx$1\\,kpc of the nucleus. If the jet consists almost entirely of electron-positron plasma on pc scales, it could still pick up enough mass to be consistent with our estimates on kpc scales.  This argument is not yet conclusive, because of the many uncertainties in estimating the stellar mass input rate, but a jet consisting initially of pair plasma would be entirely compatible with our results.  \\subsubsection{Jet composition on parsec scales}  We expect the amount of thermal material to increase from parsec scales to the flaring point. We  therefore  compare our results with those derived for pc scales using the methods of \\citet{Rey96}.  These authors used VLBI and X-ray observations of M\\,87 to argue that its parsec-scale jet is composed primarily of electron-positron plasma, although they could not exclude an electron-proton jet with a low-energy cut-off.  We have repeated their analysis for 3C\\,31.  An upper limit to the magnetic field strength is derived from the surface-brightness of the self-absorbed core.  For an observing frequency of 4.973~GHz, an angular diameter of $<$0.56~milliarcsec and a flux density of 0.071~Jy for the core \\citep{Lara97}, we deduce $B \\la $ 2.4 $\\times 10^{-4}$\\,T for 3C\\,31 if $\\theta =$ 52.4$^\\circ$.  Consideration of the absorption coefficient at the point where the jet becomes optically thick gives $n_{\\rm rad}B^2 > 0.02\\gamma_{\\rm min}^{-1} D_{\\rm max}^{-2}$ where $n_{\\rm rad}$ (in m$^{-3}$) is the number density of radiating particles, $D_{\\rm max} = 1/\\sin\\theta = 1.26$ is the maximum Doppler factor and $B$ is in T. Consequently, $n_{\\rm rad}B^2 \\ga 0.0125$ for $\\gamma_{\\rm min} = 1$.  We use the value of the kinetic luminosity estimated earlier for the reference model ($\\Phi = 1.1 \\times 10^{37}$\\,W) to solve for the particle number density assuming e$^-$e$^+$ or e$^-$p$^+$ jets.  For a bulk Lorentz factor of 3, as assumed by \\citet{Rey96}, we derive $n_{\\rm rad} \\approx 7.5 \\times 10^7$\\,m$^{-3}$ for a pure e$^-$e$^+$ jet and $n_{\\rm rad} \\approx 8.1 \\times 10^5$\\,m$^{-3}$ for an e$^-$p$^+$ jet (note that $n_{\\rm rad}$ consistently includes all radiating species).\\footnote{In \\citet{Rey96}'s discussion of e$^-$p$^+$ jets, the quantity $n$ is used in different places for the number densities of all radiating particles, and for electrons alone.}  Finally, we deduce a lower limit to the magnetic field, $B \\ga 2.9 \\times 10^{-5}$\\,T using the X-ray core flux density at 1\\,keV from \\citet{Hard_3C31} as an upper limit to the synchrotron self-Compton emission.  The constraints are plotted in Fig.~\\ref{Rey}.  \\begin{figure} \\epsfxsize=8.5cm \\epsffile{fig22.eps} \\caption{Constraints on the $B$ -- $n_{\\rm rad}$ plane imposed by synchrotron self-absorption, total kinetic luminosity and synchrotron self-Compton constraints, as in fig.\\,1 of \\citet{Rey96}.  Hatched areas are excluded, as described in the text.  The two bold horizontal lines represent the e$^-$e$^+$ and e$^-$p$^+$ cases for $\\gamma_{\\rm min} = 1$. \\label{Rey} } \\end{figure}  The conclusions from this analysis are slightly weaker than those of \\citet{Rey96} for M\\,87.  For 3C\\,31, an e$^-$p$^+$ jet with $\\gamma_{\\rm min} = 1$ just satisfies the constraints for a bulk Lorentz factor $\\Gamma = 3$, whereas it was formally ruled out for M\\,87. Our values for the jet composition and $\\gamma_{\\rm min}$ at the flaring point would be consistent with the constraints shown in Fig~\\ref{Rey} even in the absence of any changes along the jet.    If the momentum flux is allowed to exceed $\\Phi/c$ by a large factor, the solutions are much less well constrained (Section~\\ref{mom-flux-var}). The jets can be much denser, and entrain more rapidly than those having $\\Pi = \\Phi/c$ (Fig.~\\ref{consfitb}).  Although we cannot rule out these solutions from our data alone, they are incompatible with the need to decelerate from high Lorentz factors on parsec scales and require extremely high entrainment rates even where the jet Mach number ${\\cal M} \\approx 5$.  We suggest that these solutions are highly unlikely.  \\subsubsection{The deceleration mechanism}  The large over-pressure at the beginning of the flaring region suggests the presence of a stationary shock, perhaps associated with reconfinement of the jet.  The amount of mass lost by stars inside the jets and the degree of mixing of the ejecta are both very uncertain, but our best estimate is that stellar mass input is within a factor of two of the rate needed to slow the jet at the beginning of the flaring region. At larger distances, the required entrainment rate is much higher than could be supplied by stars and also increases with distance from the nucleus in a region where the stellar density falls rapidly.  We conclude that another mass source (presumably entrainment from the large-scale galactic atmosphere across the boundary layer of the jet) must dominate at large distances; perhaps everywhere, but that stellar mass input could still significantly affect the initial deceleration. A second piece of evidence in favour of entrainment of external gas across the jet boundary is the (approximate) isotropy of the magnetic field at the edge of the flaring region \\citep{LB02}, which is most easily interpreted as the effect of disordered motions in a turbulent entraining flow.  \\subsubsection{Jet composition}  Our estimate of stellar mass injection within 1\\,kpc of the nucleus is most consistent with the hypothesis that the jets consist primarily of pair plasma on parsec scales and that most of their mass at the flaring point is in the form of entrained thermal plasma.  A jet consisting entirely of electron-positron plasma at the flaring point would require a very large low-energy excess over a power-law energy spectrum. Given the uncertainties in our estimates, we cannot rule out an electron-proton composition; this would require a minimum Lorentz factor of $\\gamma_{\\rm min} \\approx$ 20 -- 50 for the radiating electrons.  \\subsection{Further work}  \\subsubsection{Observations}  The next step in this work is to carry out kinematic modelling and X-ray observations of other sources and to investigate how the jet behaviour depends on galaxy properties and luminosity.  Particularly important questions include: \\begin{enumerate} \\item Is flaring and recollimation always associated with a steep external pressure gradient? \\item Is there a difference in the entrainment rate for sources whose jets propagate entirely within their radio lobes (presumably much less dense than the external medium) compared with those, like 3C\\,31, where the jets appear to be in direct contact with the hot gas? \\item How does the deceleration process depend on jet power? \\item What is the stellar density close to the nucleus? (This will require infra-red imaging at high spatial resolution). \\item Is there morphological evidence for the reconfinement shock structure we have suggested? \\item What limits can we set on the energy spectrum of the relativistic electrons from low-frequency radio and high-frequency (optical -- X-ray) observations? \\item Can we refine the constraints on pc-scale jet composition by higher-resolution VLBI imaging or measurements of circular polarization \\citep{Ward}? \\end{enumerate}  \\subsubsection{Theory}  Our analysis also poses a number of challenging theoretical problems: \\begin{enumerate} \\item Is it possible to simulate entrainment into a decelerating, relativistic, magnetized jet with the very low density contrast we infer and in a realistic galactic atmosphere? \\item Is the required over-pressure at the flaring point consistent with the shock structure in a reconfining jet? \\item What is the viscosity mechanism? How is momentum transported across the jet?  How can we constrain this using estimates of the velocity profile? \\item Are turbulent velocities significant? What are their effects on energy and momentum transport and magnetic fields? \\item Is an ultra-relativistic equation of state an adequate approximation everywhere? \\item How is the entrained material mixed and heated? \\item Can better estimates be made of the mass input rate from stars inside a jet? \\end{enumerate}"
    },
    "0207/astro-ph0207082_arXiv.txt": {
        "abstract": "{Our recent VLBI observations of the prominent FR\\,II radio galaxy Cygnus\\,A with the EVN and the VLBA reveal a pronounced two-sided jet structure. At 5\\,GHz, we now have 4 epochs from 1986, 1991 (Carilli et al., 1991 \\& 1994), 1996 and 2002 from which we could derive the kinematics of the jet and counter-jet. On the jet side and on mas scales, the jet seems to accelerate from $\\beta_{\\rm app}\\approx 0.1-0.2$ (Krichbaum et al. 1998) at core-separations near 1\\,mas to $\\beta_{\\rm app}\\approx 0.4-0.6$ at $r \\geq 4$\\,mas ($H_0=100$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $q_0=0.5$). For the first time we also measure significant structural variability on the counter-jet side. For this, we derive a motion of $\\beta_{\\rm app}=0.35\\pm0.2$ at $r=9.5$\\,mas. The flat spectrum of the inner region of the counter-jet (free-free absorption) and the frequency dependence of the jet to counter-jet ratio support strong evidence for an obscuring torus in front of the counter-jet (Krichbaum et al. 1998). } ",
        "introduction": "Cygnus\\,A is one of the first and strongest objects detected in the radio (\\cite{bolton}). It is the closest ($z=0.057$) strong FR\\,II radio galaxy and therefore a key object for detailed studies of FR\\,II nuclei. In the radio bands, Cygnus\\,A is characterized by two strong lobes separated by $\\sim$\\,2$^\\prime $ in the sky, and two highly collimated jets connecting the lobes with the core (\\cite{perley}, \\cite{carilli1991}). On VLA scales, the jet is oriented along P.A. $\\sim 285^\\circ$ and the fainter counter-jet along P.A. $\\sim 107^\\circ$. Due to the large inclination of the jet with respect to the observer and correspondingly reduced relativistic effects, Cygnus\\,A is an ideal candidate for detailed studies of its jet physics, which is thought to be similar to those in the more beamed quasars (e.g. Barthel 1989). A detailed description of the arc second structure of Cygnus\\,A is given in Carilli \\& Harris (1996) and Carilli \\& Barthel (1996).  We present and discuss the preliminary results of our new VLBI images of the core region of Cygnus\\,A from an 1.6\\,GHz EVN observation (1998.1), two 4.9\\,GHz VLBI observations with the EVN (1996.8) and the VLBA (2002.0) and an 8.4\\,GHz EVN observation performed simultaneously to the first 4.9\\,GHz observation. ",
        "conclusions": "Although our results are still preliminary, the spectral properties give a detailed view on the physics at the core of Cygnus\\,A. We conclude that the small jet to counter-jet ratio and correspondingly the large inclination with respect to the observer, which reduces relativistic effects, contradict the observed difference in the jet and counter-jet apparent velocities and question the applicability of simple jet-models, in which straight and intrinsically symmetrical jets are assumed. In the near future we will have two more observations spaced by only 6 months and we also proposed phase-referencing observation to solve the question of the location of the true nucleus."
    },
    "0207/astro-ph0207647_arXiv.txt": {
        "abstract": "{  High-energy cosmic ray air-showers have been known for over 30 years to emit strong radio pulses in the regime between a few and several 100~MHz (Allan \\cite{Allan}). To date, however, a thorough analysis of the emission mechanisms has not yet been conducted. Adopting a simplified shower geometry and electron-positron energy distribution, we calculate theoretical pulse spectra in the scheme of synchrotron emission from highly relativistic e$^{\\pm}$ pairs gyrating in the Earth's magnetic field. These calculations will play an important role for the calibration of observational data of radio-emission from cosmic ray air-showers acquired with LOPES and later LOFAR and SKA. }  \\authorrunning{T. Huege and H. Falcke} \\titlerunning{Radio-Emission from Cosmic Ray Air-Showers - Theoretical Perspective} ",
        "introduction": "The pulsed radio-emission from high-energy cosmic ray air-showers allows to study their physics with forthcoming digital radio-interferometers such as LOFAR -- an approach offering a number of advantages over other methods (Falcke \\& Gorham \\cite{Falcke}, hereafter FG02). The aim of the LOPES project is to develop and test the necessary hardware, software and techniques for future implementation in LOFAR. Additionally, a thorough understanding of the underlying emission processes is necessary to interpret and calibrate the observational data. Past modeling efforts for radio-emission from cosmic ray air-showers have concentrated on scenarios such as charge-separation and transverse currents induced by the Earth's magnetic field (Kahn \\& Lerche \\cite{Kahn}). An equivalent, but conceptually more attractive and flexible approach is the scenario of coherent synchrotron emission from e$^{-}$ and e$^{+}$ gyrating in the Earth's magnetic field. ",
        "conclusions": "We have shown that known properties of radio-emission from extended cosmic ray air-showers can be successfully reproduced using the approach of coherent synchrotron emission from e$^{\\pm}$ pairs gyrating in the Earth's magnetic field. The comparison with the FG02 approximation and empirical results is encouraging, and the next step will be to incorporate a more realistic geometry for the air-shower. All predictions show that LOPES (or a single LOFAR station) should easily be able to detect the radio pulses of a $10^{17}$~eV shower."
    },
    "0207/astro-ph0207192_arXiv.txt": {
        "abstract": "{Recently, 46 low-luminosity object transits were reported from the Optical Gravitational Lensing Experiment. Our follow-up spectroscopy of the 16 most promising candidates provides a spectral classification of the primary. Together with the radius ratio from the transit measurements, we derived the radii of the low-luminosity companions. This allows to examine the possible sub-stellar nature of these objects. Fourteen of them can be clearly identified as low-mass stars. Two objects, \\object{OGLE-TR-03} and \\object{OGLE-TR-10} have companions with radii of 0.15\\,R$_\\odot$ which is very similar to the radius of the transiting planet HD\\,209458\\,B. The planetary nature of these two objects should therefore be confirmed by dynamical mass determinations. ",
        "introduction": "The detection of planets outside our solar system was a longstanding goal of astronomy. After the first detections (\\citealt{latham:89,wolszczan:92,mayor:95}), an intensive search with various methods began (see \\citealt{schneider:02} for an overview). Out of the currently 102 known planets, 100 have been detected with Doppler velocity measurements of the planets host stars. All these planets were found around solar like stars. The other two are planets around pulsars and were found by periodic pulse modulation measurements.  The Doppler method is subject to several selection effects which are problematic for a more general understanding of planet formation and evolution. It is mainly applied to solar like stars (spectral type F---K) because they provide sufficient lines to measure the radial velocity with the required precision of the order of m/sec. Radial velocity detections favor close-in and massive planets.  Therefore, many Jovian planets are found within Mercury-like orbits.  Regardless of the selection effects, the detection of extra-solar planets has already had a large impact on the understanding and evolution of planetary systems. Establishing a less biased sample would, however, be a big step forward.  No planet has yet been found by photometric monitoring. The (currently) unique planetary companion of HD\\,209458 has an orbital inclination which allows the measurement of the eclipse of the host star by the planet \\citep{charbonneau00,henry:00}. This planetary companion was, however, known before from Doppler measurements \\citep{mazeh:00}. Recently 46 transiting planet candidates were announced by the OGLE (Optical Gravitational Lensing Experiment) consortium \\citep{udalski:02}. These candidates were extracted from a sample of about 5 million stars observed during a 32-day photometric monitoring. In a sub-sample of 52\\,000 stars with a photometric accuracy better than 1.5\\%, these 46 candidates exhibit light curves indicating the presence of a transiting low-luminosity companion. From the analyses of the light curves, the radii of the visible primaries and of the invisible secondaries were derived. Up to now, no spectroscopic information of the primary is available. The goal of this project is to provide this information and to infer the nature of these low-luminosity companions.  We will describe the observations, data reduction and discuss the determination of the spectral types of the primaries in Sect.\\, 2. The results are discussed  in Sect.\\, 3. ",
        "conclusions": "The range of secondary radii is displayed in the Hertzsprung-Russell-Diagram (Fig.\\,\\ref{Fhrd}) together with evolutionary tracks of \\citet{baraffe:98}, \\citet{chabrier:00}, and \\citet{baraffe:02}. The thick lines indicate the (pre-)main sequence evolution of low-mass stars. The position at an age of 5\\,Gyr is indicated in the figure. We also display the evolution of {\\em isolated} contracting brown dwarfs (dashed) and gas giants (dotted). The tracks of the sub-stellar models end at an age of 1\\,Gyr for 0.05\\,M$_\\odot$ and 5\\,Myr for 0.002\\,M$_\\odot$, respectively, and therefore represent very young objects.  \\begin{figure*}[th] \\vspace{10.0cm} \\special{psfile=Ef072_f3a.eps hscale=85 vscale=85 hoffset= -69 voffset=-319 angle=0} \\special{psfile=Ef072_f3b.eps hscale=85 vscale=85 hoffset= 232 voffset= 130 angle=0} \\caption[]{Companion radii compared to evolutionary tracks of \\citet{baraffe:98}, \\citet{chabrier:00}, and \\citet{baraffe:02} in the HRD. Thick lines: stellar models, dashed lines: brown-dwarf models, dotted lines: gas-giant models. Note that the sub-stellar models are for {\\em isolated} objects. Masses and radii are given in solar units. The inset figure shows the mass-radius relation for low-mass stars at an age of 5\\,Gyr.} \\label{Fhrd} \\label{Fmrr} \\end{figure*}  For the following discussion we assume that the OGLE-transits are undisturbed from blends of very nearby stars on the sky and that the transits are no grazing-incident eclipses. Even though these possibilities can not be completely ruled out, the former scenario seems unlikely because we do not detect an additional spectral contribution, the latter one because the photometry indicates flat-bottomed light curves.  All low-mass companions are found to have radii consistent with low-mass stars of about M0V or later \\citep{allen}. For all except two objects our relatively large radii do not allow an interpretation as sub-stellar objects. This list of low-mass star companions includes the best planetary companion candidate, \\object{OGLE-TR-40}, from \\citet{udalski:02}, who derived a companion radius of 0.1\\,R$_\\odot$. Modeling the eclipse light curve, they derived a primary radius of 0.73\\,R$_\\odot$, which can be clearly excluded from our spectroscopic determination. These systems are, however, also interesting. As indicated in Tab.\\,\\ref{Tspec}, the mass ratio for these binary stars is quite extreme. he formation of a close binary out of a common proto-stellar disk favors typically a mass ratio of about unity (e.g. \\citealt{bate:97}). These low-mass objects in eclipsing binaries can also be used to calibrate the mass-radius relation of these stars, providing constraints for evolutionary models. This seems to be required since discrepancies are reported by \\citet{torres:02}.  For two objects, \\object{OGLE-TR-03} and \\object{OGLE-TR-10}, the derived radius of 0.15\\,R$_\\odot$ does allow an interpretation as sub-stellar objects. The latter was also among the two top candidates of \\citet{udalski:02}. In this case our spectroscopic determination fits reasonably well with the light curve fit. In the case of \\object{OGLE-TR-03}, our radius is smaller than the one derived by \\citet{udalski:02}.  Fig.\\,\\ref{Fhrd} shows that sub-stellar objects can be as large as 0.15\\,R$_\\odot$, but only during a very early phase of their evolution, i.e. 0.1\\,Gyr for a 0.05\\,M$_\\odot$ brown dwarf and during 5\\,Myr for a 0.002\\,M$_\\odot$ gas giant. It should be noted that these tracks are calculated for {\\em isolated} sub-stellar objects. The separation of a few solar radii (derived from the orbital period and the assumption that the companion mass is negligible) does indicate a strong influence of the secondary. Theoretical models for sub-stellar companions taking the irradiation of the primary into account are currently worked on (e.g. \\citealt{burrows:00}) and show that the large radii result from the high residual entropy remaining from the early proximity of a luminous companion. For the presently only known transiting gas giant planet, HD\\,209458B, this effect is indeed observed. The derived radius is about 0.14\\,R$_\\odot$, despite the age of probably several Gyrs. The same is possible for \\object{OGLE-TR-03} and \\object{OGLE-TR-10}. While \\object{OGLE-TR-10} would be nearly a twin of the HD\\,209458 system regarding orbital period, spectral type of the primary, and companion radius, \\object{OGLE-TR-03} would be even more extreme. The orbital period is only 1.18 days resulting in a separation of only 5.4\\,R$_\\odot$. In combination with the earlier spectral type, the irradiation is even more drastic.  In summary, the spectroscopic follow-up of the most promising planetary transit candidates did not result in a clear identification of a new sub-stellar object, moreover most of the candidates could be identified as low-mass stars. Two objects did, however, pass this spectroscopic test and therefore continue to qualify as planetary candidates. The ultimate determination of their nature does require a detailed study of radial velocity variations with very high precision. Dynamical mass determination of the secondaries with less demanding instrumental requirement will provide more insight in the mass-radius relation at the lower end of the main sequence."
    },
    "0207/astro-ph0207471_arXiv.txt": {
        "abstract": "There is mounting observational evidence from {\\it Chandra} for strong interaction between keV gas and AGN in cooling flows. It is now widely accepted that the temperatures of cluster cores are maintained at a level of $\\sim$ 1 keV and that the mass deposition rates are lower than earlier {\\it ROSAT/Einstein} values. Recent theoretical results suggest that thermal conduction can be very efficient even in magnetized plasmas. Motivated by these discoveries, we consider a ``double heating model'' which incorporates the effects of {\\it simultaneous} heating by both the central AGN and thermal conduction from the hot outer layers of clusters. Using hydrodynamical simulations, we demonstrate that there exists a family of solutions that does not suffer from the cooling catastrophe. In these cases, clusters relax to a stable final state, which is characterized by minimum temperatures of order 1 keV and density and temperature profiles consistent with observations. Moreover, the accretion rates are much reduced, thereby reducing the need for excessive mass deposition rates required by the standard cooling flow models. ",
        "introduction": "Radiative cooling of gas in the central regions of galaxy clusters occurs on a timescale much shorter than the Hubble time. The cooling time-scale of this gas increases with distance from the cluster core. In the absence of any heating sources, this implies that the intracluster medium must accrete subsonically toward the center in order to maintain pressure equilibrium with the gas at larger radii. The mass deposition rates predicted by this cooling flow model are very high and range typically from 10 to 1000 solar masses per year. X-ray observations made prior to the launch of the {\\it Chandra Observatory} seemed to be broadly consistent with this picture \\citep{fa94}. Although both the gas temperature and cooling time are observed to decline toward cluster cores, new {\\it Chandra} and {\\it XMM-Newton} observations show a remarkable lack of emission lines from gas at temperatures below $\\sim 1$ keV in the central regions of clusters \\citep{pe01,al01}. Moreover, the cooling mass deposition rates obtained with {\\it Chandra} and {\\it XMM} using spectroscopic methods are $\\sim 10$ times smaller than earlier estimates based on {\\it ROSAT} and {\\it Einstein} observations \\citep{mc01,da01,pe01}. On the other hand, morphological cooling rates give accretion rates vastly exceeding the ones based on spectroscopic methods \\citep{da01}. The strong discrepancy between these results indicates either that the gas is prevented from cooling by some heating process or that it cools without any spectroscopic signatures \\citep{fa01}.\\\\ \\indent In this paper, we consider cooling flow models that incorporate the effects of heating by central active galactic nuclei and thermal conduction from the hot outer layers of clusters. We show that evolving density and temperature profiles can relax to stable final states, which are consistent with X-ray observations. Equilibria are characterized by minimum temperatures of order 1 keV at radii $\\sim$ 1 kpc. The paper is organized as follows. In section 2 we discuss the observational and theoretical motivation for including both AGN heating and thermal conduction. In section 3 we present the details of our model. The results of hydrodynamic simulations are presented in Section 4 and the main results are summarized in section 5. ",
        "conclusions": "We have proposed a new class of time-dependent cooling flow models where cooling is offset by a combination of central AGN heating and thermal conduction from the outer regions. Our models do not require any mass dropout rate distributed throughout the cluster. We showed that it is possible to obtain stable final equilibrium states, which do not suffer from the cooling catastrophe. We have presented a representative model, which reproduces the main features of observed cooling flows, including a floor in the temperature at about 1 keV. We have also found stable models for other parameters, which we will present later. Fine tuning is apparently not required to obtain stable equilibrium solutions. Moreover, stable models are characterized by gas accretion rates that are much smaller than the mass dropout rates predicted by standard cooling flow models.\\\\  We are grateful to Phil Armitage, Fabian Heitsch, Christian Kaiser and Daniel Proga for helpful discussions and the referee for a fast response. This work was supported in part by NSF grant AST--9876887."
    },
    "0207/astro-ph0207537_arXiv.txt": {
        "abstract": "Neutron production via ${}^{4}$He breakup and $p$${\\left( p,\\,n{\\pi }% ^{+}\\right) }$$p$ is considered in the innermost region of an accretion disk surrounding a Kerr Black Hole. These reactions occur in a plasma in Wien equilibrium, where (radiatively produced) pair production equals annihilation. Cooling of the disk is assumed to be due to unsaturated inverse Comptonization of external soft photons and to the energy needed to ignite both nuclear reactions. Assuming matter composition of $90\\%$ Hydrogen and $% 10\\%$ He, it is shown that, close to the border of this region, neutron production is essentially from $^{4}$He breakup. Close to the horizon, the contribution from $p\\,{\\left( p,\\,n{\\pi }^{+}\\right) }\\,p$ to the neutron production is comparable to that from the breakup. It is shown that the viscosity generated by the collisions of the accreting matter with the neutrons, may drive stationary accretion, for accretion rates below a critical value. In this case, solution to the disk equations is double-valued and for both solutions protons overnumber the pairs. It is claimed that these solutions may mimic the states of high and low luminosity observed in Cygnus X-1 and related sources. This would be explained either by the coupling of thermal instability to the peculiar behavior of the viscosity parameter $\\alpha $ with the ion temperature that may intermittently switch accretion off or by the impossibility of a perfect tuning for both thermal and pair equilibrium in the disk, a fact that forces the system to undergo a kind of limit cycle behavior around the upper solution. ",
        "introduction": "Accretion disks are presently thought as the scenario for hydrodynamical flows in close binary systems, galactic nuclei and quasars. Despite this fast growing recognition, models which allow for some particle processes, taking into account the detailed nature of the flow, as well as realistic processes that may act as the source for the required viscosity, are still lacking. Concerning particles processes, most of the theoretical work on the field has focused on the production of positron-electron pairs. Steady state scenarios, under the assumption of production-annihilation equilibrium, have been first tackled by \\citet{bis71}, who found that, if particle-particle dominates pair creation, equilibrium is only possible for lepton temperature below $20m{c}^{2}$, \\citet{poz77}, \\citet{stoe77}, and by \\citet{lian79}, who considered pair production dominated by gamma-gamma interactions, in a plasma under Wien equilibrium, finding multivalued solutions to the disk structure. These results, however, are applicable only to a small region of parameter space or to a subset of important reactions. More general studies have been subsequently carried out by many authors \\citep{lig82,sve82,sve84,zdz84,tak85,lig87,whi89}, which contributed to a better understanding of radiative processes in very hot astrophysical plasmas, opening the way to the understanding of the topology of the disk equations solution, which at fixed $r$ in the ${\\dot{M}}\\,{\\Sigma }$ plane forms an S-shape kink corresponding to its multiple valued nature \\citep{kus90,lian91,bjo91,bjo92,min93,kus92,kus95,kus96}. This multivalueness consists of three branches, two of them being hot (one pair dominated and the other pair deficient) and the third being cool and pair deficient. The hot pair dominated branch is very promising to explain some especial features in some black hole candidates, like the MeV bump of Cygnus X-1, the radio plasmoid bipolar outflow of the jet sources, and the alleged annihilation line features \\citep{lin87,lian98, mir92,mir94,hje95,paul91}. It may happen that, under the conditions prevailing in the innermost parts of some systems, the time needed for the matter to cross these regions is comparable to, or even less than, the time the electrons need to extract energy from the protons, through collisional energy exchange. These regions will grow in the direction perpendicular to the plane of the disk, becoming geometrically thick (swelling of the disk may also occur for systems with Eddington or super Eddington luminosities \\citep{abr80,pac82}, not considered in this paper). From angular momentum transport point of view, collective processes are efficiently operative to generate viscosity (due to the protons), but with a deficient cooling (due to the electrons). This drives protons and electrons out of thermal equilibrium, the ion temperature being much greater and close to the virial one. For a disk around a maximally (synchronously) rotating Kerr black hole the inner radius of the disk may reach the horizon, with the ion temperature approaching the mass of the proton. A two temperature geometrically thicker and optically thin disk model was first developed by \\citet{shap76} to explain the hard X-ray from Cygnus X-1. These authors, however, have not considered radial advection of energy and entropy which, in that situation, are very important and the local approximation breaks down. Advection of energy in optically thin disks was first considered by \\citep{ichi77}, in the context of the bimodal behavior of the X-ray spectra from Cygnus X-1, \\citep{lian80} who considered advection in an optically thin, two-temperature disk, and by \\citep{whi90} who studied advection of energy and pairs in optically thin disks. \\citep{pac81,muc82,abr88,hon91,wal91,chen93}, have unraveled the basic physics of advection in accretion flows with their studies on the properties and structure of optically thick disks. However, only recently, advective cooling in optically thin accretion flows has been recognized thanks to a systematic work by \\citet{nar94,nar95}, \\citet{nar95a,nar95b}, \\citet{abr95}, \\citet{nar95a,nar95b}, \\citet{nar96}, \\citet{las96a,las96b}, and \\citet{chen95}. Neglecting pairs, these authors have studied the topology of the solution in the ${\\dot{M}}\\,{\\Sigma }$ plane, showing the existence of a maximum accretion rate, above which no steady state solution is allowed, and that, below it, there are two optically thin solutions, namely the radiative cooling dominated and the advective cooling dominated states. A general description of flows in accretion disks, taking advection into account and neglecting pairs, was given by \\citet{chen95}. They have shown that, at a given radius, there exist exactly four physically distinct types of accretion disks. Two of these correspond to values of ${\\alpha }$, the viscosity parameter, smaller than a critical value, and the other two, to values of ${\\alpha }$ greater than this critical value. Inclusion of pairs have been considered by \\citet{bjo96}, \\citet{kus96}, \\citet{esin96,esin97}, who have shown that pairs do not modify the topology of the solution, besides being negligible for ${% \\alpha }<1$. It should be remarked that, though the advective model constitutes an improvement as far as the standard Shakura and Sunyaev disk model is considered, till the moment, no global solution exists for the disk \\citep{bjo96}. At this point, we must realize the importance of the viscosity parameter and, yet, we don't know the physical process that may be the source for such a viscosity. Since the seminal paper by \\citet{shak73}, a lot of mechanisms have been proposed to account for viscosity: shear turbulence generated by the Keplerian rotation of the disk \\citep{shak73,zel81,dub90,zahn91}, turbulence driven by convection \\citep{lin80,tay80,bis77,shak77,ryu92,cab92,sto96,mei91a,mei93,mei97}, tangled magnetic fields sheared by differential rotation \\citep{lyn69,shak73,ear75,ichi77,cor81}, angular momentum transport by waves \\citep{pap84a,pap84b,vis89,vis90}, Velikov-Chandrasekhar magnetic instability \\citep{bal91,bal92,bal96,vis92}, ion viscosity and neutron viscosity \\citep{gue90,mei93}, radiative viscosity \\citep{loe91}. None of these processes, however, is immune to criticism: concerning shear turbulence, disks satisfy Rayleigh's criterion for stability; convection may transport angular momentum inwards rather than outwards; magnetic field may be removed by magnetic buoyance and ohmic dissipation; concerning wave turbulence, it seems that the most stable mode have low wavenumbers, giving rise to structures of the order of the size of the system, being very sensitive to the Coriolis force which hinders the appearance of smaller structures; the Velikov-Chandrasekhar instability is highly dependent on the radial structure of the azimuthal component of the magnetic field, and can only occur if the Alfv\\'{e}n velocity is of the order of the Keplerian one, or if the typical scale of the magnetic field is smaller than the Keplerian one; as far as ion viscosity is concerned, besides needing high temperatures (the system needs another process to achieve these high temperatures), any small magnetic field present in the flow will decrease the strength of the viscosity by orders of magnitude; neutron viscosity needs high temperature to ignite nuclear reactions capable{\\em \\ }to produce neutrons, needing, therefore, some other process{\\em \\ }to heat up particles till these high temperatures; Keplerian thin disks cannot be supported by radiative viscosity, because the required energy density of the photon gas should be one or two orders of magnitude larger.  In view of all these uncertainties concerning the hydrodynamics of the flow, as well as the angular momentum transport in accretion disks, we would like to consider neutron viscosity, assuming neutron production through Helium break up and through pion production due to proton-proton collisions. The threshold for Helium breakup is about 20 MeV and about 290 MeV for the production of pions (and neutrons) in the proton-proton collision. Close to the horizon, the energy in the rest frame of two colliding protons is about 2 GeV, which means that, at least theoretically, these reactions are energetically viable. The real difficulty is to find conditions in parameter space, if they exist, such that the plasma is no longer ruled by electromagnetic interactions alone and nuclear interactions start playing a role. Effectiveness of strong interactions depends not only on the value of ion temperature itself but also on the electron temperature. This is so because electromagnetic interactions are restricted to the Debye sphere and the number of electrons within it decreases with electron temperature. High electron temperature will favor, not only the production of electron-positron pairs, but also the possibility of nuclear reactions. One has to realize we are here dealing with a very intricate situation highly dependent on the viscosity. It is well known, that for a given density, electron and proton temperatures, drift time decreases with increasing viscosity parameter while $t_{ep}$, the electron-proton collision time, has inverse behavior. However, physical variables in the disk are very sensitive functions of the viscosity parameter and we have to make a self-consistent calculation to find a region in parameter space where the nuclear reactions effects are maximized.  Account for nuclear reactions in the accretion disks may have interesting observational consequences such as the lines spectral features of Hydrogen-Helium plasmas at high temperatures and $\\gamma $-ray lines production. From the theoretical point of view, Helium breakup and pion production reactions produce neutrons whose collisions with the accreting matter may be a source for the viscosity in the disk, which, in turn, will depend upon electron temperature, proton temperature and accretion rate. This may have interesting consequences as far as the topology of the solution is concerned. This is a problem we want to tackle in a future work. However, in the present paper, it should be stressed, our main concern is to find out if neutron production through proton-proton collisions is viable in the innermost region in accretion disks, and if its collisions with the accreting matter can drive accretion on. We will, therefore, be interested only in order of magnitude estimates. Keeping this on mind, we will make some reasonable approximations, mainly on the radiative transport, pair production and on the hydrodynamics, which, at the right moment, will be justified.  A potentially promising application of the afore mentioned reactions is the modelling of some X-ray systems, like Cygnus X-1, that exhibit multimodal behavior. In a previous paper \\citep{mei93}, it was shown that although these reactions, in steady state accretion, can not supply the disk with enough neutrons to make their collisions with the accreting matter the main source of viscosity, they may have strong implications as far as the temporal behavior of the disks is concerned. We have considered, however, only the production of neutrons through the ${}^{4}$He breakup reaction $$ {}^{4}He+28.296 MeV \\rightarrow 2p+2n ,\\eqno(1) $$% for ion temperatures greater than $3$ MeV. Assuming that a steady accretion can be achieved with the drift time equal to the nuclear reaction time, we were able to show the existence of a critical accretion rate, below which there is no steady state accretion onto the hole. Above the critical accretion rate it is possible, under special circumstances, for the disk to choose between the two states of steady accretion.  This kind of procedure has to be criticized on the grounds that equality between drift and nuclear reaction times is a very stringent constraint and, at the temperature range considered, electron-positron pair production should be taken into account. Besides, as we approach the hole, ion temperature can be much greater and even exceed the threshold temperature for the reaction $$ p+p+290 MeV\\rightarrow p+n+\\pi \\,\\,.\\eqno(2) $$  We claim that accounting of both reactions in the inner parts of an accretion disk together with allowance for thermal instability may explain the transitions observed in Cygnus X-1, between the states of high and low luminosity. ",
        "conclusions": "We don't know by sure what physical processes may be operating in the outermost parts of the disk generating viscosity and switching accretion on. Putting this flaw aside, we realize that under certain conditions protons and electrons are out of thermal equilibrium, the ion temperature being much greater and close to the virial one and, for a matter of self-consistency, we are compelled to consider some nuclear reactions, since we have available energy up to twice the proton rest mass. However, besides energy considerations, we had a very strong motivation to consider $^{4}$He breakup and $p\\,{\\left( p\\,n\\,{\\pi }^{+}\\right) }\\,p$. This is related to the viscosity problem itself: these reactions produce neutrons whose collisions with the accreting matter may be a source of the required viscosity to drive accretion. We have considered these reactions in an environment of protons, electrons, neutrons, photons and pairs. Concerning radiative cooling and pair production, we have made some simplifications that overestimated these processes and, as a consequence, underestimated the nuclear reaction rate because the total energy available is constant. The treatment we have adopted to calculate the nuclear reaction rate consists of an improvement over a previous one \\citep{mei93} due to the abandon of the assumption of equality between nuclear reaction and dynamical times, as well as for the inclusion of both pion and pair production. To emphasize the role of the nuclear reactions, as far as the requirement of huge ion temperatures needed to ignite pion production, we have assumed a maximally synchronously rotating Kerr black hole, the inner radius of the disk equal to its horizon, $R_{h}$. The region of the disk we were concerned is the one extending from the horizon to about $20\\,R_{h}$. In the outer parts of the inner region, the contribution of the reaction $p{\\left( p,n{\\pi }^{+}\\right) }\\,p$ to the production of neutrons is negligible, this being dominated by${}^{4}$He breakup, which we use as a kind of boundary condition. Neutrons produced by this reaction contribute to the viscosity and the plasma heats up as we approach the hole, making possible the production of pions which, in turn, increase the viscosity by concomitant production of neutrons.  The radiative cooling of the disk was assumed to be due to unsaturated inverse Comptonization of soft external photons impinging upon the disk. Since we have not considered the reaction $p{\\left( p,p{\\pi }^{0}\\right) }p$% , we have not taken into account the radiative cooling due to photons coming from the ${\\pi }^{0}$ decay. Part of those soft photons are upscattered in energy, reaching the Wien peak where they interact with themselves producing pairs. Radiative cooling and pair production decrease the electron temperature, an effect that hinders the efficiency of nuclear reactions, even if the ion temperature is great enough for them to occur. When $T_{e}$ decreases, the number of interacting electrons with one proton increases sharply, due to less shielding. Nuclear forces only act at short distances and the high ion temperature needed to overcome the Coulomb barrier also hinders the occurrence of nuclear events due to very short time in the nuclear range. As a matter of fact, reaction rate due to electromagnetic interactions prevails over the rate due to nuclear interactions if $T_{e}$ satisfies $$ T_{e}<28.56\\,{T_{i}}^{1/3}\\,\\,.\\eqno(43) $$% However, the energy transferred in the (electromagnetic) scattering will be in the KeV range, while for the nuclear reaction it will be in the MeV range. Therefore, the previous inequality changes roughly to $$ T_{e}<0.2856\\,{T_{i}}^{1/3}\\,\\,,\\eqno(44) $$% when one also considers the energy transferred in the event.  Aiming at the application to Cygnus X-1 and related sources, we have set ${% \\dot{M}}_{17}=1$ and ${\\dot{M}}_{17}=0.9$ which are much less than the critical value for ${\\dot{M}}_{17}$ calculated close to the horizon, i.e., \\[ {{\\dot{M}}_{17}}^c=0.183\\,{T_i}^{1.85}\\,.\\, \\] For both values of the accretion rate $T_e$ is larger than the value given by the ineq. (43). As a result, for every solution we have found, $z<1$. This is surprising, since it is known that, in the absence of these reactions, one of the solutions for the two-temperature disk has $z>1$. Another surprising result we have found concerns the sensitivity of the number of solutions to the value of the accretion rate. Calculating at $r=1$ and $r=10$, for ${\\dot{M}}_{17}=1$ we have found the solution is unique. Decreasing to ${\\dot{M}}_{17}=0.9$ there are two solutions. For these values of the accretion rate the neutron abundance in the inner region is fairly high, going from $y_n=1/19$ in the outer border, with total ${}^4$He depletion, to about twice that value close to the horizon. This implies $% \\alpha $ quite large, i.e., large viscosity. This large viscosity, nevertheless, comes along with a high nuclear cooling, being comparable to the radiative cooling.  It is known that the innermost region of the accretion disk is secular and thermally unstable \\citep{shak76}in the absence of nuclear reactions. Taking into account these nuclear reactions, the disk behaves in the same manner at the very onset of these instabilities. However, it is worthwhile observing the behavior of the viscosity parameter with $T_i$, because two scenarios may emerge as far as the time evolution of the inner region is concerned. ${% \\alpha }$ starts growing, reaches a maximum somewhere between $T_i=15$ and $% T_i=20$, then decreases and may become very small, even null, with increasing $T_i$. At that moment, accretion starts switching off and the cooling, that goes with ${\\alpha }^{-2}$, is practically instantaneous. Outside this region, however, some other mechanism for viscosity generation is still operating and accretion, there, keeps going on. As the matter reaches the border of the inner region it gets piled up there. As this piling up goes on, this region will be subject to several instabilities. Once one of these instabilities starts to grow it triggers accretion in the inner region. As the accretion proceeds, the accretion rate decreases till it reaches a value for which there are two steady states for the disk. Which one will the disk choose? At that moment, do both thermal and pair equilibrium hold in that region? A less drastic scenario is the one somehow similar to that proposed by \\citet{kus91}. Physical processes in the disk are characterized by time scales, the one chosen by the system being the least scale. As we have seen, thermal and pair equilibrium only hold under special circumstances. It may happen that due to the nuclear reactions and instabilities in the disk, the system may undergo a kind of limit cycle behavior around the upper solution.  However, to have a better understanding of the time evolution of these systems we should make better treatment of both radiative cooling and pair production and a more detailed stability analysis of the problem as well. That is what we intend to do next, in a future contribution, taking also into account the reaction $p\\,{\\left( p\\,\\,p\\,{\\pi }^{0}\\right) }\\,p$."
    },
    "0207/astro-ph0207298_arXiv.txt": {
        "abstract": "I am currently analyzing the emission line spectra of the $\\sim$600 galaxies from the sample of Cohen \\etal\\ (2000) and Cohen (2001) in the region of the HDF-North with $z<1.5$. A progress report on this effort of the Caltech Faint Galaxy Redshift Survey is presented. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207251_arXiv.txt": {
        "abstract": "We present an analysis of a highly asymmetric cluster merger from a {\\it Chandra} observation of Abell~85. The merger shows significant disruption of the less massive subcluster from ram pressure effects. Nevertheless, a cold core, coincident with the cD galaxy, is observed to persist in the subcluster. We derive dynamical information from the motion of the cold core through the main cluster's ICM. Multiple derivations of the velocity of the core suggest a Mach number of $\\mathcal{M} \\approx 1.4$, or $v \\sim 2150$~km~s$^{-1}$, though with substantial uncertainty. We construct a consistent kinematic model for the merger based on this dynamical analysis. As has been found for other such ``cold fronts,'' conduction appears to be suppressed across the front. Thermal conduction may be suppressed by a magnetic field with a significant component perpendicular to the subcluster's direction of motion. The effect of the merger interaction in creating and shaping the observed radio sources is also discussed. It appears most likely that the radio source is due to distorted and detached lobes from the subcluster cD galaxy, rather than being a radio halo. ",
        "introduction": "\\label{sec:south_intro}  Mergers of clusters of galaxies are highly energetic events, releasing a total kinetic energy of $\\sim 10^{63}$~ergs into the intracluster medium (ICM). When clusters merge, shocks are driven into the ICM, dissipating the kinetic energy of the merger and heating the gas. These shocks also have nonthermal effects, including the generation of turbulence in the ICM and acceleration of charged particles to relativistic, or cosmic ray, energies. Observations with {\\it Chandra} of merging clusters have provided new insights into the cluster merger process, including the unpredicted discovery of the persistence of cold cores from pre-merger cooling flows well into the lifetime of a merger \\citep[``cold fronts:''][]{mpn+00,vmm01b}.  Abell~85 is in the early stages of merging with two subclusters, each much less massive than the main cluster. One subcluster is merging from the southwest while the other subcluster is merging from the south. The south subcluster will be the focus of our discussion here, while the other subcluster and its associated radio relic will be discussed in a later paper. Abell~85 is unusual in being one of the few clusters known to be in the process of a merger while maintaining a moderate \\citetext{$107~M_\\odot$~yr$^{-1}$; \\citealp{pfe+98}} cooling flow. Presumably, this implies that the merging subclusters have not yet penetrated the inner few hundred kiloparsecs of the cluster and have therefore not yet been able to disrupt the cooling flow.  The south subcluster is more massive than the southwest subcluster. There has been some uncertainty in the past as to whether or not this subcluster is in fact merging with the main cluster or is merely seen against the main cluster in projection. Using data from {\\it ASCA}, \\citet{mfs+98} determined that the temperature in the region of the subcluster is the same as or slightly greater than that of the rest of the main cluster at the same radius. If the subcluster were not merging and were only seen it projection, its smaller mass would give it a lower temperature than that of the main cluster. Thus, the higher temperature indicates that the subcluster is almost certainly interacting.  The redshifts of the galaxies in the southern subcluster are slightly larger that those of the main cluster \\citep{bfh+91,dfl+98}. This suggests that the southern subcluster is either a background cluster or that it is slightly in front of the main cluster and its excess redshift comes from its peculiar motion as it falls into the main cluster. Based on the analysis of \\citet{mfs+98} and the observations presented in the present paper, we believe that it is merging with the main cluster. Thus, we will assume that the southern subcluster is at essentially the same distance as the main cluster, and that any difference in their observed redshifts is caused by their relative motion along the line of sight as they merge.  The {\\it Chandra} observation and basic data reduction are discussed in \\S~\\ref{sec:south_data}. The X-ray image is presented in \\S~\\ref{sec:south_xray_image}. In \\S~\\ref{sec:south_xray_spectra}, we analyze the spectra of interesting regions associated with the southern subcluster. The profiles of the X-ray surface brightness and temperature within the subcluster and in the region ahead of the subcluster are extracted in \\S~\\ref{sec:south_xray_profiles}. We discuss the evidence for a merger and X-ray determinations of the merger Mach number in \\S~\\ref{sec:south_hydro}. The pressure increase at the cold front and properties of the bow shock are used to derive the merger velocity in \\S~\\ref{sec:south_hydro_stag} and \\ref{sec:south_hydro_bow}. We construct a consistent kinematic model for the merger in \\S~\\ref{sec:south_kine}. The suppression of conduction across the cold front is discussed briefly in \\S~\\ref{sec:south_trans_conduction}. There have been claims of a possible radio relic in this cluster as well \\citep{bpl98}, which we discuss in \\S~\\ref{sec:south_radio}. Our results are summarized in \\S~\\ref{sec:south_summary}. We assume $H_0 = 50$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0 = 0.5$ throughout this paper. At the cluster redshift of $z = 0.0538$, 1\\arcsec\\ corresponds to 1.43~kpc. All of the errors quoted are at the 90\\% confidence level. ",
        "conclusions": "\\label{sec:south_summary}  Our analysis of the south subcluster in Abell~85 from $\\sim 37$~ksec of {\\it Chandra} data has revealed several interesting features. The most obvious is a confirmation that the subcluster is indeed merging with the the main cluster. The subcluster contains a remnant cold core which has survived the early stages of the merger. It is smaller and more discrete than similar structures found in other clusters such as Abell~2142 \\citep{mpn+00}, Abell~3667 \\citep*{vmm01b}, and RX~J1720.1+2638 \\citep{mmv+01}, and is perhaps more akin to the ``bullet'' in 1E0657-56 \\citep{mgd+02} or the ``tongue'' seen in Abell~133 \\citep{fsk+02}.  Based on the ratio of the pressure at the stagnation point of the cold front to that far upstream, on the standoff distance of a possible bow shock, and on the shock compression from the bow shock, we find a consistent Mach number and velocity for the merger of $\\mathcal{M} \\approx 1.4$ and $v \\approx 2150$~km~s$^{-1}$. By comparing this velocity to the radial velocity of the subcluster relative to that of the main cluster, we have determined that the merger velocity is about $19^\\circ$ from the plane of the sky. We find a consistent kinematic model for the merger in which the subcluster is in front of and falling into the main cluster. This model is consistent with the expected merger velocity if the subcluster and main cluster have fallen towards one another due to gravity from their turn-around distance in the Hubble flow.  The X-ray observations indicate that this is an offset merger with a finite impact parameter and a significant angular momentum. A crude estimate based on our consistent kinematic model suggests an angular momentum parameter of $\\lambda \\sim 0.2$, which is somewhat larger than the median values expected due to tidal torques.  Magnetic fields in the cold core may be responsible for suppressing thermal conduction across the cold front. This would explain the sharpness of the front, which should be smeared out by conduction in the absence of a magnetic field. The magnetic fields in the cold core may be high as a result of a cooling flow or the AGN located in the central cD galaxy.  We confirm the assertion that the diffuse radio structure in the subcluster is not a cluster radio halo or relic, but is more likely to be a tailed galaxy with a weak or dead nucleus. We also show that its morphology has been shaped by ram pressure in the merger interaction."
    },
    "0207/astro-ph0207584_arXiv.txt": {
        "abstract": "In the standard paradigm for cosmological structure formation, clustering develops from initially random-phase (Gaussian) density fluctuations in the early Universe by a process of gravitational instability. The later, non-linear stages of this process involve Fourier mode-mode interactions that result in a complex pattern of non-random phases. We present a novel mapping technique that reveals mode coupling induced by this form of nonlinear interaction and allows it to be quantified statistically. The phase mapping technique circumvents the difficulty of the circular characteristic of $\\phi_{\\bi k}$ and illustrates the statistical significance of phase difference at the same time. This generalized method on phases allows us to detect weak coupling of phases on any $\\Delta{\\bi k}$ scales. ",
        "introduction": "The morphology of the large-scale structure in the Universe is that of a complex hierarchy of nodes, filaments and sheets interlocking large voids. The Fourier-space description of such a pattern is dominated by the properties of the phases rather than the amplitudes of the Fourier modes \\cite{c3}. According to the prevailing theoretical ideas this pattern developed by a process of gravitational instability from an amorphous pattern of density fluctuations characterized by a Gaussian field with random phases. Since the non-random phases of the present structure have grown from random-phase initial perturbations then there is strong motivation for understanding how phase information develops within this paradigm and to construct a statistical description of galaxy clustering that could be used as a test of the basic idea.   Unfortunately, quantifying the properties of Fourier phases is difficult for a number of technical reasons, so their use in statistical studies has so far been limited. Ryden \\& Gramann (1991), Soda \\& Suto (1992) and Jain \\& Bertschinger (1996) focused on the evolution of individual phases away from their initial values but since the initial phases are unknown these studies can not be used as the basis of a statistical descriptor. The pattern of association between phases is subtle and hard to visualize which makes a statistical test hard to construct {\\it a priori}.   As the first step in a different approach towards quantifying phase information, Coles \\& Chiang (2000) proposed a colour representation method to visualize phase coupling that at least reveals qualitatively how phase information arises during the evolution of $N$-body experiments but does not in itself constitute a statistical descriptor. In a related study, Chiang \\& Coles (2000) quantified phase information using a statistic derived from the Shannon entropy of the distribution of successive phase differences. This study displayed interesting relationships between phase entropy and gravitational clustering but still did not provide a general statistical description.   In this paper we use a generalization of the concept of a return map \\cite{may,cc1} to transfer the phases of different Fourier modes on to a bounded square upon which simple statistical tests can be applied. In this way, we build upon the earlier studies \\cite{cc1,cc2} to construct a method that allows us to transform the phase information in a clustering pattern into a more useful form. ",
        "conclusions": "We have generalized a method based on phase mapping on the return map. This simple, easy-to-implement method can detect phase coupling at any scales $\\Delta {\\bi k}$ in  $k$ space.  We apply this method to two-dimensional simulations of gravitational clustering and the result has shown that even when the evolution is in the mild non-linear regime, phase coupling on certain scale is revealed through the $\\meanchi$ statistics on the $(m,n)$ plane.   In contrast to other methods, such as the Shannon entropy of the distribution of neighbouring phase differences \\cite{cc1}, this method does not require large number of phases. Moreover, this approach can detect the {\\it scale} of phase coupling through the phase mapping as shown in Fig.~\\ref{mapping}.   With the systematic $N$-body simulations shown in Fig.~\\ref{simulation}, we have also demonstrated in Fig.~\\ref{maxchi23}-\\ref{maxchi67} that the scale of phase coupling differs according to the clustering morphology: modes between {\\it small} $\\Delta{\\bi k}$ for large-scale filaments, {\\it large} for small clumps.   This method reveals a signature of non-linear gravitational instability, but also offers the opportunity to provide a general test of Gaussianity that could be applied to cosmic microwave background temperature maps. In future work we shall evaluate the effectiveness for such method."
    },
    "0207/astro-ph0207067_arXiv.txt": {
        "abstract": "X-ray observations have shown that the chemical abundance in the starburst galaxy M82 is quite rich in Si and S compared with oxygen. Such an abundance pattern cannot be explained with any combination of conventional Type I and II supernova yields. Also the energy to heavy element mass ratio of the observed hot plasma is much higher than the value resulted from normal supernovae. We calculate explosive nucleosynthesis in core-collapse hypernovae and show that the abundance pattern and the large ratio between the energy and the heavy element mass can be explained with the hypernova nucleosynthesis. Such hypernova explosions are expected to occur for stars more massive than $\\gsim 20-25 M_\\odot$, and likely dominating the starburst, because the age after the starburst in M82 is estimated to be as short as $\\sim 10^6 - 10^7$ yr. We also investigate pair-instability supernovae ($\\sim 150-300 M_\\odot$) and conclude that the energy to heavy element mass ratio in these supernovae is too small to explain the observation. ",
        "introduction": "M82 is the most active nearby starburst galaxy.  Recently several exciting discoveries have been made for M82.  X-ray observations have revealed the presence of intermediate mass black holes with masses $10^3-10^6$ M$_{\\odot}$ in M82 (Matsumoto et al. 2001; Kaaret et al. 2001), whose locations were found to coincide with the star clusters by SUBARU observations (Harashima et al. 2002).  Also the very energetic expanding molecular super-bubble has been discovered (Matsushita et al. 2000).  These findings have created lots of interest in the formation of black holes of various masses and the evolution of star burst galaxies (e.g., Ebisuzaki et al. 2001).  The critically important information for understanding the evolution of star burst galaxies is the chemical abundances.  The abundance information has also been provided by X-ray observations. Tsuru et al. (1997) observed M82 with ASCA in the $0.5-10$ keV X-ray band and found that its spectrum can be fit with a three component model: a point-like hard component and extended soft-medium components. From the observed emission lines they also obtained abundances of O, Ne, Mg, Si, S and Fe, and found that the abundance pattern is peculiar: Si and S are much abundant than O and Fe compared with the solar ratio. Also the O/Fe ratio is almost solar. Tsuru et al. (1997) concluded that this abundance pattern cannot be reproduced with any combination of the previous Type Iabc and Type II supernova (SN Iabc and SN II) yields (Nomoto, Thielemann \\& Yokoi 1984; Woosley \\& Weaver 1995; Thielemann, Nomoto \\& Hashimoto 1996; Nomoto et al. 1997a,b). They also discussed that the energy to heavy element mass ratio of the observed hot plasma is too large to be of the supernova origin.  We should note that Tsuru et al. (1997) assumed that all stars above 8$M_\\odot$ have already exploded. However, the age of the starburst is estimated to be $\\sim 10^6$yr from the radio observation (Matsushita et al. 2000) and $\\sim 10^7$yr from the size of the X-ray halo observed with ROSAT (Strickland, Ponman \\& Stevens 1997). Only the very massive stars have exploded in such a short time.  Moreover, the previous SN II and SN Ibc yields adopted in Tsuru et al. (1997) were obtained only for the explosion energy of $10^{51}$ erg. Recently some massive supernovae have been found to explode much more energetically as ``hypernovae'' than normal SNe II (e.g., Iwamoto et al. 1998; Nomoto et al. 2001).  Therefore, we need to reconsider the previous conclusions by Tsuru et al. (1997).  In this paper we re-examine whether the abundance and energetics of M82 are consistent with the supernova models taking account of the starburst age and hypernova explosions.  Because of the short age, the abundance is likely to be determined mainly by relatively massive supernovae. The contributions of high energy explosions would affect the energy to the heavy element mass ratio and the abundance pattern in the galactic winds. Recently we have found that in nucleosynthesis of hypernovae, the Si and S abundances are much enhanced relative to O; and also large Fe/O ratio ([Fe/O] $\\sim 0$) can be realized (Umeda, Nomoto \\& Nakamura 2000; Nakamura et al. 2001). These patterns are consistent with those observed in M82, which has motivated us to calculate detailed nucleosynthesis in massive energetic core-collapse SNe for comparison with the M82 data. We also investigate pair instability supernovae (PISNe) of $150 - 300M_\\odot$ stars, because PISNe also yields relatively abundant Si and S compared with O, and large amount of Fe. ",
        "conclusions": "X-ray observations have shown that the chemical abundance in the starburst galaxy M82 is quite rich in Si and S compared with oxygen. Such an abundance pattern cannot be explained with any combination of conventional Type I and II supernova yields. Also the energy to heavy element mass ratio of the observed hot plasma is much higher than that of normal supernovae. We have calculated explosive nucleosynthesis in core-collapse hypernovae to show that the abundance pattern and the large energy of the hot plasma of M82 can be explained with such hypernovae as ($M$, $E_{51}$) = (25, 10) and (30, 30).  More massive and more energetic core-collapse explosions would also satisfy the observed constraints. We have also investigated pair-instability supernovae ($M \\sim 150-300 M_\\odot$) and conclude that the energy to heavy element mass ratio in these supernovae is too small to explain the observation.  Such ``hypernova'' (energetic core-collapse SN) explosions are expected to occur for stars more massive than $\\gsim 20-25 M_\\odot$ (Mazzali et al. 2002). The upper mass limit of core-collapse SNe is still uncertain. Too massive stars may collapse to black holes without explosion.  The question is how the abundance in M82 can be dominated by such hypernovae. One possible explanation is the age effect. As mentioned in Introduction, the age after the beginning of star-burst is estimated to be $\\sim 10^6 - 10^7$ years. On the other hand, the lifetime of our Pop III stars are 1.40, 1.10, 0.81 and 0.70 $\\times 10^{7}$ years for 15, 20, 25 and 30$M_\\odot$ models, respectively (Umeda et al. 2000).  This is consistent with the assumption that only such massive stars as $M \\gsim 20-25M_\\odot$ have exploded in M82.  There have been some suggestions that hypernovae might be correlated with gamma-ray bursts (GRBs) (e.g., Iwamoto et al. 1998). However, how much fraction of hypernovae is associated with GRBs is still unknown. Paczynski (2001) discussed that the SN rate (including the energetic SNe) is much higher than the GRB rate. On the other hand, we have shown that in order to explain the observed abundance pattern of M82, dominant fractions of massive stars ($M > \\sim 25 M_\\odot$) needs to be hypernovae. This implies that hypernovae are much more frequent than GRBs. If GRBs are associated with hypernovae, it is likely that they occur only for certain special cases of hypernova explosions. For example, the explosion energy needs to exceed a certain value, the explosion needs to be extremely aspherical, i.e., the ejecta needs to be strongly beamed, and both hydrogen and helium envelopes need to be stripped off before the explosion. In fact, GRB980425 associate with SN1998bw is the exceptionally weak GRB (Galama et al. 1998), while SN1998bw is among the most energetic hypernovae so far (Nomoto et al. 2001; Mazzali et al. 2002).  What is the implication of this work to the intermediate mass black holes recently discovered in M82?  One possible scenario for the formation of such black holes is the merging of massive stars in the dense stellar clusters (Portegies et al. 1999). This scenario is consistent with the results of this paper, because such merging increases the number of massive star explosions. However, our results constrain the merging history. The typical mass range for the merged stars cannot be in the range for PISN ($M \\sim 150-300M_\\odot$) to be consistent with the energy to heavy element mass ratios observed in M82.   \\bigskip  This work has been supported in part by the grant-in-Aid for Scientific Research (12640233, 14047206, 14540223) of the Ministry of Education, Science, Culture, Sports, and Technology in Japan."
    },
    "0207/gr-qc0207020_arXiv.txt": {
        "abstract": "In a recent {\\it Letter}, Schnittman and Rasio \\cite{sr} argue that they have ruled out chaos in compact binary systems since they find no positive Lyapunov exponents. In stark contrast, we find that the chaos discovered in the original paper under discussion, J.Levin, PRL, {\\bf 84} 3515 (2000), is confirmed by the presence of positive Lyapunov exponents. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207317_arXiv.txt": {
        "abstract": "The Sloan Digital Sky Survey (SDSS) automatically targeted as a quasar candidate the recently discovered, gravitationally lensed, extremely reddened $z=2.2$ quasar \\pmnj. The SDSS spectrum exhibits \\caii\\ absorption at $z=0.76451$, which we identify as the redshift of a lensing galaxy. {\\em Hubble Space Telescope} imaging shows that components CDE of the system are significantly redder than components A or B and detects faint galaxy emission between D and A+B. The redshift of the dust responsible for the reddening remains unconstrained with current data. However, we outline a model wherein lensing and differential reddening by a $z=0.76451$ galaxy pair can entirely explain this system. ",
        "introduction": "\\label{INTRO}  The Sloan Digital Sky Survey\\footnote{The SDSS Web site is http://www.sdss.org/.} \\markcite{yor00}(SDSS; {York} {et~al.} 2000) is using a drift-scanning imaging camera \\markcite{gun98}({Gunn} {et~al.} 1998) to image 10$^4$\\,deg$^2$ of sky on the SDSS $ugriz$ AB magnitude system \\markcite{fuk96,sdss82,sdss105}({Fukugita} {et~al.} 1996; {Hogg} {et~al.} 2001; {Smith} {et~al.} 2002). Two multi-fiber, double spectrographs on a dedicated 2.5m telescope are used to obtain spectra for $\\sim$10$^6$ galaxies to $r=17.8$ and $\\sim$10$^5$ quasars to $i=19.1$ ($i=20.2$ for $z>3$ candidates).  As discussed in \\markcite{sdssqtarget}{Richards} {et~al.} (2002), quasar candidates are targeted for spectroscopy because they are outliers from the stellar locus or because they are unresolved objects with radio emission detected by the FIRST survey \\markcite{bwh95}({Becker}, {White}, \\& {Helfand} 1995). Due to these inclusive criteria and its area and depth, the SDSS is effective at finding quasars with unusual properties and colors (\\markcite{sdss123}{Hall} {et~al.} 2002).  Quasars heavily reddened by dust form one population of quasars with `unusual' colors, at least for optically selected, magnitude limited samples.  Recently, \\markcite{gre01}{Gregg} {et~al.} (2002), hereafter G02, reported on two heavily reddened quasars from a survey of FIRST radio sources with red counterparts in the Two-Micron All-Sky Survey \\markcite{skr97}({Skrutskie} {et~al.} 1997).  One of these --- \\pmnj\\ --- is gravitationally lensed, as discovered independently by G02 and \\markcite{win02}{Winn} {et~al.} (2002), hereafter W02.  W02 identify six radio components of \\jay, five of which have the same spectral slope (not measurable for the faintest component, F), and two of which (C and E) have lower surface brightnesses than the others. W02 present two lens models.  In one, \\jay\\ is a six-image lens, with C and E differentially broadened by interstellar scattering (\\markcite{jon96}{Jones} {et~al.} 1996). This model requires more than one lensing galaxy and even then might not reproduce the image configuration.  In the second, despite the similar spectral slopes of all the radio components, C and E are a foreground object or objects while A+B and D+F are the double images of a core+jet source. This model requires almost perfect source-lens alignment to match the flux ratios.  However, components A, B and D+F (the D+F separation is only 0\\farcs05) have near-IR (NIR) counterparts with very different flux ratios than in the radio.  More data are clearly needed for a viable lens model.  \\jay\\ was independently targeted as a quasar candidate in the SDSS. Here we investigate this system further using SDSS and {\\em Hubble Space Telescope} ({\\em HST}) data. ",
        "conclusions": ""
    },
    "0207/astro-ph0207135_arXiv.txt": {
        "abstract": "{We have discovered three certain (SAX\\,J1324.5$-$6313, 2S\\, 1711$-$339 and \\SAXII) and two likely (SAX\\,J1818.7$+$1424 and \\SAXI) new thermonuclear X-ray burst sources with the BeppoSAX Wide Field Cameras, and observed a second burst ever from a sixth one (2S\\,0918$-$549). Four of them (excluding 2S\\,1711$-$339 and 2S\\,0918$-$549) are newly detected X-ray sources from which we observed single bursts, but no persistent emission. We observe the first 11 bursts ever from 2S\\,1711$-$339; persistent flux was detected during the first ten bursts, but not around the last burst. A single burst was recently detected from 2S\\,0918$-$549 by Jonker et~al. (2001); we observe a second burst showing radius expansion, from which a distance of 4.2\\,kpc is derived.  According to theory, bursts from very low flux levels should last $\\gtap100$\\,s. Such is indeed the case for the last burst from 2S\\,1711$-$339, the single burst from \\SAXII\\ and the two bursts from 2S\\,0918$-$549, but not for the bursts from SAX\\,J1324.5$-$6313, SAX\\,J1818.7$+$1424 and \\SAXI. The bursts from the latter sources all last $\\sim${20}\\,s. We suggest that SAX\\,J1324.5$-$6313, SAX\\,J1818.7$+$1424, \\SAXII\\ and \\SAXI\\ are members of the recently proposed class of bursters with distinctively low persistent flux levels, and show that the galactic distribution of this class is compatible with that of the standard low-mass X-ray binaries. ",
        "introduction": "About 40\\%\\ of the low mass X-ray binaries in our Galaxy occasionally show (so-called Type\\,I) X-ray bursts, which are thermonuclear flashes due to unstable helium and/or hydrogen burning of matter accreted on a neutron star surface (for a review see e.g. Lewin et al. 1993).  A typical burst shows a fast rise ($\\simeq1$\\,s) and exponential decay, softening during the decay (interpreted as cooling of the neutron star photosphere), and a spectrum which can be well described with black-body radiation.  At the moment, about 70 X-ray bursters are known, approximately 20 of which have been discovered with BeppoSAX (e.g., in 't Zand 2001).  Most X-ray bursters are detected with persistent X-ray flux observed before and after bursts. But some sources have been detected during bursts only, with upper-limits on the persistent emission. These limits vary widely, as they depend on the sensitivity of the instrument used.  For example, the X-ray sources in the globular clusters Terzan 1 and Terzan 5 were detected with the Hakucho satellite during bursts only, but EXOSAT and ROSAT also detected the persistent emission (Makishima et al.\\ 1981, Warwick et al.\\ 1988, Verbunt et al.\\ 1995).  The study of X-ray bursts serves various purposes. First, applying the theory of X-ray bursts to observations provides information about the neutron star (e.g., its radius when the distance is known) and about its companion (e.g., whether the matter transferred from the companion is hydrogen-rich or not).  Second, the bursts  unambiguously decide which of the low-mass X-ray binaries contain a neutron star as opposed to a black hole, and provide (an upper limit to) the distance of the binary, from the condition that its luminosity should be less than the Eddington luminosity.  Third, new low-mass X-ray binaries may be discovered by the detection of bursts, in cases where the persistent flux is too low.  The last two points help in forming a more accurate view of the total number and distribution of low mass X-ray binaries in our Galaxy, and of the fraction that contains a neutron star.  In this paper we describe the observation with the BeppoSAX Wide Field Cameras of  six type I burst sources. Four of these are new X-ray sources with a persistent flux below the detection threshold of the Wide Field Cameras of $\\simeq10^{-10}$ $\\ergcms$. The fifth source is a known X-ray source, from which we detect bursts for the first time.  The sixth source is also a previously known X-ray source, the first burst of which was recently discovered by Jonker et al. (2001); we describe a second burst from this source and use it to determine its distance.  In Sect.\\,2 we describe the observations and data reduction; the results are described in Sect.\\,3. Because of their diversity, the sources are discussed in separate subsections, each of which is accompanied by a sub-subsection in which comparison with other observations are made. Finally, in Sect.\\ 4 we discuss some implications of our results for the theory of bursts (Sect. 4.1) and for various sub-populations of the low-mass X-ray binaries with neutron stars (Sect. 4.2).   \\begin{table*} \\caption{Results of the BeppoSAX Wide Field Cameras observations.  For each burst source the table gives the time of the burst, the position with error $\\delta$, the total exposure time on the source between August 1996 and December 2001 $t_{\\rm tot}$, the exposure time $t_{\\rm exp}$ of the pointing in which the burst was detected, and the hydrogen absorption column $\\nh$, persistent flux $F_{\\rm pers}$ between 2-28 keV, and the distance $d$ derived from the burst peak flux.  For the bursts the table gives the e-folding times $\\tau$ for the total flux, and for the hard and soft energies, where we choose bands 2-x and x-28 keV such that both bands have similar countrates. Spectral fits have been made for counts integrated over $t_{\\rm fit}$; we give the black body temperature $kT_{\\rm bb}$, radius $R$ at distance (limit) $d$, the average flux $F$ in two bands, the bolometric peak flux $F_{\\rm peak}$, and the total burst fluence $E_{\\rm b}$; and the temperature $kT_{\\rm brems}$ and photon index $\\Gamma$ for bremsstrahlung and power law fits, respectively. All fluxes are corrected for absorption. \\label{results}} \\begin{tabular}{l@{\\hspace{0.12cm}}c@{\\hspace{0.23cm}}c@{\\hspace{0.23cm}}c@{\\hspace{0.23cm}}c@{\\hspace{0.23cm}}c@{\\hspace{0.23cm}}c@{\\hspace{0.23cm}}} \\hline & SAX  &SAX & SAX &SAX&2S&2S\\\\ & J1324.5$-$6313 & J1818.7+1424 & J1828.5-1037 & J2224.9$+$5421 & 1711$-$339 & 0918$-$549\\\\ \\hline \\multicolumn{6}{l}{\\bf Source parameters}\\\\ Burst time (MJD) & 50672.151 & 50683.770 & 51988.863 & 51488.454 &b1-11, Table\\,\\ref{1711} & 51335.049\\\\ RA (J2000)  & $13^{\\rm h} 24^{\\rm m} 27^{\\rm s}$ & $18^{\\rm h} 18^{\\rm m} 44^{\\rm s}$ & $18^{\\rm h} 28^{\\rm m} 33^{\\rm s}$ & $22^{\\rm h} 24^{\\rm m} 52^{\\rm s}$ & $17^{\\rm h} 14^{\\rm m} 17^{\\rm s}$ & $09^{\\rm h} 20^{\\rm m} 37^{\\rm s}$\\\\ Dec (J2000) & $-63^\\circ 13\\farcm4$ & $14^\\circ 24\\farcm2$ & $-10^\\circ 37\\farcm8$  & $+54^\\circ 21\\farcm9$ & $-34^\\circ3\\farcm3$  & $-55^\\circ 13\\farcm9$ \\\\ $\\delta$ (99\\% confidence.)       & $1\\farcm8$& $2\\farcm9$ & $2\\farcm8$ & $3\\farcm2$ & $1\\farcm5$& $0\\farcm7$\\\\ $l_{II},b_{II}$ & $306\\fdg6, -0\\fdg6$ & $42\\fdg2,+13\\fdg7$ & $20\\fdg9,0\\fdg2$ & $102\\fdg6,-2\\fdg6$ & $352\\fdg1, +2\\fdg8$ & $275\\fdg9, -3\\fdg8$ \\\\ $t_{\\rm tot}$ (day) & 58 & 31 & 25 & 65  & 66 & 62  \\\\ $t_{\\rm exp}$ (ks) & 18.5 & 16.2 & 12.7 & 40.2 & 314 & 36 \\\\ $\\nh$ ($10^{22}$ $\\cmsq$)  & 1.5$^a$ & 0.1$^a$ & 1.9$^a$ & 0.5$^a$ & 1.5$^b$  & 0.24$^c$ \\\\ $F_{\\rm pers}$ ($10^{-10}$ $~\\ergcms$)& $<0.80$ ($3\\sigma$)& $<1.7$ ($3\\sigma$) & $<1.9$ ($3\\sigma$) & $<0.35$ ($3\\sigma$) &$6.3\\pm0.6$& $3.8\\pm0.6$\\\\ $d$ (kpc)        & $<6.2$  & $<9.4$ & $<6.2$ & $<7.1$ &$<7.5$ & 4.2 \\\\ \\hline \\multicolumn{5}{l}{\\bf Burst parameters}\\\\ $\\tau_{2-28 \\rm keV}$ (s)             & $6.0\\pm0.1$       & $4.5\\pm0.1$ & 11.2$\\pm$0.6 & 2.6$\\pm$0.2 & $7.1\\pm0.2$& $48.5\\pm0.2$\\\\ x     & 8                 & 4           & 7 & 6 & 6          & 5\\\\ $\\tau_{2-x\\rm keV}$  (s)              & $9.7\\pm0.4$       & $5.7\\pm0.2$ & 21.5$\\pm$1.3 &  3.9$\\pm$0.7 & $7.6\\pm0.3$& $80.6\\pm0.7$\\\\ $\\tau_{x-28 \\rm keV}$ (s)             & $2.6\\pm0.2$       & $1.17\\pm0.04$ & 4.7$\\pm$0.6 & 1.8$\\pm$0.3 & $5.9\\pm0.3$& $36.5\\pm0.2$\\\\ \\multicolumn{6}{l}{\\bf black-body fit}\\\\ $t_{\\rm fit}$ (s)    & 9.3    & 6.0  & 25.1 & 4.0 & -          & 86.4\\\\ $kT_{\\rm bb}$ (kev)  & $2.5\\pm0.2$   & $1.1\\pm0.14$ & 2.3$\\pm$0.2 & 2.5$\\pm$0.3  & 1.6$\\pm$0.1 & 2.26$\\pm$0.05\\\\ $R$ (km) at $d$      & $4.5\\pm0.5$   & $24.^{+5}_{-8}$& 4.7$\\pm$0.9 & 4.7$\\pm$0.5 & 5.5$-$11.9 & 6.3$\\pm$0.2\\\\ $F_{2-10 \\rm keV}$ ($10^{-8}$ $~\\ergcms$) & $1.23\\pm0.10$  & $1.02\\pm0.03$ & 1.08$\\pm$0.40& 1.00$\\pm$0.42 & $0.4\\pm0.1^d$ & 4.05$\\pm$0.27\\\\ $F_{2-28 \\rm keV}$ ($10^{-8}$ $~\\ergcms$)& $2.17\\pm0.07$  & $1.07\\pm0.05$ & 1.65$\\pm$0.73 & 1.66$\\pm$0.89 & $0.5\\pm0.1^d$ & 6.1$\\pm$0.5\\\\ $F_{\\rm peak}$ ($10^{-8}$ $~\\ergcms$) & $4.3\\pm0.2$ & $1.9\\pm0.1$ & 4.3$\\pm$1.6 & 3.3$\\pm$1.5 & $3.0\\pm1.0^d$ & $9.4\\pm1.7$\\\\ $E_{\\rm b}$  ($10^{-7}$ $~\\ergcm$) & $\\simeq$2.6  &  $\\simeq$0.86 & $\\simeq$4.3 & $\\simeq$0.67 & -  &  $\\simeq$52 \\\\ $\\chi^2_\\nu$ (d.o.f.) & 1.0 (26) & 0.6 (26) & 0.7(26) & 1.0(26) & 0.9 (270)   & 1.6 (26)\\\\ \\multicolumn{6}{l}{\\bf bremsstrahlung  fit}\\\\ $kT_{\\rm brems}$ (keV)& $56^{+47}_{-34}$& $5.9\\pm1.7$ &  44.$^{+146}_{-22}$ & 50.$^{+149}_{-25}$ & -           &  64$\\pm$14\\\\ $\\chi^2_\\nu$ (d.o.f.) & 1.4 (26)    & 0.5 (26) & 1.1(26) & 1.1(26) & -           & 9.7 (26)\\\\ \\multicolumn{5}{l}{\\bf power law fit}\\\\ $\\Gamma$              & $1.4\\pm0.1$ & $2.3 \\pm 0.21$ & 1.4$\\pm$0.2 & 1.4$\\pm$0.2 & -       & 1.31$\\pm$0.04\\\\ $\\chi^2_\\nu$ (d.o.f.) & 1.5 (26)          & 0.5 (26) & 1.1(26) & 1.1(26) & -           & 10 (26)\\\\ \\hline \\end{tabular} $^a$\\,Interpolated from Dickey \\&\\ Lockman (1990); $^b$From NFI, see Sect. 3.2.1; $^c$\\,From Christian \\& Swank (1997); $^d$Values for b8. \\end{table*} ",
        "conclusions": "Our observations add at least three and possibly five bursters to the list of X-ray bursters in our Galaxy and provide a distance estimate for a recently discovered burster.  In this section we discuss the implications of our results for the theory of bursts (in Sect.\\ 4.1) and make a comparison between the class of bursters with low persistent luminosities -- to which we have added four members -- with other low-mass X-ray binaries (in Sect. 4.2).  \\subsection{Comparison with burst theory}  We compare the properties of the bursts with theory. Fujimoto et\\,al.\\ (1987, see also Bildsten 2000) propose three classes of bursts.  At the lowest accretion rates, $10^{-14}$ M$_\\odot{\\rm yr}^{-1} \\ltap \\dot M \\ltap 2\\times10^{-10}$ M$_\\odot$ yr$^{-1}$, a burst is triggered by thermally unstable hydrogen burning, and can last between 10$^2$ to 10$^4$ s. At intermediate accretion rates, $2\\times10^{-10}\\ltap\\dot M\\ltap10^{-9}$ M$_\\odot$ yr$^{-1}$, a pure helium burst occurs with a duration of order 10 s.  In the high accretion regime, $10^{-9}\\ltap\\dot M\\ltap2.6\\times 10^{-8} $M$_\\odot$ yr$^{-1}$, a burst with a duration of tens of seconds may occur in a mixed He/H environment. At even higher or lower accretion rates no bursts are expected to occur.  To consider our observations, we first converted accretion rates to fractions of the Eddington limit.  The pure helium bursts occur when the accretion rate is in the range 0.014 - 0.070 of the Eddington accretion rate.  The ratio of (the upper limit to) the persistent flux and the peak flux during the burst, where the latter is (a lower limit to) the Eddington flux, provides an estimate of the fraction of the Eddington limit at which a source is accreting. This assumes that the emission is isotropic and that the persistent flux in the range 2-28 keV is close to the bolometric flux. With the values listed in Table\\,\\ref{results} we obtain $<0.002$ for SAX\\,J1324.5$-$6313, $<0.009$ for SAX\\,J1818.7+1424, $<0.004$ for \\SAXII, $<0.001$ for \\SAXI, $0.02$ for 2S\\,1711$-$339 at bursts b1-b10, and $<0.002$ at burst b11, and $0.004$ for 2S\\,0918$-$549.  We thus note that, with the exception of 2S\\,1711$-$339 during bursts b1-b10, all sources are in the low accretion regime, and thus according to theory should emit bursts lasting longer than about 100\\,s.  Source 2S\\,1711$-$339 follows this prediction nicely, showing short ($\\simeq$10-20 s) bursts b1-b10 when accretion was in the intermediate range, and a longer ($\\simeq$60 s) burst b11 when the accretion had dropped to the low regime. Also the burst from \\SAXII\\ lasts about 60 s, pointing towards the low accretion regime. Similarly, the bursts observed by Jonker et al.\\ (2001) and by us for 2S\\,0918$-$549 are long ($\\simeq$150 s), as predicted from the low accretion rate.  In remarkable contrast, the bursts from SAX\\,J1324.5$-$6313, SAX\\,J1818.7+1424 and \\SAXI\\  are all short ($\\simeq$10-20 s), even though these systems appear to be in the low accretion regime. Can it be that the true accretion rate is higher than we estimate? One possibility is that the emission is anisotropic. However, to our knowledge no indication has been found for anisotropies in other burst systems. A second possibility would be that most of the persistent flux is outside the observed 2-28\\,keV range. However, we estimate that more than 50\\% of the persistent flux is in this range. We therefore consider it unlikely that these effects are sufficient to bring especially SAX\\,J1324.5$-$6313 and \\SAXI\\ to the intermediate accretion regime.  A third possibility is that the accretion is limited to a small area of the neutron star, e.g.\\ a ring connected with the accretion disk (Popham \\& Sunyaev 2001).  This enhances the local accretion rate, which is the parameter determining the properties of the bursts (as discussed by Bildsten 2000).  This would imply that the accreting surfaces in SAX\\,J1324.5$-$6313 and \\SAXI\\ are less than about 15\\%\\ and 8\\% of the surface of the neutron star, respectively.  This possibility cannot be excluded {\\em a priori}, but raises the interesting question why the accreting surface areas would be so different between bursters -- the rotation period of the neutron star could affect the area over which the accreted matter spreads out, for example.  A fourth possibility is that the persistent flux at the time of the burst is not representative of the time-averaged flux in the months before the burst. In transients like e.g. Aql~X-1, Cen~X-4, XTE J1709$-$267 and SAX\\,J1750.8$-$2900, X-ray bursts were detected during the decline of the outburst, at times when the persistent flux, easily detectable at $L_{\\rm x}\\gtap10^{36}$ $\\ergs$, was at an accretion rate of ordinary burst sources (Matsuoka et al.\\ 1980, Koyama et al.\\ 1981, Cocchi et al. 1998, Natalucci et al. 1999). Also the transient SAX~J1808.4$-$3658 showed a $\\sim$$100$ s long burst 30 days after the peak of an outburst, when the persistent flux had declined below the detection limit of the Wide Field Cameras, $<10^{36}$ $\\ergs$ (in 't Zand et al.\\ 2001). However, the RXTE/ASM lightcurves show no detection of \\SAXIII\\ and SAX\\,J1818.7+1424, at an upper-limit of $\\simeq10^{36}$ $\\ergs$ making a transient outburst very unlikely.  One might propose the possibility that these systems are old and the companion has only pure helium left. In this case only helium bursts can occur independent of the accretion rate. However, calculations on bursts due to pure helium accretion show that at low accretion rates the burst duration increases to $\\sim100$ s (Bildsten 1995).  For pure helium bursts, the energy released during the burst due to nuclear fusion is about 1\\%\\ of the accretion energy released when the same matter accreted onto the neutron star before the burst (see e.g.\\ Lewin et al.\\ 1993).  From the observed burst fluences and the (upper limits to) the persistent flux, we can therefore derive (lower limits to) the interval to the previous burst.  The computed waiting time of 16\\,d is sufficiently long to explain that only one burst was detected for 2S\\,0918$-$549, whose WFC exposure times totalled for all observations between August 1996 and December 2001 is about 62\\,d. For SAX\\,J1324.5$-$6313, SAX\\,J1818.7+1424, \\SAXII\\ and \\SAXI\\ the total observation times between August 1996 and December 2001 are about 58\\,d, 30\\,d, 25\\,d and 65\\,d, respectively.  For these sources the waiting times are $>3.7$\\,d, $>0.6$\\,d, $>2.6$\\,d and $>2.2$\\,d. The chance probability of observing at most one burst for these sources is then 0.07\\% or (much) less. The fact that only one burst was observed for each system suggests that the persistent emission levels are much lower than the upper-limits derived.  \\subsection{Low persistent emission bursters}  Gotthelf \\&\\ Kulkarni (1997) discovered a burst from a low-luminosity source in the globular cluster M\\,28, with a peak luminosity that is only 0.02\\%\\ of the Eddington limit. This low peak flux discriminates it from the bursters discussed by Cocchi et al. (2001) and in this paper, that have fluxes close to the Eddington limit: if their peak fluxes were as low as that of the M\\,28 source, they would be a local population near the Sun, which is clearly incompatible with their galactic length and latitude distributions.  As discussed by Cocchi et al.\\ (2001), a class of bursters with low persistent emission has emerged in recent years. The four sources discussed in Sect. 3.1 also appear to be member of this class, strenghtening its existence.  Whereas most bursters emit their bursts at persistent luminosities $\\gtap10^{36}$ $\\ergs$, most of the members of this new class emit bursts at luminosities below the RXTE/ASM detection-limit of $\\simeq10^{36}$ $\\ergs$. How much lower is not clear, and we briefly consider three possibilities. One is that the sources are steady in the range $10^{34-35}$ $\\ergs$, as suggested for the bursters 1RXS\\,J171824.2$-$402934 (Kaptein et al.\\ 2000) and \\SAXII\\ (this paper), whose persistent emission levels were detected at this level with ROSAT a few years before the burst. The second possibility is that the sources are steady at the level $10^{32-33}$ $\\ergs$, the quiescent level of soft X-ray transients with neutron stars; and the third possibility is that they are usually at this low level, but emit their bursts during or soon after faint ($\\ltap10^{36}$ $\\ergs$) outbursts, as suggested by the case of 2S\\,1711$-$339.  More sensitive X-ray observations are required to discriminate between these various possibilities.  \\begin{table} \\caption{Overview of the burst sources at low persistent emission as observed with the Wide Field Cameras. \\label{nopersis}} \\begin{tabular}{lc@{\\hspace{0.12cm}}l@{\\hspace{0.12cm}}l@{\\hspace{0.12cm}}c@{\\hspace{0.12cm}}} \\hline Name & $l_{II}$ & $b_{II}$ & $F_{\\rm peak}/F_{\\rm pers}$  & $\\tau$ (s)\\\\ \\hline SAX\\,J1324.5-6313        & $306\\fdg64$ & $-0\\fdg59$ & $>$540   & 6.0   \\\\ RX\\,J171824.2-402934$^a$ & $347\\fdg28$ & $-1\\fdg65$ & $>$90    & 47.5  \\\\ GRS\\,1741.9-2853$^b$     & $359\\fdg96$ & $0\\fdg12$  & $>$130   & 8.8   \\\\ &             &            & $>$180   & 11.0  \\\\ &             &            & $>$100   & 16.0  \\\\ SAX\\,J1752.4-3138$^c$    & $358\\fdg44$ & $-2\\fdg64$ & $>$120   & 21.9  \\\\ SAX\\,J1753.5-2349$^d$    & $5\\fdg30$   & $1\\fdg10$  & $>$180   & 8.9   \\\\ SAX\\,J1806.5-2215$^d$    & $8\\fdg15$   & $-0\\fdg71$ & $>$200   & 4.0   \\\\ &             &            & $>$210   & 9.0   \\\\ \\SAXII                   & $20\\fdg88$  & $+0\\fdg18$ & $>$226   & 11.2  \\\\ SAX\\,J1818.7+1424        & $42\\fdg32$  & $13\\fdg65$ & $>$110   & 4.5   \\\\ \\SAXI                    & $102\\fdg56$ & $-2\\fdg61$ & $>$940   & 2.6   \\\\ \\hline \\multicolumn{5}{l}{$^a$ Kaptein et al. (2000); $^b$ Cocchi et al. (1999).}\\\\ \\multicolumn{5}{l}{$^c$ Cocchi et al. (2001); $^d$ in 't Zand et al. (1998).}\\\\ \\end{tabular} \\end{table}  However, a first test can be made on the basis of the spatial distributions. In Figure\\,\\ref{kstest} we compare the distributions of galactic length and latitude for the bursters with low persistent luminosity -- listed in Table\\,\\ref{nopersis} -- with those of the low-mass X-ray binaries. For the latter we use exponential distributions in galactic longitude and latitude with scale angles of $45^{\\circ}$ and $8\\fdg3$, respectively, as determined by van Paradijs \\&\\ White (1995). Kolmogorov-Smirnov tests indicate that the bursters with low persistent luminosity may indeed be drawn from the distribution of low-mass X-ray binaries.  A new class of faint transients has been discovered with BeppoSAX: transients whose outbursts are rather fainter (peaking below $10^{37}$ $\\ergs$) and often also shorter (lasting less than a month) than the outbursts of the ordinary soft X-ray transients which reach the Eddington limit and may last months (Heise et al. 2000). It is tempting to assume that the bursters at low persistent emission are an extension of this class of faint transients.  Remarkably, this new class of faint transients is more concentrated towards the galactic center than the ordinary low-mass X-ray binaries (In 't Zand 2001).  A Kolmogorov-Smirnov test shows that the galactic longitudes of bursters with low persistent emission cannot be drawn from the longitude distribution of the faint transients, as illustrated in Fig.\\,\\ref{kstest}.  We conclude that the bursters at low persistent emission are probably not from the same class as the faint transients, which makes the transient explanation even more unlikely.  \\begin{figure}[t] \\psfig{figure=H3274.KS.eps,width=9.0cm} \\caption{The top two panels show the cumulative Galactic longitude $l_{II}$ and latitude $b_{II}$ distributions (solid lines) of the low persistent emission bursters compared to the exponential (Galactic) distribution of the low-mass X-ray binaries weighted  with observation times (dashed lines). The probability according to a two-sided Kolmogorov-Smirnov test that they have the same distribution is given in the lower right corners. For comparison we also show an isotropic distribution weighted  with observation times (dotted lines). In the bottom two panels the low persistent emission bursters (solid line) are compared to the faint transients (dashed line). \\label{kstest}} \\end{figure}"
    },
    "0207/astro-ph0207629_arXiv.txt": {
        "abstract": "The primary ultrahigh energy particles which produce giant extensive air showers in the Earth's atmosphere present an intriguing mystery from two points of view: (1) How are these particles produced with such astounding energies, eight orders of magnitude higher than those produced by the best man-made terrestrial accelerators? (2) Since they are most likely extragalactic in origin, how do they reach us from extragalactic distances without suffering the severe losses expected from interactions with the 2.7 K thermal cosmic background photons -- the so-called GZK effect?  The answers to these questions may involve new physics:  violations of special relativity, grand unification theories, and quantum gravity theories involving large extra dimensions. They may involve new astrophysical sources, \"zevatrons\". Or some heretofore totally unknown physics or astrophysics may hold the answer. I will discuss here the mysteries involving the production and extragalactic propagation of ultrahigh energy cosmic rays and some suggested possible solutions. ",
        "introduction": "About once per century per km$^2$ of the Earth's surface, a giant shower of charged particles produced by a primary particle with an energy greater than or equal to 16 joules (100 EeV = $10^{20}$ eV) plows through the Earth's atmosphere. The showers which they produce can be detected by arrays of scintillators on the ground; they also announce their presence by producing a trail of ultraviolet flourescent light, exciting the nitrogen atoms in the atmosphere. The existence of such showers has been known for almost four decades \\cite{li63} (Linsley 1963). The number of giant air showers detected from primaries of energy greater than 100 EeV has grown into the double digits and may grow into the hundreds as new detectors such as the ``Auger'' array and the ``EUSO'' (Extreme Universe Space Observatory) and ``OWL'' (Orbiting Wide-Angle Light Collectors) satellite detectors come on line. These phenomena present an intriguing mystery from two points of view: (1) How are particles produced with such astounding energies, eight orders of magnitude higher than are produced by the best man-made terrestrial accelerators? (2) Since they are most likely extragalactic in origin, how do they reach us from extragalactic distances without exhibiting the predicted cutoff from interactions with the 2.7K cosmic background radiation? In these lectures, I will consider possible solutions to this double mystery. ",
        "conclusions": ""
    },
    "0207/astro-ph0207303_arXiv.txt": {
        "abstract": "During the course of our deep optical imaging survey for Ly$\\alpha$ emitters at $z \\approx 5.7$ in the field around the $z=5.74$ quasar SDSSp J104433.04-012502.2, we have found a candidate strong emission-line source.  Follow-up optical spectroscopy shows that the emission line profile of this object is asymmetric, showing excess red-wing emission. These properties are consistent with an identification of Ly$\\alpha$ emission at a redshift of $z=5.687 \\pm 0.002$.  The observed broad line width, $\\Delta v_{\\rm FWHM} \\simeq 340$ km s$^{-1}$ and excess red-wing emission also suggest that this object hosts a galactic superwind. ",
        "introduction": "Recent progress in deep optical imaging with 8-10 m class telescopes has enabled new searches for star-forming galaxies beyond redshift 5. In particular, imaging surveys using narrow-passband filters have proved to be a particularly efficient way to find such galaxies (Hu \\& McMahon 1996; Cowie \\& Hu 1998; Steidel et al. 2000; Kudritzki et al. 2000; Hu et al. 2002).  Indeed the most distant Ly$\\alpha$ emitter known to date is HCM 6A at $z=6.56$ (Hu et al. 2002), and more than a dozen Ly$\\alpha$ emitters beyond $z=5$ have been discovered (Dey et al. 1998; Spinrad et al. 1998; Weymann et al. 1998; Hu et al. 1998, 1999, 2002; Dawson et al. 2001, 2002; Ellis et al. 2001), most by using this technique.  One interesting object is J123649.2+621539 at $z=5.190$ which was found serendipitously in the HDF-North flanking fields (Dawson et al. 2002). Its Ly$\\alpha$ emission-line profile shows a sharp blue cutoff and broad red wing emission, both of which are often observed in star-forming systems with prominent wind outflows. These features are also expected from radiative transfer in an expanding envelope. Therefore, Dawson et al. (2002) suggested that the Ly$\\alpha$ profile of J123649.2+621539 is consistent with a superwind with a velocity of $\\sim$ 300 km s$^{-1}$. Galactic superwinds are now considered to be one of the key issues for understanding the interaction and evolution of both galaxies and intergalactic matter (e.g. Heckman 1999; Taniguchi \\& Shioya 2000). In order to improve our knowledge of galactic superwinds at high redshift, a large sample of superwind candidates at $z > 3$ is needed. During the course of our new search for Ly$\\alpha$ emitters at $z \\approx 5.7$, we have found a candidate superwind galaxy at $z = 5.69$.  In this {\\it Letter}, we report its observed properties.  We adopt a flat universe with $\\Omega_{\\rm matter} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $h=0.7$ where $h = H_0/($100 km s$^{-1}$ Mpc$^{-1}$) throughout this {\\it Letter}. ",
        "conclusions": "\\subsection{Star Formation and Superwind Activities}  The observed Ly$\\alpha$ flux is $f$(Ly$\\alpha$) = $(1.49 \\pm 0.33) \\times 10^{-17}$ ergs cm$^{-2}$ s$^{-1}$ based on the ESI spectrum. Given the cosmology adopted in this {\\it Letter}, we obtain an absolute Ly$\\alpha$ luminosity of $L$(Ly$\\alpha$) $\\simeq (5.3 \\pm 1.2) \\times 10^{42} ~ h_{0.7}^{-2}$ ergs s$^{-1}$.  This Ly$\\alpha$ luminosity is comparable to those of other $z > 5$ galaxies; 1) $3.3 \\times 10^{42}$ ergs s$^{-1}$ for HCM 6A at $z=6.56$ (Hu et al. 2002), 2) $6.1 \\times 10^{42}$ ergs s$^{-1}$ for SSA22-HCM1 at $z=5.74$ (Hu et al. 1999), 3) $3.4 \\times 10^{42}$ ergs s$^{-1}$ for HDF 4-473.0 at $z=5.60$ (Weymann et al. 1998), and 4) $8.5 \\times 10^{42}$ ergs s$^{-1}$ for J123649.2+621539.5 at $z=5.19$ (Dawson et al. 2002). Note that all the above luminosities are estimated by using the same cosmology as that used here.  We note that approximately half of the intrinsic Ly$\\alpha$ emission from LAE J1044-0130 could be absorbed by intergalactic atomic hydrogen (e.g., Dawson et al. 2002). In order to reproduce the observed Ly$\\alpha$ emission-line profile a two-component profile fit was made using the following assumptions:\\ 1) the intrinsic Ly$\\alpha$ emission line profile is Gaussian, and 2) the optical depth of the Ly$\\alpha$ absorption increases with decreasing wavelength shortward of the rest-frame Ly$\\alpha$ peak. The resulting fit is shown in Figure 4 (thick curve), which corresponds to the following emission and absorption line parameters -- 1) Ly$\\alpha$ emission: the line center, $\\lambda_{\\rm c, em} = 8030.70$ \\AA, the line flux, $f_{\\rm em} \\simeq 2.42 \\times 10^{-17}$ ergs s$^{-1}$ cm$^{-2}$, and the line width, $FWHM_{\\rm em} \\simeq 650$ km s$^{-1}$; 2) Ly$\\alpha$ absorption: the line center, $\\lambda_{\\rm c, abs} = 8122.73$ \\AA, the optical depth at the absorption center, $\\tau_{\\rm abs} \\simeq 9.85$, and the line width, $FWHM_{\\rm abs} \\simeq 175$ km s$^{-1}$.  This analysis suggests that the total Ly$\\alpha$ emission-line flux amounts to $1.73 \\times 10^{-17}$ ergs s$^{-1}$ cm$^{-2}$.  This is larger by a factor of 1.16 than the observed flux, giving a total Ly$\\alpha$ luminosity of $L$(Ly$\\alpha$) $\\sim 6.1 \\times 10^{42} ~ h_{0.7}^{-2}$ ergs s$^{-1}$. We then estimate the star formation rate of LAE J1044$-$0130.  Using the relation $SFR = 9.1 \\times 10^{-43} L({\\rm Ly}\\alpha) ~ M_\\odot {\\rm yr}^{-1}$ (Kennicutt 1998; Brocklehurst 1971) and the total Ly$\\alpha$ luminosity, we obtain $SFR = 5.6 \\pm 1.1 ~h_{0.7}^{-2} ~ M_\\odot$ yr$^{-1}$.  Note that $SFR$ may be overestimated because part of the Ly$\\alpha$ emission may arise from shock-heated gas if the superwind interpretation is applicable to this object.  Although the signal-to-noise ratio of our ESI spectrum is not high enough to analyze the profile shape in great detail, the presence of the excess red-wing emission seems secure (Fig. 4).  The FWHM of the Ly$\\alpha$ emission is measured to be 340$\\pm$110 km s$^{-1}$ and the full width at zero intensity (FWZI) is estimated to be 890$\\pm$110 km s$^{-1}$. These properties are similar to those of the Ly$\\alpha$ emitter at $z=5.190$, J123649.2+621539, found by Dawson et al. (2002).  \\subsection{Comments on Possible Association with the Quasar SDSSp J104433.04-012502.2 and the Lyman Limit System at $z=5.72$}  The observed redshift of LAE J1044$-$0130, $z=5.687$, is close both to the quasar redshift, $z=5.74$ (see footnote 15) and to that of a Lyman limit system (LLS) at $z_{\\rm LLS} =5.72$ in the quasar spectrum reported by Fan et al. (2000). The redshift difference between LAE J1044$-$0130 and the quasar corresponds to the velocity difference of $\\Delta v \\approx 2370$ km s$^{-1}$ and that between LAE J1044$-$0130 and the LLS corresponds to $\\Delta v \\approx 1476$ km s$^{-1}$. The angular distance between LAE J1044$-$0130 and the quasar is approximately 330 arcsec, giving a co-moving separation of $\\sim$ 13 $h_{0.7}^{-1}$ Mpc. This separation seems too large to identify LAE J1044$-$0130 as a counterpart of the LLS. It seems also unlikely that that LAE J1044$-$0130 is associated with the large-scale structure in which the quasar SDSSp J104433.04-012502.2 resides.  \\begin{deluxetable}{lcccc} \\tablenum{1} \\tablecaption{Journal of imaging observations} \\tablewidth{0pt} \\tablehead{ \\colhead{Band} & \\colhead{Obs. Date (UT)} & \\colhead{$T_{\\rm int}$ (sec)\\tablenotemark{a}}  & \\colhead{$m_{\\rm lim}$(AB)\\tablenotemark{b}} & \\colhead{$FWHM_{\\rm star}$ (arcsec)\\tablenotemark{c}} } \\startdata $B$         & 2002 February 17      &  1680 & 27.1 & 1.2 \\\\ $R_{\\rm C}$ & 2002 February 15, 16  &  4800 & 26.8 & 1.4 \\\\ $I_{\\rm C}$ & 2002 February 15, 16  &  3360 & 26.2 & 1.2 \\\\ $NB816$     & 2002 February 15 - 17 & 36000 & 26.6 & 0.9 \\\\ $z'$        & 2002 February 15, 16  &  5160 & 25.4 & 1.2 \\\\ \\enddata \\tablenotetext{a}{Total integration time.} \\tablenotetext{b}{The limiting magnitude (3$\\sigma$) within a 2$^{\\prime\\prime}$ aperture.} \\tablenotetext{c}{The full width at half maximum of stellar objects in the final image} \\end{deluxetable}  \\vspace{0.5cm}  We would like to thank both the Subaru and Keck Telescope staff for their invaluable help. We would also like to thank the referee for useful comments. This work was financially supported in part by the Ministry of Education, Culture, Sports, Science, and Technology (Nos. 10044052, and 10304013)."
    },
    "0207/astro-ph0207559_arXiv.txt": {
        "abstract": "We have obtained Hubble Space Telescope UV spectra of the white dwarf in GW Lib, the only known non-radially pulsating white dwarf in a cataclysmic variable, and the first known DAZQ variable. The UV light curve reveals large amplitude (10\\%) pulsations in the UV with the same periods (646, 376 and 237 s) as those seen at optical wavelengths, but the mean spectrum fits with an average white dwarf temperature (14,700K for a 0.6M$_{\\odot}$ white dwarf) that is too hot to be in the normal instability strip for ZZ Ceti stars. A better fit is achieved with a dual temperature model (with 63\\% of the white dwarf surface at a temperature of 13300K and 37\\% at 17100K), and a higher mass (0.8M$_{\\odot}$) white dwarf with 0.1 solar metal abundance. Since the blue edge of the instability strip moves to higher temperature with increasing mass, the lower temperature of this model is within the instability strip. However, the presence of accretion likely causes abundance and atmospheric temperature differences in GW Lib compared to all known single white dwarf pulsators, and the current models that have been capable of explaining ZZ Ceti stars may not apply. ",
        "introduction": "With the discovery (van Zyl et al. 2000) of GW Lib as the first non-radially pulsating white dwarf in a cataclysmic variable (CV), asteroseismology could be applied, for the first time, to understand the internal structure of an accreting white dwarf. However, GW Lib does not easily relinquish its secrets. The six optical observing runs of van Zyl et al. (2002; VZ2002) showed that the pulsation spectrum was highly unstable on timescales of months, which is common in cool, hydrogen atmosphere white dwarf pulsators (ZZ Ceti stars or DAVs). While they could identify clusters of signals with very close (2s) spacing, and some frequencies that were repeatedly present in the optical runs (periods of 236, 376 and 648 s), they could not disentangle the various modes to interpret the pulsations. The presence of fine structure, linear combination modes and changes in modes on monthly timescales all indicate that, in the optical at least, GW Lib is typical of ZZ Ceti stars. However, cool pulsators usually have the largest amplitudes. The relatively low amplitude of oscillation (5-17 mmag) in GW Lib could be due to some dilution of the white dwarf optical light by an accretion disk. VZ2002  concluded that they would need a much longer baseline of data (e.g. Kleinman et al. 1998) to solve the problem.  Since GW Lib is too faint for a Whole Earth Telescope (Nather et al. 1990) campaign, a consortium of larger telescopes will be needed to make progress from the ground. But the ultraviolet offers unique opportunities, as the white dwarf usually contributes close to 100\\% of the light in this portion of the spectrum for low mass accretion rate systems (Szkody et al. 2002a; S2002a), and comparison of the UV and optical amplitudes of pulsation can identify the modes of DAVs (Robinson et al. 1995; R1995, Nitta et al. 2000; N2000). Thus, GW Lib became an integral part of our Hubble Space Telescope study of white dwarfs in short period dwarf novae. This project uses the Space Telescope Imaging Spectrograph (STIS) to obtain UV spectra which can be modelled to determine the temperature, gravity, mass, rotation and composition (see G\\\"ansicke et al. 2001; Howell et al. 2002, S2002a, Szkody et al. 2002b; S2002b for results).  GW Lib is one of the WZ Sge type dwarf novae with very infrequent and extreme amplitude outbursts (Howell, Szkody \\& Cannizzo 1995; HSC1995). Its only known outburst occurred in 1983 (Duerbeck 1987). Recent ground based time-resolved spectra (Szkody, Desai \\& Hoard 2000; SDH2000) revealed an orbital period near 79 min, one of the shortest of disk accreting CVs (Warner 1995). Subsequent data with longer time-coverage refines this period to 76.78 min (Thorstensen et al. 2002). The optical spectra show broad absorption troughs (from the white dwarf) surrounding narrow Balmer emission lines (from the low inclination disk), consistent with a low mass transfer rate system. SDH2000 fit the absorption with an 11,000$\\pm$1000K white dwarf at a distance of 114 pc. This temperature was within the general location of the ZZ Ceti instability strip (11,200-12,900K; Koester \\& Holberg 2001, although the edges vary slightly with different authors) but dependent on the estimated disk contribution to the optical light. Surprisingly, the HST data reveal a much hotter white dwarf, as well as the first known DAZQ variable. ",
        "conclusions": "Our STIS data on the only apparent non-radial pulsator in an accreting close binary show that the pulsations in the UV are of large amplitude and have the same periods as seen in the optical. These periods match with those evident in both hot and cool single non-radially pulsating DAs, implying general similarities in structure to single ZZ Ceti stars, whereas the UV spectrum indicates a DAZQ white dwarf with an inhomogenous temperature distribution over its surface, or with a change in temperature during its pulsation. The average UV/optical pulse amplitude ratio is similar to what has been observed and predicted for DAVs in the $\\ell = 1$ mode. While the fit of the spectrum to white dwarf models substantiates a dual-temperature structure (13,300K from 63\\% of the white dwarf surface and 17,100K from 37\\% of the surface), both temperatures are hotter than evident for known single DAVs of 0.6M$_{\\odot}$. However, the cooler of the two temperatures is within the instability strip for more massive white dwarfs ($\\ge$0.8M$_{\\odot}$). Thus, GW Lib may represent a system containing a more massive and spotted white dwarf as compared to typical single DAVs. The mass difference may help to explain why other white dwarfs in CVs that contain similar or cooler temperatures than GW Lib are not pulsating -- e.g. HV Vir, EG Cnc, VY Aqr (S2002b, Howell et al. 2002). The identification of GW Lib begs for time-dependent, non-adiabatic, non-radial pulsation models of white dwarfs undergoing accretion at a low rate."
    },
    "0207/astro-ph0207073_arXiv.txt": {
        "abstract": "Several unexpected astrophysical observations can be explained by gravitationally captured massive axions or axion-like particles, which are produced  inside the Sun  or other stars and are accumulated over cosmic times. Their radiative decay  in  solar outer space would give rise to a  `self-irradiation' of the  whole star, providing the time-independent component of the corona heating source (we do not address here the flaring Sun). In analogy with the Sun-irradiated Earth atmosphere, the temperature and density gradient in the corona$-$chromosphere transition region is suggestive for an omnipresent irradiation of the Sun, which is the strongest evidence for the generic axion-like scenario. The same mechanism is compatible with phenomena like the solar wind, the X-rays from the dark-side of the Moon, the  X-Ray  Background Radiation,  the diffuse X-ray excesses (below $\\sim 1$ keV),  the non-cooling of oldest Stars, etc. A temperature of $\\sim 10^6$ K  is observed in various places, while the radiative decay of  a population of such elusive particles mimics a hot gas, which fits unexpected astrophysical X-ray observations. Furthermore, the recently reconstructed quiet solar X-ray spectrum  during solar minimum supports this work, since it covers the expected energy range, and  it is consistent with the result of a simulation based on  Kaluza-Klein axions above $\\sim$ 1 keV. The derived axion luminosity ($L_a\\approx 0.16L_{\\odot}$) fits the cosmic energy density spectrum and is compatible within 2$\\sigma $ with the recent SNO result, showing the important interplay between any exotic energy loss mechanism and neutrino production. At lower energies, using also a ROSAT observation, only $\\sim 3\\%$ of the X-ray intensity is explained. Data from orbiting X-ray Telescopes provide upper limits for particle decay rates 1 AU from the Sun, and suggest  new types of searches on Earth or in space. In particular, X-ray observatories, with an unrivalled equivalent fiducial volume of $\\sim 10^3~m^3$ for the 0.1  - 10 keV range, can search  for the radiative decay of new particles even from existing data. This work introduces the elongation angle of the X-ray Telescope relative to the Sun as a relevant new parameter. ",
        "introduction": "The direct detection of dark-matter (DM) particles has proved elusive since the first gravitational observation of  non-luminous matter in the Universe. So far, the outcome of the intense experimental and theoretical  work in the field of dark matter during the last $\\sim 20$ years was the birth of the new discipline of astroparticle physics.   In this work, a rather old question is reconsidered as to whether a large number of as yet unexplained astrophysical phenomena occur because of the involvement of novel very weakly interacting particles or additional as yet unknown properties of existing particles. For example, in order to explain the ionization  state of the intergalactic medium, and the anomalous ionization of the interstellar medium \\cite{bowyerxxx}, speculations included a widespread source of ionizing UV-photons from the electromagnetic decay  of real or hypothetical exotic particles (clustered in haloes) \\cite{bowyer}, e.g., massive dark-matter neutrinos \\cite{sciama,abazajian}; however, observations appear to have ruled out this model \\cite{bowyerxxx,bowyer1xxx}. Considerations based on axions apply as well,  invoking either its $2\\gamma $ decay mode or its coherent conversion to a photon inside astrophysical electric/magnetic fields via the Primakoff effect \\footnote{ a) The absence of a monochromatic axion line from the night Sky expected from the $a\\rightarrow \\gamma \\gamma$ decay  of relic axions in the visible \\cite{turner} almost excluded an axion rest mass in the $\\sim  1$--10~eV range. b) Until recently, a conventionally expected thermal X-ray spectrum from solar axions converted inside the solar magnetic fields with mean energy of $\\sim 4.4$~keV could not have  been disentangled from the derived solar X-ray spectrum (e.g. Ref.~\\cite{gudel3xxx}). In any case, such a mechanism could not explain the  main observations addressed in this work, because such converted axions give rise to photons always emitted away from the Sun. Moreover, even in the scenario of massive axions from the Sun, the possibility to detect photons emitted in flight through axion decays was considered to be \"presumably indistinguishable from the general background radiation\" \\cite{dienes}. }.  Recently, in order to explain the as yet unknown underlying mechanism(s) of the Gamma Ray Bursts, massive axions with properties far beyond the widely accepted theoretical axion concepts have been considered, providing a built-in dissipationless energy transfer mechanism from the hypothetical energy generating core to the outside layers  some 100$-$1000 km away \\cite{grb}. Even though none of such intriguing ideas has been established so far \\cite{henry1}, an additional electromagnetic energy source in the cosmos seems to be necessary. Structure formation in warm dark matter  (WDM) cosmological models \\cite{narayanan} provides a lower limit to the mass of the WDM particle candidate of $\\sim 0.75$ keV, but with a radiative decay lifetime as large as $\\sim 10^{16}\\times $Hubble time; however, halos in Galaxies and clusters of Galaxies can be an enormous \"fiducial volume\" of DM particles.  In this work, we argue that the photon emission of some hypothetical particles, referred to more generally  as `axion-like', could be involved  in different unexplained astrophysical observations. An extensive search of astrophysical literature has been undertaken, which includes also some 10--50 years old observations. In this work we focus mainly  on : \\begin{itemize} \\item[a)] the solar corona problem and related observations; \\item[b)] the  observed X-rays from the direction of the dark side of the Moon; \\item[c)] the soft-X-ray background radiation; \\item[d)] the (diffuse) soft X-ray excess phenomenon; \\item[e)] first simulation results  in the frame of an axion scenario. \\end{itemize} Following the reasoning of this work, we also suggest performing a {\\it specific} axion search in a new type of experiment, either on the surface (and in space) or underground, aiming  directly at detecting the $2\\gamma $ decay / interaction mode. The alternative case of a single photon emission seems to be beyond the present sensitivity of an Earth-bound detector and only orbiting X-ray telescopes could be considered. In addition, some astrophysical measurements could be reconsidered or re-analysed.  X-ray space detectors could operate also as sensitive orbiting axion telescopes. ",
        "conclusions": "A missing  explanation of an astrophysical observation is actually suggestive to search for an exotic approach. The framework of the celebrated dark-matter physics world is a source of possible exotic  solutions. In the cases considered  here,  a same axion-like scenario consistently explains the usually different alternative  conventional solutions (if they exist at all). This  scenario  is not supposed to abandon globally previous models, which describe actually related findings; it can  be rather complementary, providing a so far missing physics input.   A temperature (component) of a few  10$^6$K ($\\sim 0.3$ keV ), which  appears in  so diverse astrophysical  places, such as from the solar corona to Clusters of Galaxies and  probably beyond, is associated with several  unexpected significant observations. In order to  explain this in a combined way, we reach the conclusion that some new particles~  $-$we use massive axions as a generic example$-$  ~must be involved in processes occurring inside { and} outside Stars. For example, relatively short-lived massive axions  have  been considered in theories with extra dimensions \\cite{dienes,dilella5xxx}; the  two-photon decay mode remains dominant providing  theoretical support to  our purely  observationally/astrophysically motivated claim of celestial axion-like signatures in the $\\sim$ keV range. In our favoured scenario, axion-like particles escape from their place of birth, e.g., from the interior of the Sun (or that of other Stars in the Sky), get gravitationally trapped and  decay in outer space.    Alternatively, a more or less isotropic radiative decay of other hypothetical particles, e.g. massive neutrinos, could in principle also explain the astrophysical observations  considered here. Only laboratory  experiments  could clarify this issue. We  give a (theoretically) unbiased  parameter space how and where to directly search for such exotica. Fortunately, the two-photon decay mode allows to have a very high detection sensitivity inside a large TPC, because of the much suppressed uncorrelated two-prong background  events within a small distance, narrow time and same energy. High-performance low-threshold detectors developed primarily for high-energy physics experiments can  also be utilized for this kind of astro-particle physics.   Following the reasoning for the suggested axion-like scenario, we should also notice that the previous failed searches \\cite{raffelt333,carlson9xxx} for axions converted to X-rays inside the external solar/stellar magnetic fields do not contradict this work. At first, this missing signal can be due to an accordingly small coupling strength and/or the required very small axion restmass. Independent on this actually unknown value, the \"conventional\" Primakoff effect should result to radially  outwards emitted X-rays, excluding  a self-irradiation of the Sun, which we consider as the cornerstone of the reasoning of this work. The tentative solar Kaluza-Klein  model provides reasonable gravitational capture rates,  but  it fails to  completely explain the low energy part of the reconstructed solar X-ray spectrum. However, other possible sources like the solar Bremsstrahlung-axions seem to have reasonably  low energy, but with a smaller  gravitational capture rate. For  astrophysical standards, the encountered  discrepancies are actually not particularly large.  Finally, the estimated axion density due to gravitational trapping  by the Sun can exceed a critical value, which is necessary for the appearance of a BEC,  with unforeseen implications. The continuous dynamic coronal phenomen  might be a manifestation of such processes. Thus, the predicted particles orbiting around the Sun can become an invaluable clue to physics beyond the standard (solar) model, explaining first of all the as yet mysterious properties of the chromosphere and the corona. The gravitational trapping of massive particles emitted by the Star itself provides a  mechanism for the appearance of boson  clumps around a Star.  {\\sf In conclusion,} the strongest and rather direct evidence in favour of the axion scenario  comes from the sofar unexplained solar corona related  observations like its  heating mechanism, its narrow interface to the chromosphere, the chromosphere itself, together with the striking similarities of the temperature/density profiles with the Sun-irradiated Earth atmosphere. The recently experimentally reconstructed solar X-ray spectra combined with X-ray measurements from orbiting observatories are the first potential direct signatures for this work. They show how these orbiting instruments can directly search for decaying particles in (near) outer space, providing also the expected maximum rate in future investigations of this kind. The relevant new parameter introduced by this work is the as yet disregarded elongation angle of the X-ray Telescope relative to the Sun. The other astrophysical observations we have addressed, in particular when they are seen combined,  provide an additional  piece of evidence, probably of not minor importance at the end.  \\vskip2.0cm"
    },
    "0207/astro-ph0207590_arXiv.txt": {
        "abstract": "Recently  \\citet{sanwal02} reported  the first  clear detection  of absorption features  in an  isolated  neutron star,  \\pulsar.  Remarkably their  spectral modeling demonstrates that the  atmosphere cannot be Hydrogen. They speculated that  the  neutron star  atmosphere  is indicative  of  ionized  Helium in  an ultra-strong ($\\sim1.5\\times10^{14}$  G) magnetic  field. We have  applied our recently developed atomic model \\citep{mori02} for strongly-magnetized neutron star atmospheres  to this problem.  We find that  this model, along  with some simple atomic physics arguments,  severely constrains the possible composition of  the atmosphere. In  particular we  find that  the absorption  features are naturally  associated with  He-like  Oxygen or  Neon  in a  magnetic field  of $\\sim10^{12}$  G, comparable  to  the  magnetic field  derived  from the  spin parameters of  the neutron  star. This interpretation  is consistent  with the relative line strengths and widths  and is robust. Our model predicts possible substructure in  the spectral  features, which has  now been reported  by \\xmm \\citep{mereghetti02}. However  we show  the Mereghetti et  al. claim  that the atmosphere is  Iron or  some comparable high-Z  element at $\\sim10^{12}$  G is easily ruled out by the \\chandra and \\xmm data. ",
        "introduction": "A major goal of neutron star  (NS) research has remained unrealized despite 30 years of  effort - to determine  the fundamental properties  of the superdense matter  in a  NS  interior, in  particular  its equation  of  state (EOS)  and composition. There  are many approaches which  can be employed  in this effort (\\citet{vankerkwijk01}  and   references  therein),   but  one  of   the  most intensively studied is to exploit the thermal radiation from isolated NS. This radiation can  be used to deduce  information about the EOS  from neutron star cooling theory \\citep{tsuruta02,yakovlev99}.  Extracting information about the NS interior is not straightforward,  however, since the observed spectrum does not represent the  NS surface emission, but is  modified by radiative transfer effects in the NS atmosphere.   The problem of unfolding the observed spectrum and understanding NS interiors thus depends on deducing the composition of the atmosphere.  NS atmospheres are interesting in their  own right. We know little about their nature,  other   than  that  some   NS  probably  have   Hydrogen  atmospheres \\citep{pavlov02_2},  since they  fit  well to  the sophisticated  H-atmosphere models that have been developed \\citep{pavlov95_1, zavlin02}. The observations have been silent on the question of whether non-Hydrogen atmospheres exist.  The  NS  \\pulsar  was   discovered  by  \\citet{helfand84}  with  the  Einstein Observatory  and   is  associated  with   the  SNR  PKS   1209-51.  Subsequent observations established it as a  prototypical radio silent NS associated with a SNR \\citep{bignami92, caraveo96}. ROSAT  and ASCA observations were fit with a   black-body  spectrum  \\citep{mereghetti96,   vasisht97}  and   a  Hydrogen atmosphere    model   \\citep{zavlin98_2}.     No    X-ray   pulsations    were detected.    \\chandra   detected   X-ray    pulsations   with    $P=0.424$   s \\citep{zavlin00}. While  this is perhaps  not surprising in light  of previous work, the  small period derivative estimated  \\citep{pavlov02_1} certainly was unexpected.  The inferred surface B-field is $\\sim(2-4)\\times10^{12}$ G.  Even  more  remarkable,  \\citet{sanwal02}  (hereafter  SZPT)  have  discovered features at $\\sim0.7$ keV and  $\\sim1.4$ keV are required to obtain acceptable fits  to the  spectrum (which  was fit  with a  $2.6\\times 10^6$  K underlying black-body continuum).  A feature of marginal significance  and unclear origin was   also  noted   at  $\\sim2$   keV.  A   subsequent  observation   by  \\xmm \\citep{mereghetti02} has  provided some hint of substructure  in the features, and has marginally detected the $\\sim2$ keV feature seen by \\chandra.  Three different phenomenological fits indicative of either absorption lines or edges  were  used  by SZPT  to  fit  the  two  strong features.   Of  enormous significance,  the  spectrum  was  shown  to  be  inconsistent  with  Hydrogen atmosphere  models.  The intense  effort  which  has  gone into  the  Hydrogen atmosphere  models  and their  success  in  explaining  other NS  observations \\citep{pavlov02_2}  lends  credence  to  the inevitable  conclusion  that  the atmosphere  of \\pulsar is  something other  than Hydrogen.   Unfortunately, as pointed out  elsewhere \\citep{zavlin02},  work on non-Hydrogen  atmospheres at high B-fields is much less developed.  SZPT argued on various grounds that the features  could  not  arise  via  cyclotron lines.  Instead  they  tentatively suggested emission  from a  once ionized-Helium atmosphere  with a  B-field of $1.5\\times 10^{14}$  G.  This B-field  is inconsistent with that  derived from the  spin parameters,  but SZPT  argue this  could be  due to  an off-centered B-field or  glitches affecting the  $\\dot{P}$ measurement. Others  have argued for   a    cyclotron   line   solution   at    lower   B-field   \\citep{xu02}. \\citet{mereghetti02} claimed an atmosphere of Iron or other high-Z elements at a B-field  of $\\sim10^{12}$ G, although  they did not actually  fit their \\xmm data to a model.  In  this   paper  we  offer   alternate  interpretations  of   these  spectral features.  All  our  interpretations  involve atomic  transitions,  mainly  in He-like  ions of  mid-Z  elements  at $B\\sim10^{12}$  G,  consistent with  the B-field  derived  from  the NS  spin  properties.  The  most likely  of  these interpretations  is that  the  NS  atmosphere contains  Oxygen  or Neon;  most noteworthy  is that {\\it  all} our  models, whether  considering just  the two strong features or including the third weak feature, demand mid-Z elements for an acceptable solution.  Our model, combined with the  \\xmm and \\chandra data, easily rule out the Iron and high-Z solutions of \\citet{mereghetti02}.  Some comments are in order on  our approach. The atomic spectroscopy data used in this analysis is based on  a novel approach for obtaining fast and accurate solutions to  the Schr{\\\"o}dinger equation  for B-fields in the  Landau regime (appropriate    for   all    cases   considered    here).     This   approach, multiconfigurational,  perturbative,  hybrid,  Hartree,  Hartree-Fock  theory, allows rapid  computation of transition energies and  oscillator strengths for arbitrary atom,  ion, excitation state and  B-field (\\citet{mori02}, hereafter MH02a).  This  permits  a   complete  search  of  all  possible  spectroscopic transitions consistent with the given line or edge energies.  While  it may  appear that  this  approach produces  an uninterestingly  large number  of   potential  solutions,  we   demonstrate  in  a   companion  paper (\\citet{mori02_2}, hereafter MH02b)  that this is not the  case.  We show that the  presence of  two or  more  line or  edge features  provides a  remarkable robustness  to  a  host   of  poorly-understood  atomic  physics  effects  and unambiguously restricts the atmosphere composition to mid-Z elements. ",
        "conclusions": "Our most important conclusion is  that {\\it all} viable solutions for B-fields comparable  to  that  inferred  from  the NS  spin  parameters  require  mid-Z atmospheres in  both the  two and three  spectral feature case.  Our solutions which most  closely correspond to  the inferred B-field demand  He-like Oxygen and Neon.  In contrast  to this paper,  SZPT proposed ultrahigh B-field  solutions. Their Helium  atmosphere appears  problematic, however.  In order  to get  the right feature positions  SZPT's H-like  Helium features must  of necessity  arise in transitions from electrons in the $(0,0)$ or $(1,0)$ quantum states to $(0,1)$ and $(1,1)$ states.  The $(0,1)$ state is heavily depopulated  at the types of ion densities one expects and the $(1,1)$ state is actually autoionized due to the  large nuclear  cyclotron energy.  In fact  the motional  Stark  shift and pressure effects in  Helium will likely destroy all  the loosely-bound states. Consequently one  is forced to assume  the features are  absorption edges, and this  leads  to  severe  difficulty  explaining the  relative  edge  strengths (MH02b).  \\citet{mereghetti02} did  not fit their  \\xmm data, but  proposed nevertheless that the  spectral features could be  attributed to Iron or  some other high-Z element  at  a  field  $\\sim10^{12}$  G.  Our current  work  rules  out  their proposal. There is no combination  of B-field, redshift, element or ionization states, even  with unphysical, arbitrary  pressure shifts, that  is consistent with the \\xmm and \\chandra data.  The interpretation we present here makes an important prediction: the observed features may show substantial  substructure when observed with higher spectral resolution.   Transition  energies from  the  tightly  bound  ground state  to loosely bound  final excited states  differing in longitudinal  quantum number are  rather close  to each  other, and  may be  blended. Details  of predicted energies and  transition assignments are presented in  MH02b. Whether detailed atmospheric models  would lead to a  washing out of  this substructure through broadening is  not clear, but its  detection would certainly  confirm that the features are atomic transitions since we would not expect similar substructure in cyclotron  lines. In fact, \\citet{mereghetti02} claim  a marginal detection of such substructure.  Concerning the  B-field derived from the  free parameter fits to  the data, we note that our Neon  and Oxygen line solutions are all within  a factor of two or three of the B-field derived from spin parameters. The discrepancy may lie in the  NS B-field geometry or  line blending. Line blending  will affect our B-field   (and  redshift)  estimate,   but  will   not  affect   our  element identifications.  The analysis here, along with that of  SZPT, begs the question of why there is anything in the  atmosphere other than Hydrogen.  Oxygen  or Neon are expected constituents  near the  NS mass  cut in  many models  of  supernova explosions \\citep{arnett96}. Indeed due to turbulent mixing in such explosions Oxygen can be found  throughout the  core \\citep{herant94}.  But  given that a  very thin layer of Hydrogen  accreted from the interstellar medium  is enough to produce an optically  thick Hydrogen atmosphere, we  are left to  ponder mechanisms by which  such  accretion can  be  inhibited,  or  at least  permit  simultaneous emission from  a mid-Z element. One  possibility is that the  timescale for an optically thick  layer to accrete  onto the surface  is shorter than  that for gravitational stratification in  the atmosphere \\citep{vankerkwijk01}. In that case,  even trace  amounts of  mid-Z elements  could lead  to  strong emission \\citep{pavlov02_2}. Optical  observations of  the associated SNR  indicate the presence of  Oxygen and it  has been argued  that the unusually  high Galactic latitude  of this  NS implies  the Oxygen  originated in  the  progenitor star \\citep{ruiz83}. Thus  a rich  source of Oxygen  is available.  Alternately the mid-Z  atmosphere may  be associated  with the  NS mass  cut, and  a propeller effect may be severely inhibiting  accretion of Hydrogen. We note that several of our proposed solutions allow a Hydrogen-rich environment (MH02b).  For  more progress to  be made  on \\pulsar  it will  be necessary  to develop complete models  of NS  atmosphere for mid-Z  elements incorporating  the more sophisticated atomic data bases we have utilized here. This source has much to tell us in the years  to come, especially when higher resolution spectroscopic observations become available."
    },
    "0207/astro-ph0207245_arXiv.txt": {
        "abstract": "We present the first detailed and homogeneous analysis of the $s$-element content in Galactic carbon stars of N-type. Abundances of Sr,Y, Zr (low-mass $s$-elements, or ls) and of Ba, La, Nd, Sm and Ce (high-mass $s$-elements, hs) are derived using the spectral synthesis technique from high-resolution spectra. The N-stars analyzed are of nearly solar metallicity and show moderate $s$-element enhancements, similar to those found in S stars, but smaller than those found in the only previous similar study (Utsumi 1985), and also smaller than those found in supergiant post-AGB stars. This is in agreement with the present understanding of the envelope $s$-element enrichment in giant stars, which is increasing along the spectral sequence M$\\rightarrow$MS$\\rightarrow$S$\\rightarrow$SC$\\rightarrow$C during the AGB phase. We compare the observational data with recent $s$-process nucleosynthesis models for different metallicities and stellar masses. Good agreement is obtained between low mass AGB star models ($M\\la 3 M_\\odot$) and $s$-elements observations. In low mass AGB stars, the $^{13}$C($\\alpha, n)^{16}$O reaction is the main source of neutrons for the $s$-process; a moderate spread, however, must exist in the abundance of $^{13}$C that is burnt in different stars. By combining information deriving from the detection of Tc, the infrared colours and the theoretical relations between stellar mass, metallicity and the final C/O ratio, we conclude that most (or maybe all) of the N-stars studied in this work are intrinsic, thermally-pulsing AGB stars; their abundances are the consequence of the operation of third dredge-up and are not to be ascribed to mass transfer in binary systems. ",
        "introduction": "It is known that the chemical composition of the interstellar medium is oxygen-rich. Hence, the overwhelming majority of the stars are formed with a carbon to oxygen ratio lower than unity, and most of them do not change this property during their evolution. However, there are exceptions: several classes of stars are known, whose carbon to oxygen ratio in the envelope is larger than unity (by number of atoms). These stars are named carbon (C) stars. The existence of carbon stars must be related to some specific mechanism, acting on a limited population of objects. There are basically three possibilities: i) the carbon enrichment is the result of a deep mixing process suitable to pollute the photosphere with the carbon synthesized by shell He-burning in the stellar interior; ii) it is due to mass transfer of carbon-rich material after the stellar birth; iii) it dates back to the star's birth, due to an anomalous composition of the parental cloud, more enriched in carbon than in oxygen. Concerning the last hypothesis, so far no such interstellar clouds have been detected in Galactic disk environments, though the hypothesis is not completely ruled out for certain low metallicity stars (Beveridge \\& Sneden 1994). In the second case, the origin of the carbon-rich material is simply moved to another place, i.e. the carbon-rich primary component of a binary system. For instance, the stars of the CH spectral type (Luck \\& Bond 1991; Vanture 1992a,b,c) most likely owe their carbon enhancement to the transfer of carbon-rich material from a companion (now a white dwarf). These stars are usually named {\\it extrinsic} asymptotic giant branch (AGB) stars to distinguish them from those deriving their carbon enhancement from nuclear processing in their interiors (hypothesis i), which are called {\\it intrinsic}. Here, we shall deal with a specific subclass of carbon stars, those of spectral type N. We shall see later that most probably all of them can be classified as intrinsic C-rich objects. They are formed through the mixing into the envelope of newly produced $^{12}$C from He burning. Stellar evolution then limits the candidates to evolved stars with masses $1\\la M/M_\\odot\\la 8$, in particular those ascending AGB. The structure of an AGB giant is characterized by a degenerate CO core, by two shells (of H and He) burning alternatively, and by an extended convective envelope. In the HR diagram, these stars lay close to the brightest part of the Hayashi line. They become long period variable of the irregular, semiregular or Mira types, presenting large mass-loss rates: $10^{-8}$ to $10^{-4}$ M$_\\odot/$yr (Wallerstein \\& Knapp 1998). As a consequence, a thick circumstellar envelope eventually forms, sometimes developing detached shells. Depending on its chemical composition and optical thickness, this  circumstellar material can obscure partially or completely the central star at optical wavelengths (Knapp \\& Morris 1985; Olofsson et al. 1993; Marengo et al. 2001). Schwarzschild \\& H\\\"arm (1965) early showed  that thermal instabilities in the He-shell (thermal  pulses, TP) occur periodically during the advanced phases of AGB evolution.  During a TP the whole region between the H shell and the He shell (called ``the He intershell'') becomes convective. After each TP, the convective envelope penetrates downward dredging-up material previously exposed to incomplete He-burning conditions. This phenomenon is called {\\it third dredge up} (TDU), and its main consequence is the increase of the carbon content in the envelope so that, eventually, the C/O ratio can exceed  unity and the star becomes a carbon star. In such a way, the carbon content in the envelope is expected to increase along the spectral sequence M$\\rightarrow$MS$\\rightarrow$S$\\rightarrow$SC$\\rightarrow$C, stars of spectral class C showing C/O$>1$ (see e.g. Iben \\& Renzini 1983; Smith \\& Lambert 1990).  Another important consequence of TDU is the  enrichment of the envelope in {\\it s}-elements. The necessary neutrons  for the $s$-process are released by two reactions: $^{13}$C$(\\alpha,n)^{16}$O, which provides the bulk of the neutron flux at low neutron densities ($N_n\\la 10^7$ cm$^{-3}$), and $^{22}$Ne$(\\alpha,n)^{25}$Mg, which is activated at temperatures $T\\ga 3.0\\times  10^8$ K, providing a high peak neutron density ($N_n\\sim 10^{10}$ cm$^{-3}$) and is responsible of the production of $s$-nuclei controlled by reaction branchings (see Wallerstein  et  al. 1997; Busso, Gallino, \\& Wasserburg 1999 and references therein). The abundances of $s$-nuclei are known to increase along the above mentioned spectral sequence, as the star gradually ascends the AGB. Evidence of this was provided during the last few decades by several studies on AGB stars of different spectral types:  MS and S stars (C/O$<1$) (Smith \\& Lambert 1985, 1986, 1990); SC stars (C/O$\\sim 1$) (LLoyd-Evans 1983; Abia \\& Wallerstein 1998, hereafter Paper I) and even post-AGB supergiants of spectral types A and F  (Van Winckel \\& Reyniers  2000; Reddy, Bakker, \\& Hrivnak 1999).  The above studies found consistent enhancements of $s$-elements with respect to ``normal'' red giants assumed as comparisons (or with respect to the Sun): from [s/Fe]$\\approx +0.3$ in MS giants to [s/Fe]$>+0.5$ in S stars\\footnote{We adopt the usual notation [X/Y]$\\equiv$ log (X/Y)$_{\\rm{program star}}-$ log (X/Y)$_{\\rm{comparison star}}$ for the stellar value of any element ratio X/Y.}. The enrichment seems to continue along the AGB phase until the planetary nebula ejection. Indeed, post-AGB supergiants show high enhancements, [s/Fe]$> +1.0$. In this scenario, however, normal carbon stars (N-type)\\footnote{There are carbon stars of R-type and J-type, which probably owe  their carbon enhancement to a mechanism different than the TDU. These stars  are not significantly enhanced in $s$-elements (Dominy 1985; Abia \\& Isern 2000)} have still to find a place. This is so because of the complex spectra of N stars, which are so crowded with molecular and atomic absorption features (many of which unidentified) that abundance analysis has been strongly limited.  The situation has not improved much with the advent of spectrum synthesis techniques. Indeed, the only abundance studies available to date are those by Kilston (1975) and Utsumi (1970, 1985), still based on abundance indexes and on low-resolution photographic spectra, respectively. These works suggested that N stars were $s$-element rich, showing overabundances by a factor of ten with respect to the Sun. As far as the AGB phase is concerned, this figure is not in direct contradiction with the accepted general scenario of $s$-nuclei enhancement (see e.g. Busso et al. 1995). More likely, such high production factors might be difficult to reconcile with the enrichment in post-AGB supergiants showing similar, or even smaller, abundances. However, the large observational uncertainties for C-rich red giants, and the lack of adequate model atmospheres has so far prevented a solid theoretical interpretation.  In a previous  work (Abia et al.  2001, hereafter Paper  II) we presented high-resolution spectroscopic  observations for  a sample of  N stars focusing  our attention on light $s$-elements (ls), sited at the abundance peak near the neutron-magic number $N=50$, around the $^{85}$Kr branching point of the $s$-process path. We showed how the analysis of the abundance ratios of Rb (a neutron-density sensitive element, see e.g. Beer \\& Macklin 1989) relative to its neighbors (Sr, Y and Zr) yields information on important details of the $s$-process mechanism operating, and on the initial stellar mass. We concluded that $s$-processing suggests low mass stars (LMS, $M\\la 3 M_\\odot$) as the likely parents of C(N) giants.  In LMS the major neutron source is $^{13}$C, which burns radiatively in a tiny layer during the interpulse phase (Straniero et al. 1995) at relatively low temperatures ($\\approx$ 8 keV), as a consequence of the formations of a $^{13}$C-rich {\\it pocket} in the intershell region. In the rarer AGB stars of intermediate mass (IMS) ($M\\ga 4 M_\\odot$), $^{22}$Ne would instead be favored as a neutron source, by the higher temperature in thermal pulses. Because of the very different neutron density provided by the two neutron-producing reactions, different compositions are expected from them, especially for the ls mixture. It is on this basis that in Paper II we drew our conclusions for the initial masses.  In Paper II we also discussed the $^{12}$C/$^{13}$C ratios measured in the sample stars, showing that  most of them cannot be  explained by canonical stellar models on the AGB phase, requiring probably the operation of an ad-hoc mixing mechanism. This mechanism is often indicated with the term 'cool bottom process' (CBP). It is expected to occur in low-mass stars during the red giant branch and perhaps, also during the AGB phase (Wasserburg, Boothroyd, \\& Sackmann 1995; Nollett, Busso, \\& Wasserburg 2002).  In the present work we have extended our study of $s$-element nucleosynthesis in N stars to the nuclei belonging to the second $s$-process peak: Ba, La, Ce, Nd, and Sm. We also re-analyzed the $s$-element abundances already derived in Paper II, by applying the spectral synthesis method. Together, these data are compared with recent models for $s$-processing in AGB stars at the metallicities relevant  for our  sample stars. Furthermore, using the infrared properties and Tc content, we discuss the possibility that our stars be extrinsic carbon stars, concluding that this is unlikely for several reasons. The structure of the paper is the following. In $\\S~2$ we present the characteristics of the stars and the spectroscopic observations. Section $\\S~3$ describes the method of  analysis and the  sources of error.  Our abundance results are reported in $\\S~4$ together with a comparison with similar studies in other AGB stars and with models of AGB nucleosynthesis. Finally, in $\\S~5$ we summarize the main conclusions that can be drawn from this work. ",
        "conclusions": "In this paper we have presented for the first time a rather complete sample of $s$-process abundance measurements in N-type carbon stars, extended to light and heavy species, across the $s$-process peaks with neutron numbers $N=50$ and $N=82$. Making use of a large line list, and of up-to-date model atmospheres for carbon rich giants, we have deduced the $s$-element abundances through spectrum synthesis technique. In doing so, we have also verified previous results (Paper II) for the light $s$-elements across the $^{85}$Kr branching point of the $s$-path. N stars turn out to be characterized by s-element abundances very close to, or slightly higher than, those found for S stars. Compared with the only analysis performed before (Utsumi 1985), our abundances are clearly smaller; they are smaller also with respect to those of post-AGB supergiants (Reddy, Bakker, \\& Hrivnak 1999; Van Winckel \\& Reyniers 2000), and this gives rise to a general picture in which the surface enrichment continuously increases along the evolutionary sequence producing MS, S, and N stars, and subsequently yellow post-AGB supergiants (Reddy et al. 2002). We have compared our data with model envelope compositions obtained from previously published calculations of AGB nucleosynthesis and mixing. Good agreement is obtained between low mass star models and $s$-element observations. Several pieces of evidence (from the detection of Tc lines and infrared colors, to the theoretical relations between the initial metallicity, the mass, and the final C/O ratio) lead us to conclude that most (or perhaps all) our sample stars are intrinsic TP-AGB stars, so that their abundances are locally produced by the occurrence of third dredge-up during the TP-AGB phases, and not generated by mass transfer in binary systems.   Acknowledgements. Data from the VALD  database at Vienna were used for the preparation of this paper. K. Eriksson, and the stellar atmosphere group of the Uppsala Observatory are thanked for providing the grid of atmospheres. The 4.2m WHT and the  2.5m NOT are operated on the island of La  Palma by the  RGO in the Spanish Observatory of the Roque de los Muchachos of the Instituto de Astrof\\'\\i sica  de Canarias. This work was  also based  in part on  observations  collected with  the 2.2m telescope at the German-Spanish  Astronomical Centre, Calar Alto. It was partially  supported by the spanish  grants AYA2000-1574, FQM-292, by  the  Italian   MURST-Cofin2000 project  `Stellar Observables  of Cosmological  Relevance'  and by  the French-Spanish  International Program for Scientific Collaboration, PICASSO HF2000-0087."
    },
    "0207/astro-ph0207523_arXiv.txt": {
        "abstract": "Nonthermal radiation is observed from clusters of galaxies in the radio, hard X-rays, and possibly in the soft X-ray/UV bands. While it is known that radiative processes related to nonthermal electrons are responsible for this radiation, the sites and nature of particle acceleration are not known. We investigate here the acceleration of protons and electrons in the shocks originated during mergers of clusters of galaxies, where the Fermi acceleration may work. We propose a semi-analytical model to evaluate the Mach number of the shocks generated during clusters mergers and we use this procedure to determine the spectrum of the accelerated particles for each one of the shocks produced during the merger history of a cluster. We follow the proton component, accumulated over cosmological time scales, and the short lived electron component. We conclude that efficient particle acceleration, resulting in nonthermal spectra that compare to observations, occurs mainly in minor mergers, namely mergers between clusters with very different masses. Major mergers, often invoked to be sites for the production of extended radio halos, are found to have on average too weak shocks and are unlikely to result in appreciable nonthermal activity. ",
        "introduction": "\\label{sec:intro} Rich clusters of galaxies are strong X-ray sources with luminosity typically in the range $L_X\\sim 10^{43}-10^{45} erg/s$. The X-ray emission is well explained as bremsstrahlung radiation of the very hot ($T \\sim 10^8 K$), low density ($n_e \\sim 10^{-3} cm^{-3}$) and highly ionized intracluster electron gas.  There is now compelling evidence for the existence, besides the thermal electron gas, of a nonthermal population of particles, responsible for extended synchrotron radio halos in a growing fraction of the observed clusters (see Feretti et al., 2000 for a recent rewiew), as well as for hard X-ray (HXR) and extreme ultra violet (EUV) excesses detected in a few clusters (see e.g. Fusco Femiano et al., 1999, 2000; Lieu et al., 1996). While it is clear that radio emission is due to synchrotron radiation from relativistic electrons, it is not as clear how these particles are accelerated. Several models have been proposed, based on shock acceleration of electrons in merger shocks (Roettiger, Burns \\& Stone, 1999; Sarazin 1999; Takizawa \\& Naito, 2000; Fujita \\& Sarazin 2001) or models in which electrons are secondary products of hadronic interactions (Dennison 1980, Colafrancesco \\& Blasi 1998; Blasi \\& Colafrancesco 1999; Dolag \\& Ensslin 2000) and finally models in which electrons are continuously reenergized by turbulence (Schlickeiser, Sievers, \\& Thiemann 1987; Brunetti et al. 2001; Ohno, Takizawa, \\& Shibata 2002).  HXR and EUV radiation in excess of the thermal emission may be generated by inverse Compton scattering (ICS) of relativistic electrons off the photons of the cosmic microwave background radiation. When applied to the Coma cluster, these models require values of the volume averaged magnetic field which are smaller than those measured through Faraday rotation, which are typically of several $\\mu G$ (Eilek 1999; Clarke, Kronberg \\& B\\\"{o}ringer 1999). This conclusion can be possibly avoided only by constructing models in which a cutoff in the electron spectrum is tuned up in order to reduce the corresponding synchrotron emission. In these cases the magnetic field can be as high as $0.3-0.4\\mu G$ (Brunetti et al. 2001). In the case of a secondary origin for the radiating electrons, the small magnetic fields imply a large cosmic ray content in the intracluster gas. In the case of the Coma cluster, the gamma ray upper limit found by Sreekumar et al. (1996) is exceeded by the gamma ray flux from the decay of neutral pions, as shown by Blasi \\& Colafrancesco (1999).  The HXR excess might also be the result of bremsstrahlung emission from a population of thermal electrons whose distribution function is slightly different from a Maxwell-Boltzmann (MB) distribution (Ensslin, Lieu \\& Biermann 1999; Blasi 2000; Dogiel 2000; Sarazin \\& Kempner 2000). A tail might in fact be induced in the MB distribution by the presence of MHD waves that resonate with part of the thermal electrons (Blasi, 2000). This model requires an energy input comparable with the energy budget of a cluster merger, and implies a substantial heating of the intracluster gas (this was shown by Blasi (2000) by solving the full Fokker-Planck equations, including Coulomb scattering). If the process lasts for too long a time (larger than a few hundred million years) the cluster is heated to a temperature well in excess of the observed ones, and the model fails. In this case the arguments presented by Petrosian (2001) apply.  The presence of tails in the MB electron distribution can be tested through observations of the Sunyaev-Zeldovich (SZ) effect, as proposed by Blasi, Olinto \\& Stebbins (2000) (see Ensslin \\& Kaiser (2000) for a general discussion of the SZ effect including nonthermal effects). Clearly, by simply observing radio radiation and hard X-ray radiation from clusters, it is extremely difficult, if not impossible to discriminate among classes of models. The study of the SZ effect allows one to partly break the degeneracy. An even more powerful tool is represented by gamma ray astronomy. Some of the models in the literature predict gamma ray emission to some extent, while others (this is the case of nonthermal tails in the MB distribution) do not make precise predictions about the gamma ray emission, and in fact do not require it. Clusters of galaxies are among the targets for observations by the GLAST satellite. These observations will open a new window onto the nonthermal processes occurring in the intracluster gas, and will allow one to understand the origin of the observed radiation at lower frequency (for a recent review see (Blasi 2002)).  As stressed above, there is at present no compelling evidence in favor of any of the proposed acceleration sites for the nonthermal particles in clusters. Nevertheless, energetic events in the history of a cluster represent good candidates, and cluster mergers, that build up the cluster itself hierarchically, fit the description. It seems therefore reasonable to associate the existence of nonthermal particles to some process occurring during these mergers. This argument is made stronger by the fact that mergers are also thought to be responsible for the heating of the intracluster gas. A possible observational evidence for a correlation between major mergers and radio halos has recently been found by Buote (2001). In particular, the correlation exists between the radio emission at 1.4 GHz and the degree of departure from virialization in the shape of clusters, interpreted as a consequence of a recent or ongoing merger that visibly changed the dark matter distribution in the cluster core.  The shocks that are formed in the baryon components of the merging clusters are able to convert part of the gravitational energy of the system into thermal energy of the gas, as shown by direct observations (e.g. (Markevitch, Sarazin and Vikhlinin 1999)). It has been claimed that if these shocks are strong enough, they can efficiently accelerate particles by first order Fermi acceleration (Fujita \\& Sarazin 2001; Miniati et. al. 2001a,b; Blasi 2000). The consequences of these shocks on the nonthermal content of clusters of galaxies may be dramatic, and deserve to be considered in detail. Both electrons and protons (or nuclei) are accelerated at the shock surfaces during mergers, but the dynamics of these two components is extremely different: high energy electrons have a radiative lifetime much shorter than the age of the cluster, so that they rapidly radiate most of their energy away and eventually pile up at lorentz factors $\\sim 100$. On the other hand, protons lose only a small fraction of their energy during the lifetime of the cluster, and their diffusion time out of the cluster are even larger, so that they are stored in clusters for cosmological times (Berezinsky, Blasi \\& Ptuskin 1997; Volk, Aharonian, \\& Breitschwerdt 1996). In other words, while for high energy electrons only the recent merger events are important to generate nonthermal radiation that we can observe, in order to determine the proton population of a cluster (that can generate secondary electrons) we need to take into account the all history of the cluster.  In this paper we simulate merger histories of galaxy clusters and we calculate the properties of the shocks generated during the merger events and the related spectrum of particles accelerated at the shocks. We also account for the re-energization of particles preexisting within the merging clusters. We demonstrate that major mergers (mergers between clusters with comparable masses) which are supposed to be the most energetic events and that are often thought to be responsible for nonthermal activity, generate shocks that are typically weak and cannot account for the spectral slopes of the observed nonthermal radiation. Shocks in small mergers are also considered, and we find that they may play an important role to generate the nonthermal radiation by primary electrons accelerated recently. As far as the proton component is concerned, the time integrated spectra are energetically dominated by major mergers, so that the resulting spectra are very steep, due to the weakness of the corresponding shocks.  Throughout the paper we assume a flat cosmology ($\\Omega_0=1$) with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=1-\\Omega_m=0.7$ and a value for the Hubble constant of $70\\, \\rm{km/s/Mpc}$.  The paper is organized as follows: in \\S \\ref{sec:simula} we discuss our simulations for the reconstruction of the merger tree; in \\S \\ref{sec:accelera} we discuss the basics of shock acceleration and reacceleration in clusters of galaxies, and the physics of cosmic ray confinement. In \\S \\ref{sec:results} we describe our results, and we conclude in \\S \\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  We investigated the possibility that the nonthermal activity observed from some clusters of galaxies may originate through radiative losses of electrons either accelerated at shocks during cluster mergers or produced as secondary products of the inelastic collisions of protons, in turn accelerated at the same merger shocks.  While the spectrum of protons at any time is the result of all the merger history of a cluster, due to cosmic ray confinement (Berezinsky, Blasi \\& Ptuskin 1997; Volk, Aharonian \\& Breitschwerdt 1996), primary electrons that are able to radiate radio or X-ray photons at present need to be accelerated in very recent times, so that only the last mergers are relevant (Fujita \\& Sarazin 2001).  The merger history of clusters can be simulated by using a PS approach, which allows one to obtain, for each merger event, the mass of the subclusters. In the reasonable assumption of binary mergers, it is also easy to calculate the relative velocity of the two merging clusters, and if clusters are assumed to be virialized structures, also the Mach numbers of the two approaching subclusters. Once the Mach numbers are known, it is possible to calculate the spectra of the particles accelerated by first order Fermi acceleration.  We discussed separately the case of secondary and primary electrons. The secondary electrons, generated in $pp$ scattering have a spectrum that approximately reproduces the spectrum of the parent protons (due to Feynman scaling for the cross section). The spectrum of the protons is calculated by taking into account the acceleration process at each merger, and the reacceleration of protons confined in the merging clusters. The spectrum at the present time is the result of these processes along the all history of the cluster. The spectra that we obtained from our calculations are typically steeper than those required to explain the observed nonthermal radiation. This result is the consequence of the weakness of the shocks associated to major mergers, where the relative motion of the two subclusters occurs at almost free-fall velocity, which implies Mach numbers only slightly larger than unity. Minor mergers produce stronger shocks, but they are energetically subdominant.  Mach numbers larger than those calculated here may be achieved if the merger event occurs in an overdense region, where the infall motion of the two clusters may be dominated by the gravitational well of the surrounding matter rather than by the mutual interaction between the two clusters. This kind of situations can indeed either increase or decrease the Mach numbers compared with the binary case. In \\S \\ref{sec:results} we have found that in order to obtain Mach numbers larger than $\\sim 3$ the overdensity must be such that for rich clusters (mass larger than $5\\times 10^{14}$ solar masses) the probability of sitting in such a potential is pretty slim, and the binary merger model should represent an accurate description of reality, in a statistical sense.  As stressed above, primary electrons can generate observable nonthermal radiation only if accelerated in recent mergers, occurred less than $10^9$ years ago. Our simulations of 500 clusters with mass $10^{15}$ solar masses show that only $\\sim 6\\%$ of them seem to have nonthermal activity with the same spectral features observed in the Coma cluster. This number should be compared with the statistics of radio halos (Feretti et al. 2000) which seems to suggest that $\\sim 30\\%$ of the clusters with X-ray luminosity larger than $10^{45}$ erg/s have such radio halos. This comparison should however be taken with caution. In fact, if the radio halos found by Feretti et al. (2000) have steep spectra, then our statistics increases appreciably and the disagreement may be attenuated. Unfortunately the spectrum is not available for all radio halos. A more detailed analysis of recent mergers accounting in detail for the time dependent electron losses is being currently carried out (Gabici \\& Blasi, in preparation).  We can summarize our conclusions as follows: \\begin{itemize} \\item[1)] the diffuse nonthermal activity should not be directly associated to protons accelerated at merger shocks within the cluster volume; \\item[2)] the nonthermal activity should not correlate directly with major cluster mergers, unless the turbulence induced by mergers is responsible for particle acceleration; if a correlation is confirmed between radio halos and clusters that suffered major mergers, as seems to emerge from the analysis of Buote et al. (2001), then the natural conclusion is that the nonthermal particles are not accelerated at shocks but rather energized by other processes, possibly related to resonant wave-particle interactions; \\item[3)] if electrons directly accelerated at merger shocks are the sources of radio halos and HXR emission, then only about $6\\%$ of the clusters with mass $10^{15}$ solar masses are expected to have such nonthermal activity (note that even in these cases the acceleration is not expected to occur at shocks formed during major mergers). \\end{itemize}"
    },
    "0207/astro-ph0207009_arXiv.txt": {
        "abstract": "\\rightskip 0pt \\pretolerance=100 \\noindent We present sequential optical spectra of the afterglow of GRB 010222 obtained one day apart using the Low Resolution Imaging Spectrometer (LRIS) and the Echellette Spectrograph and Imager (ESI) on the Keck telescopes. Three low-ionization absorption systems are spectroscopically identified at $z_{1}=1.47688$, $z_{2}=1.15628$, and $z_{3}=0.92747$. The higher resolution ESI spectrum reveals two distinct components in the highest redshift system at $z_{1a}=1.47590$ and  $z_{1b}=1.47688$. We interpret the $z_{1b}=1.47688$ system as an absorption feature of the disk of the host galaxy of GRB 010222. The best fitted power-law optical continuum and [Zn/Cr] ratio imply low dust content or a local gray dust component near the burst site. In addition, we do not detect strong signatures of vibrationally excited states of $H_{2}$. If the GRB took place in a superbubble or young stellar cluster, there are no outstanding signatures of an ionized absorber, either. Analysis of the spectral time dependence at low resolution shows no significant evidence for absorption-line variability. This lack of variability is confronted with time-dependent photoionization simulations designed to apply the observed flux from GRB 010222 to a variety of assumed atomic gas densities and cloud radii. The absence of time dependence in the absorption lines implies that high-density environments are disfavored. In particular, if the GRB environment was dust free, its density was unlikely to exceed $n_{H}=$$10^{2}$ cm$^{-3}$. If depletion of metals onto dust is similar to Galactic values or less than solar abundances are present, then $n_{H}$ $\\geq$ 2 $\\times$ $10^{4}$ cm$^{-3}$ is probably ruled out in the immediate vicinity of the burst. ",
        "introduction": "Years after the serendipitous discovery of gamma-ray bursts (GRBs) by the Vela spacecraft (Klebesadel et al. 1973), the nature of the progenitors responsible for generating the initial explosion remains uncertain, while an elegant set of ideas about the production and evolution of the afterglow has developed (M\\'esz\\'aros \\& Rees 1997; Sari \\& Piran 1997).  Aside from recent breakthroughs in X-ray spectroscopy of GRBs (\\eg\\ Piro et al. 2000) a large fraction of the observational information about the environments of GRBs derives from optical spectroscopy and imaging of the host galaxies. Optical spectroscopy of the integrated light and calibrated emission lines has been used to derive the star-formation rate (SFR) for a number of the host galaxies (see for instance Bloom et al. 1998; Djorgovski et al. 1998). As a complement to spectroscopy, high-resolution optical imaging and astrometry have helped pinpoint two-dimensional locations for GRBs with respect to the host galaxy. The distribution of these GRB locations has yielded a number of statistical constraints in the progenitor scenarios (\\eg\\ Bloom, Kulkarni, \\& Djorgovski 2000).  A promising technique that has been applied with less success so far is the study of time dependence in absorption lines from metal ions (Perna \\& Loeb 1998; B\\\"ottcher et al. 1999) or ${\\rm H}_2$ vibrational levels (Draine 2000) that might be excited by the UV radiation generated during the evolution of the burst. The detection of absorption-line variability could provide important clues about the physical dimensions of the photoionized region and the density in the immediate environment of the GRB. Absorption-line variability can be quantified by measuring changes in the equivalent widths  as a function of time or by identifying the appearance of new absorption features. Vreeswijk~et al. (2001) studied the time evolution of the Mg II doublet in the optical spectra of GRB 990510 and GRB 990712 but failed to find any significant changes. In this work we present a study of the time evolution of absorption systems in the optical spectrum of GRB 010222, and discuss possible implications for its progenitor environment. The outline of the paper is as follows: \\S 2 describes the optical spectroscopy,   \\S 3 describes the absorption line identification and continuum fitting. In \\S 4 we detail column density determinations, arguments in favor of the $z=1.47688$ redshift for GRB 010222, kinematics of the host galaxy, and a study of time evolution of absorption lines.  A description of the photoionization code and results is given in \\S 5 and \\S 6. Finally, the implications of our results and conclusions are presented in \\S 7 and \\S 8. ",
        "conclusions": "The use of high-resolution spectroscopy has resolved the host galaxy of GRB 010222 into multiple components. We conclude that the GRB took place in a galactic disk that gives rise to the strongest absorption system. The power-law index of the optical continuum indicates that the dust content at the burst place is low or that the reddening follows a gray dust model. Under these conditions the bulk of the X-ray excess (in 't Zand et al. 2001) can be attributed to inverse Compton scattering.  Low resolution spectroscopy obtained over the span of two days failed to reveal any significant evidence for time dependence of the observed absorption lines. Signatures of strong $H_{2}$ absorption and/or fluorescence are also absent from our spectra. Photoionization models constructed to test a range of initial conditions for this burst show that dense, compact media would be photoionized fairly quickly while less dense media with large radii evolve slower with time. We argue that prior to a GRB, an environment partially ionized either by the progenitor or by nearby sources would prevent the full manifestation of predicted absorption-line evolution. The physical attributes of GRB 010222 are consistent with a low-density molecular cloud, superbubble, or young stellar cluster as the possible environment of the GRB.  In the future some consideration should be given to a prompt optical-UV flash like the one accompanying GRB 990123. The intensity of such an event could in principle lead to dust destruction, but it might also be responsible for the photoionization of a fraction of the neutral gas. The full effect of such a flash has been ignored in our models. Preferentially, the observation of time-dependent absorption lines would work best in bursts occurring in the Earth-facing side of the host galaxy and in compact regions. Our analysis suggests that this technique applied with high-resolution spectroscopy could potentially distinguish the metallicity of progenitor environments. In that regard shortening the interval between burst localization and initial spectroscopy becomes crucial to exploring earlier, less photoionized times. The next generation of high-resolution imagers might directly resolve the environment of GRBs in a relatively nearby host galaxy.  A prospect raised by the observations of GRB 010222 is the study of the properties of galactic dust at high-redshift using bursts as beacons to illuminate the dust nearby. Further exploration of this issue is encouraged as spectroscopy is obtained for a larger sample of GRBs."
    },
    "0207/astro-ph0207186_arXiv.txt": {
        "abstract": "{The determination of the heavy element abundances from giant extragalactic \\ion{H}{2} regions is based on collisionally excited lines. We argue that in the presence of temperature variations the abundances determined are lower limits to the real heavy element abundances. To determine the real abundances it is necessary to take into account the temperature variations present in these nebulae. We discuss the relevance of obtaining accurate line intensities of recombination lines of H, He, C, and O to determine the chemical composition of extragalactic \\ion{H}{2} regions.  We suggest that Pagel's method to derive the O/H ratio should be calibrated by using recombination lines instead of photoionization models or abundances derived from collisionally excited lines. }  \\listofauthors{M.~Peimbert \\& A.~Peimbert} \\indexauthor{Peimbert, M.} \\indexauthor{Peimbert, A.}  \\begin{document} ",
        "introduction": "The chemical abundances of extragalactic \\ion{H}{2} regions are paramount for: a) the study of the heavy elements enrichment of a given galaxy as a function of position, and b) the enrichment of the heavy elements in the universe as a function of time by looking at galaxies at different redshifts.  A powerful tool to determine O/H values is provided by Pagel's method \\cite{pag79} which is based on the [\\ion{O}{2}] and [\\ion{O}{3}] collisionally excited lines, these lines are very strong and are easily detected in objects at different redshifts. This method has to be calibrated by matching the [\\ion{O}{3}]/H$\\beta$ and [\\ion{O}{2}]/H$\\beta$ line intensity ratios with abundances derived from empirical methods or with photoionization models. Pagel's method might become paramount for the study of the heavy elements enrichment of the universe as a function of time.  In this review we discuss briefly some observational evidence in favor of the presence of temperature variations and how large the $t^2$ values are. It has been shown that in the presence of temperature variations the collisionally excited lines of a given element provide a lower limit to the abundance of this element \\cite{pei67}.  We argue that Pagel's method has to be calibrated with recombination lines of \\ion{O}{1} and \\ion{O}{2}. We propose that the recombination lines of H, He, C and O should be observed for a set of extragalactic \\ion{H}{2} regions with different heavy element abundances to establish a proper calibration of Pagel's method.  Recent reviews on the temperature structure of gaseous nebulae are those by Peimbert (1995, 2002), \\scite{pei01}, \\scite{est02z}, Liu (2002a, b), \\scite{sta02}, \\scite{tor02}, and \\scite{pei02w}. ",
        "conclusions": "To be able to constrain the models for the evolution of galaxies as a function of redshift it is crucial to have good determinations of their heavy element abundances, and Pagel's method might be the best tool to determine these abundances.  Some determinations of the O/H values based on recombination lines are already available for giant extragalactic \\ion{H}{2} regions. They yield values in the $8.2 < log \\,\\,{\\rm O/H} + 12 < 8.8$ range. In the near future it will be possible to increase the quality of these determinations and to increase the available range of O/H values.  We propose to calibrate Pagel's method using O recombination lines.  The O recombination abundances are from 2 to 3 times higher than those derived from $R_{23}$ and $T_e(4363/5007)$. GRANTECAN can be used for this calibration. A spectrograph with a resolution higher than 5000 would be needed for this project."
    },
    "0207/astro-ph0207653_arXiv.txt": {
        "abstract": "name{Samenvatting}% \\def\\bibname{Bibliografie}\\def\\chaptername{Hoofdstuk}% \\def\\appendixname{Bijlage}\\def\\contentsname{Inhoudsopgave}% \\def\\listfigurename{Lijst van figuren}\\def\\listtablename{Lijst van tabellen}% \\def\\indexname{Index}\\def\\figurename{Figuur}\\def\\tablename{Tabel}% \\def\\partname{Deel}\\def\\enclname{Bijlage(n)}\\def\\ccname{Ter attentie van}% \\def\\headtoname{Aan}\\def\\headpagename{Pagina}% \\def\\today{\\number\\day\\space\\ifcase\\month\\or januari\\or februari\\or maart\\or% april\\or mei\\or juni\\or juli\\or augustus\\or september\\or oktober\\or% november\\or december\\fi \\space\\number\\year}% \\typeout{ >>>>> use hlatex209 for Dutch hyphenation <<<<< }} \\hyphenation{Schrij-ver Krij-ger Kuij-pers Bal-le-gooij-en}   \\def\\warningoverprint #1{\\special{!userdict begin /bop-hook{gsave 100 600 translate -45 rotate /Times-Roman findfont 100 scalefont setfont 0 0 moveto 0.9 setgray (#1) show grestore}def end}}  \\newcounter{onefig} \\newcounter{fignumber} \\newcount\\nocaptions \\newcount\\nofigures \\newcount\\figwidth \\newcount\\viewgraphs \\def\\figsfile{} \\def\\figspath{} \\def\\paper{} \\def\\format{} \\def\\figlabel{} \\long\\def\\nextfig#1{\\setcounter{figure}{\\value{fignumber}} \\addtocounter{fignumber}{1} \\ifnum \\viewgraphs=1 \\newpage \\pagestyle{empty} \\fi \\ifnum\\value{onefig}=0 #1 \\fi \\ifnum\\value{onefig}=\\value{fignumber} #1 \\fi} \\def\\figwidths#1#2{\\ifnum \\nocaptions=1 #2mm \\else #1mm \\fi} \\def\\paper#1{}  % \\long\\def\\plotfig#1#2{\\ifnum \\nofigures=1 \\else #2 \\fi} \\long\\def\\captiontext#1{\\ifnum \\nofigures=1 \\raggedright \\fi \\ifnum \\nocaptions=1 \\paper \\ifnum \\viewgraphs=0 \\newline  \\mbox{}\\hrulefill\\mbox{} \\newline \\newline label:~\\{\\figlabel\\} \\fi \\else \\ifnum \\nofigures=0 \\fi #1 \\fi}  \\newcount\\panelwidth \\newcount\\panelheight \\newcount\\bxmin \\newcount\\bymin \\newcount\\bxmax \\newcount\\bymax \\newcount\\tbxmin \\newcount\\tbymin \\newcount\\tpanelwidth \\newcount\\tpanelheight \\newcount\\tpdif \\panelwidth=70 \\panelheight=70  % \\def\\panelsize #1,#2;{\\panelwidth=#1 \\panelheight=#2} \\def\\setbb #1,#2;#3,#4;#5,#6;{% \\tbxmin=#1 \\tbymin=#2    % \\bxmin=#3 \\bymin=#4      % \\bxmax=#5 \\bymax=#6}     % \\def\\barepanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,clip=,% width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}} \\def\\labelxpanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelxypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}}  \\newcommand{\\panellabel}[3]{\\makebox[0pt][r]{\\raisebox{#2} #3 \\hspace{#1}}}  \\def\\topfraction{0.99}                      % \\def\\dbltopfraction{0.99} \\def\\textfraction{0.01} \\def\\floatpagefraction{0.99}  \\def\\CC{\\par \\vspace*{-2ex} \\footnotesize \\baselineskip=8pt \\begin{verbatim}} \\def\\EE{\\par \\vspace*{-2ex} \\normalsize}  \\long\\def\\startignore #1\\stopignore{}   %   \\newenvironment{task}{% \\begin{list}{$\\bullet$}{\\itemsep=1ex \\parindent=3ex \\parskip=1ex \\topsep=0ex}} {\\end{list}}  \\def\\koprr{\\list{}{ \\itemindent=0pt \\listparindent=0pt \\itemsep=0pt \\leftmargin=5ex}\\item[]}                     % \\let\\endkoprr=\\endlist  \\def\\kolom{\\list{}{\\rightmargin=5cm      % \\topsep=0ex \\itemsep=0cm \\parsep=0ex \\parskip=0ex \\partopsep=0ex \\listparindent=3ex \\itemindent=0pt \\leftmargin=0pt} \\item[] \\vspace*{-0ex}}                        % \\let\\endkolom=\\endlist  \\renewcommand{\\labelitemi}{{\\bf --}}                % \\def\\setlistparams{ \\topsep=0.7ex                 % \\itemsep=0.7ex                % \\leftmargini=3ex}             % \\setlistparams                  %  \\newcounter{alistindex}  \\newcounter{alistvalue}     % \\newenvironment{alist}{% \\begin{list}{\\alph{alistindex})}{\\usecounter{alistindex}} \\itemsep=1.5ex \\parskip=1ex}{\\end{list}}  \\def\\spacing #1{\\small \\renewcommand{\\baselinestretch}{#1} \\normalsize \\setlistparams}             %  \\newenvironment{itemrr}{\\vspace*{-\\parskip} \\leftmargini=3ex \\topsep=0.6ex \\begin{itemize} \\itemsep=0.5ex \\parindent=0ex \\parsep=0ex}{ \\end{itemize} \\vspace*{-\\parskip}}  \\newenvironment{conditem}{ \\leftmargini=6ex \\topsep=0.2ex \\begin{itemize} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\newenvironment{enumerr}{ \\leftmargini=3ex \\topsep=1.5ex \\begin{enumerate} \\itemsep=0.7ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newenvironment{condenum}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{enumerate} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newcounter{romenumnr} \\newenvironment{romanenum}{ \\leftmargini=5ex \\topsep=1.5ex \\begin{list}{(\\roman{romenumnr}) --}{\\usecounter{romenumnr} \\itemsep=1ex \\listparindent=3ex \\parsep=0ex}}{ \\end{list}}  \\newcommand{\\posterpage}[2]{{\\Large \\sf \\newpage \\setlength{\\minipagewidth}{\\textwidth} \\addtolength{\\minipagewidth}{-2cm} \\fboxrule=0.3ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[\\textwidth]{ \\begin{minipage}[t]{\\minipagewidth} \\parskip=3ex \\vspace*{1ex} \\begin{center} \\LARGE \\bf #1 \\end{center} \\vspace{2ex} #2 \\vspace*{3ex} \\mbox{ } \\end{minipage}} \\hspace*{\\textwidth} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.2ex} \\framebox[\\textwidth]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{\\textwidth}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}}  \\newcommand{\\vwgraph}[3]{\\newpage \\begin{center} \\shadowframe{#1}{ \\begin{center} \\vspace*{2ex} {\\Large \\bf #2}\\\\[4ex] \\end{center} {\\Large \\sf #3}}   % \\end{center}} \\newenvironment{itemvg}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{itemize} \\vspace{-1ex} \\itemsep=0.5ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\long\\def\\lectvw#1#2{\\newpage \\vspace*{-15mm} \\begin{center} \\begin{center} {\\LARGE \\sc \\underline{#1}}\\\\[4ex] \\end{center} {\\Large \\sf #2}   % \\end{center}}  \\newlength{\\minipagewidth} \\newcommand{\\shadowframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{1.0ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{4.5ex} \\rule{#1}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newsavebox{\\boxcontent} \\newcommand{\\ovalhead}[1]{ \\unitlength=1cm \\sbox{\\boxcontent}{\\mbox{~~{#1}~~}} \\begin{center} \\ifdim\\wd\\boxcontent>6ex \\ifdim\\wd\\boxcontent<8cm \\begin{picture}(8,3) \\thicklines \\put(4.0,0.8){\\oval(8,1.6)} \\put(0.0,0.7){\\parbox{8cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\ifdim\\wd\\boxcontent<12cm \\begin{picture}(12,3) \\thicklines \\put(6.0,0.8){\\oval(12,1.6)} \\put(0.0,0.7){\\parbox{12cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\begin{picture}(16,3) \\thicklines \\put(8.0,0.8){\\oval(16,1.6)} \\put(0.0,0.7){\\parbox{16cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\fi \\fi \\fi \\end{center}}  \\newcommand{\\colloqframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\parskip=3ex \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{#1}{2ex} \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newcommand{\\soloval}[1]{\\fboxrule=0.5mm \\unitlength=1cm \\hspace*{5mm} \\sol{20mm}\\\\[-20mm] \\begin{picture}(15,3) \\thicklines \\put(7.5,2){\\oval(15.60,3.60)} \\put(7.5,2){\\oval(15.55,3.55)} \\put(7.5,2){\\oval(15.50,3.50)} \\put(7.5,2){\\oval(15.45,3.45)} \\put(7.5,2){\\oval(15.40,3.40)} \\put(7.5,2){\\oval(15.35,3.35)} \\put(7.5,2){\\oval(15.00,3.00)} \\put(2.0,1.9){\\parbox{13cm}{#1}} \\end{picture}}  \\setcounter{secnumdepth}{3} \\setcounter{tocdepth}{3}  \\def\\chapterrr#1{\\chapter{#1} \\label{chap:#1} \\vskip-\\parskip} \\def{ A detailed spectroscopic study of 11 giants with spectral type from G9 to M2 is presented. The 2.38 -- 4.08\\,\\mic\\ wavelength-range of band 1 of ISO-SWS (Short-Wavelength Spectrometers on board of the Infrared Space Observatory) in which many different molecules --- with their own dependence on each of the stellar parameters --- are absorbing, enables us  to estimate the effective temperature, the gravity, the microturbulence, the metallicity, the CNO-abundances, the \\cc-ratio and the angular diameter from the ISO-SWS data. Using the Hipparcos' parallax, the radius, luminosity and gravity-inferred mass are derived. The stellar parameters obtained are in good agreement with other published values, though also some discrepancies with values deduced by other authors are noted. For a few stars ($\\delta$ Dra, $\\xi$ Dra, $\\alpha$ Tuc, H Sco and $\\alpha$ Cet) some parameters --- e.g.\\ the CNO-abundances --- are derived for the first time. By examining the correspondence between different ISO-SWS observations of the same object and between the ISO-SWS data and the corresponding synthetic spectrum, it is shown that the relative accuracy of ISO-SWS in band 1 (2.38 -- 4.08\\,\\mic) is better than 2\\,\\% for these high-flux sources. The high level of correspondence between observations and theoretical predictions, together with a confrontation of the estimated \\teff(ISO) value with \\teff\\ values derived from colours --- which demonstrates the consistency between $V-K$, BC$_K$, \\teff\\ and \\ad\\ derived from optical or IR data --- proves that both the used MARCS models to derive the stellar quantities and the flux calibration of the ISO-SWS detectors have reached a high level of reliability. ",
        "introduction": "In the fourth article of this series, we discuss the ISO-SWS (Short-Wavelength Spectrometers on board of the Infrared Space Observatory) spectra of 11 giants cooler than the Sun. These giants are part of a larger sample selected as stellar standard candles for the calibration of the detectors of ISO-SWS. In \\citet[][hereafter referred to as Paper~I]{Decin2000A&A...364..137D} a method was developed to analyse the infrared data (2.38 -- 12\\,\\mic) of cool stars: the specific dependence of the absorption pattern by the different molecules gives the possibility to estimate fundamental stellar parameters like the effective temperature (\\teff), the gravity ($\\log$ g), the microturbulence (\\vt), the metallicity, the CNO-abundances, the \\cc-ratio and the angular diameter (\\ad). We now apply this method of analysis to the ISO-SWS data of these 11 giants. Keeping the general calibration problems and the accuracy of the theoretical models and synthetic spectra  (based on the MARCS and Turbospectrum code \\citep{Gustafsson1975A&A....42..407G, Plez1992A&A...256..551P, Plez1993ApJ...418..812P}; version May 1998) in mind --- as discussed in \\citet[][hereafter referred to as Paper~II]{Decin2000b} --- the error bars on the different stellar parameters are determined and the resultant stellar parameters are confronted with other published parameters.  This paper is organised in the following way: in Sect.\\ \\ref{summary} a summary of the general discrepancies between the ISO-SWS spectra of the cool giants and the theoretical spectra, as described in Paper~II, is given. We will go more deeply into the determination of the different stellar parameters and their error bar in Sect.\\ \\ref{stelparameters}. After an overview of the deduced stellar parameters, each star is discussed in more detail. Not only are the accuracy and precision of the ISO-SWS data assessed, but the deduced stellar parameters are also evaluated with respect to other published values. A discussion on some of the stellar parameters obtained is held in Sect.\\ \\ref{discussion}, where we focus on (i) the effective temperature and (ii) the \\cc\\ ratio.  The conclusions are given in Sect.\\ \\ref{conclusions}.  The appendix of this article is published electronically. Most of the grey-scale plots in the article are printed in colour in the appendix, in order to better distinguish the different spectra or symbols. ",
        "conclusions": "\\label{conclusions}  In this paper, the atmospheric parameters of 11 giants have been determined by using the method described in \\citetalias{Decin2000A&A...364..137D} and \\citetalias{Decin2000b}. A confrontation with other published stellar parameters, other ISO-SWS data and the synthetic spectra always showed very `positive' results both for the theoretical modelling and for the calibration of the ISO-SWS detectors. The very small discrepancies still remaining in band 1 are at the 1 -- 2\\,\\% level for the giants, proving not only that the calibration of the (high-flux) sources in this band is better than $\\sim 2$\\,\\%, but also that the description of cool-star atmospheres and molecular linelists is already quite accurate at the ISO-SWS spectral resolution ($R \\sim 1000$).  In the discussion of the deduced parameters, we demonstrated that \\teff(ISO) is consistent with other \\teff($(V-K)_0$) or \\teff($K_0$, BC$_K$, \\ad(ISO)) values in within the few percent level, proving once more that the used MARCS-models to derive the quantities from the ISO-SWS spectra are very trustworthy and that the flux calibration of ISO-SWS is indeed reassuring.  The obtained \\cc\\ ratios are compatible with other sources \\citep[e.g.\\ ][]{Charbonnel1995ApJ...453L..41C}.  Like already pointed out in \\citetalias{Decin2000b}, these synthetic spectra will now be used for the --- latest --- OLP10 calibration of ISO-SWS. Also other consortia constructing instruments for ground-based telescopes --- i.e.\\ MIDI (\\,=\\, the Mid-Infrared Interferometric instrument for the VLTI with first observations foreseen in 2001) and new satellites --- i.e. SIRTF, the Space InfraRed Telescope Facility, with launch foreseen in December 2001, and Herschel, a far-infrared submillimeter telescopes with launch foreseen in 2007 --- will use this type of study and/or these SEDs for determining the required sensitivity of their instruments. A confrontation with existing SEDs, like the ones of M.\\ Cohen, is therefore planned in a forthcoming paper of this series, in which also the new atomic linelist constructed by Sauval \\citep{Sauval2000} will be implemented."
    },
    "0207/astro-ph0207379_arXiv.txt": {
        "abstract": "{ Due to the complexity of their structure, the theoretical study of interstellar clouds must be based on three-dimensional models. It is already possible to estimate the distribution of equilibrium dust temperature in fairly large 3D models and, therefore, also to predict the resulting far-infrared and sub-mm emission. Transiently heated particles introduce, however, a significant complication and direct calculation of emission at wavelengths below 100$\\mu$m is currently not possible in 3D models consisting of millions of cells. Nevertheless, the radiative transfer problem can be solved with some approximations. We present a numerical code for continuum radiative transfer that is based on the idea of a `library' describing the relation between the intensity of the local radiation field and the resulting dust emission spectrum. Given this mapping it is sufficient to simulate the radiation field at only a couple of reference wavelengths. Based on the library and local intensities at the reference wavelengths, the radiative transfer equation can be integrated through the source and an approximation of the emission spectrum is obtained. Tests with small models for which the radiative transfer problem can be solved directly show that with our method, one can easily obtain an accuracy of a few per cent. This depends, however, on the opacity of the source and the type of the radiation sources included. As examples we show spectra computed from three-dimensional MHD simulations containing up to 128$^3$ cells. The models represent starless, inhomogeneous interstellar clouds embedded in the normal interstellar radiation field. The intensity ratios between IRAS bands show large variations that follow the filamentary structure of the density distribution. The power law index of the spatial power spectrum of the column density map is -2.8. In infrared maps temperature variations increase the power at high spatial frequencies, and in a model with average visual extinction $\\langle A_{\\rm V} \\rangle\\sim$10 the power law index varies between -2.5 and -2.7. Assuming constant dust properties throughout the cloud, the IRAS ratio $\\langle I_{\\rm 60}/I_{\\rm 100}\\rangle$ decreases in densest cores only by a factor of $\\sim$4 compared with the value in diffuse medium. Observations have shown that in reality the ratio can decrease twice as much even in optically thinner clouds. This requires that most of the small grains are removed in these regions, and possibly a modification of the properties of large grains. ",
        "introduction": "Present knowledge on the large scale infrared emission from interstellar dust is based largely on the all-sky surveys performed by the IRAS and the COBE satellites. IRAS was the first to observe the sky at a resolution of a few arc minutes at wavelengths between 12 and 100$\\mu$m. Clear variations were observed in the infrared spectrum of interstellar clouds, e.g. the ratio of 60$\\mu$m vs. 100$\\mu$m dropping towards dense clouds (e.g. Laureijs et al. \\cite{laureijs96}; Abergel et al. \\cite{abergel94}). These were interpreted as the result of abundance variations between classical, large grains and very small dust particles that are transiently heated above the equilibrium temperature (e.g. D\\'esert et al. \\cite{desert90}). Emission in the mid-infrared and especially the IRAS 12$\\mu$m band, are now believed to be caused by even smaller grains or very large molecules, the most common candidates being Polycyclic Aromatic Hydrocarbons (PAHs) (L\\'eger et al. \\cite{leger89}; D\\'esert et al. \\cite{desert90}; Li \\& Draine \\cite{li01}). The PAH emission traces the warmer outer surfaces of the clouds.  The COBE/DIRBE instrument covered a wider wavelength range from near-infrared up to 240\\,$\\mu$m and the FIRAS instrument up to 1\\,cm. The spatial resolution was, however, comparatively poor i.e. $\\sim$40$\\arcmin$ for DIRBE and $\\sim$7$\\degr$ for FIRAS. Nevertheless, the extended wavelength coverage revealed surprisingly large quantities of cold dust that were not observed with IRAS. Detailed analysis of the COBE data has shown that at least two emission components are needed to explain the observed far-infrared spectra. For most of the diffuse regions the dust emission follows a modified black body law with temperature $T\\sim$17.5\\,K (Boulanger et al. \\cite{boulanger96}; Dwek et al. \\cite{dwek97}). Lagache et al. (\\cite{lagache98}) found, however, a colder component ($T\\sim15$\\,K) associated with molecular clouds. Finkbeiner et al. (\\cite{finkbeiner99}) reached similar conclusions but included an even colder, $T\\sim$9\\,K, component.  Recent observations with the PRONAOS balloon borne observatory have confirmed the lowering of the colour temperature towards molecular clouds and, e.g. in the case of the Polaris flare ($A_{\\rm V}\\sim$1), towards moderately dense regions (Stepnik et al. \\cite{stepnik_esa}). Similar effects have been seen in ISOCAM studies towards cirrus like clouds ($A_{\\rm V}\\sim$0.5; Miville-Desch\\^enes et al. \\cite{miville02}) where the temperature variations caused by extinction are quite insignificant. The observed increase in the sub-mm emission can be explained by grain growth and the formation of large dust aggregates (Stepnik et al. \\cite{stepnik_esa}; Cambr\\'esy et al. \\cite{cambresy01}) and in dense regions the very small grains seem to have disappeared almost completely. Conversely, in the more diffuse medium, large quantities of small grains can be produced by grain shattering (see e.g. Miville-Desch\\^enes et al. \\cite{miville02}).  Based on these observations and theoretical studies we have obtained a rough picture of the dust properties and infrared emission from interstellar clouds. In the outer layers we have the PAH emission and the warm, very small dust grains. Deeper in the clouds the extinction of short wavelength radiation reduces mid-infrared emission. At the same time, the grains start to form larger aggregates and/or acquire mantles as gas molecules freeze onto them. This increases the far-infrared and sub-mm emissivity relative to the absorption at short wavelengths and the physical grain temperatures also drop. The colour temperature therefore drops to approximately 10\\,K (or possibly even lower) in dense cloud cores (e.g. Juvela et al.~\\cite{juvela02a}).  Theoretical and laboratory studies have provided us with some information on the properties of likely dust grain materials. Together with radiative transfer calculations, these have made it possible to determine size distributions for the different dust components consistent with the emission properties (e.g. D\\'esert et al. \\cite{desert90}; Li \\& Draine \\cite{li01}). On the other hand, radiative transfer models can be used to study the variations of these properties in individual objects. These previous examples show that the dust emission spectrum is variable over the whole range from near-infrared to sub-mm. The changes are related to the local radiation field and the local density. So far, modelling has concentrated either on the determination of the dust emission under given radiation conditions, or at most in one-dimensional, spherical models divided into some tens of cells.  In the related field of the modelling of molecular line emission, three-dimensional models have already been used for some time (Park \\& Hong \\cite{park95}; Park et al. \\cite{park96}; Juvela \\cite{juvela98}; Padoan et al. \\cite{padoan98}; Juvela et al. \\cite{juvela01}). In that field, inhomogeneous source structure has an even stronger effect on the emerging radiation and, on the other hand, direct calculations are already possible for models with up to $\\sim$100$^3$ cells or, with approximate methods, even higher (Ossenkopf \\cite{ossenkopf02}).  The continuum emission is not as directly linked with the local, physical conditions. It is, nevertheless, affected by the variations in the radiation field (extinction of the external radiation field and internal sources) and the changes in the dust properties. So far dust emission has not been studied with general three-dimensional cloud models. The main reason is the complexity caused by the transiently heated small dust particles. In order to be able to predict the emission accurately, the distribution of the grain temperatures must be first determined in each computational cell. Each dust population must be discretized into separate grain size intervals. For each of these, a couple of hundred enthalpy bins are needed and the solution of the associated set of linear equations takes typically several seconds. This becomes significant when the model contains millions of cells and, in practice, it limits direct calculations to one- or possibly two-dimensional models.  We will discuss how this limitation can be overcome by suitable approximations without siginificantly affecting the accuracy of the results. This will be based on the assumption that the local spectral energy distribution can be deduced from the intensity at suitably selected reference wavelengths. One must first determine a mapping between these intensities and the resulting local spectral energy distribution. The radiative transfer problem needs only to be solved at the reference wavelengths, and by using the established mapping, one can solve the radiative transfer problem. The solution is approximative but can be be made sufficiently accurate within the limitations of present day computers. The method makes it possible to study explicitly the effects that the cloud structure has on the infrared emission, and the method could be useful e.g. in the simulation of the spectral energy distribution of galaxies.  We first explain the implementation of the radiative transfer program (Sect.~\\ref{sect:program}) and discuss its accuracy (Sect.~\\ref{sect:tests}). In Sect.~\\ref{sect:models} we apply the method to three-dimensional model clouds that are based on numerical simulations of supersonic magneto-hydrodynamic (MHD) turbulence and we show some observable consequences of the inhomogeneity of the clouds. Finally, in Sect.~\\ref{sect:adjust} we discuss some qualitative effects of spatial variations in the dust properties. ",
        "conclusions": "We have presented a new approximate method that can be used to compute the infrared dust emission from large, three dimensional cloud models consisting of millions of computational cells. The method is based on simple discretization of the incoming intensity using a few reference wavelengths. We have demonstrated that relative accuracy of a few percent is easily reached, even when transiently heated grains are included. Such accuracy is hard to match either by observations or the Monte Carlo sampling that is often used to estimate the intensity of the radiation field. We have tested the method using model clouds with average visual extinctions up to $A_{\\rm V}\\sim$100. However, with use of more reference wavelengths and/or higher discretization, the method can be applied to even denser clouds, and to models with several heating sources with different emission spectra.  We have used our new approximate method to compute dust emission, from near-infrared to sub-mm wavelengths, from models based on MHD simulations with supersonic and super-Alfv\\'enic turbulence. The spatial discretization was found to have little impact on the derived distributions of the intensity ratios between IRAS bands. On the other hand, spatial averaging will always reduce the range of intensities. High resolution calculations are therefore necessary, especially when the computed infrared maps are used for statistical studies or for direct comparison with observations.  The spatial power spectra at different wavelengths generally show little variation. The shorter wavelengths have, however, slightly higher relative power at higher spatial frequencies. For the selected MHD model, the power spectra follow approximately a power law with an exponent between -2.5 and -2.7, while the power spectrum of the underlying column density map has a slope of -2.8.  We qualitatively studied the effect of density dependent dust size distributions on the observable infrared maps. By reducing the relative abundance of very small grains by 50\\% in regions with densities $n>10^4$\\,cm$^{-3}$, the ratio $I_{\\rm 60}/ I_{\\rm 100}$ could be brought down to $\\sim$0.05 in dense regions where visual extinction is $A_{\\rm V}\\sim$10. In observations, the colour ratio can reach lower values even in less opaque clouds. This indicates that dust properties must undergo even larger changes, with the removal of most of the small grains and the growth of larger ones.    The numerical method presented in this work will be very useful to constrain dust properties by comparison with new observational data. Dust column density maps of large molecular cloud complexes can now be obtained from stellar extinction measurements based on the ``Two Microns All Sky Survey'' (2MASS). Padoan et al. (\\cite{padoan2002}) have recently generated extinction maps of the whole Taurus molecular clouds complex using 2MASS data, with a spatial resolution comparable or better than that of 100~$\\mu$m IRAS images, spanning a range of visual extinction from $A_{\\rm V}\\approx 0.3$~mag to $A_{\\rm V}\\approx 30$~mag. The comparison of this type of dust column density maps with IRAS images and, in the near future, with SIRTF-MIPS images will provide new observational constraints over a large range of column density and gas density values. Detailed three dimensional calculations of dust emission from physical models of interstellar clouds will be necessary for reproducing the observational data and improve our knowledge of the properties of interstellar dust."
    },
    "0207/hep-ph0207199_arXiv.txt": {
        "abstract": ": We present analytic and numerical results for the evolution of currents on superconducting strings in the classical $U(1) \\times U(1)$ model. We derive an energy functional for the currents and charges on these strings, establishing rigorously that minima should exist in this model for loops of finite size (vortons) if both charge and current are present on the worldsheet. We then study the stability of the currents on these strings, and we find an analytic criterion for the onset of instability (in the neutral limit).  This limit specifies a lower maximal current than previous heuristic estimates. We conclude with a discussion of the evolution of loops towards their final vorton state in the model under consideration. ",
        "introduction": "Topological defects are a class of exact solutions in field theories whose stability is enforced by topological reasons. In particular, strings, the class of defects associated with a non-trivial first homotopy group of the vacuum manifold, have been widely studied, since they seem to appear in a variety of generalisations of the Standard Model (GUTs, SUSY etc.). They are for example associated with the spontaneous symmetry breakdown of a $U(1)$ symmetry, like the axion or baryon symmetry (global), or electromagnetism (local) in superconductivity. Strings could therefore appear during a cosmological phase transition, and are prime candidates for a number of astrophysical puzzles, like the dark matter of the universe or the origin of the most energetic cosmic rays (for a review of cosmic defects, see ref.~\\cite{VS:CS}).  In this paper, we shall be interested in a class of defects where the string's field is coupled to another scalar field, which allows the build up of charge and currents on the worldsheet. Because of the non-dissipative properties of the currents on these strings, they are called superconducting.  The plan of this paper is as follows: in the next section, we will discuss the Lagrangian under consideration, and show how the amount of charge and currents is limited; in the following part, we will carry out an analysis of the energy functional of the condensate, to get a very simple analytic expression for it. We will then use our results to discuss the possibility of forming stable loops of superconducting strings. Finally, we will study the stability of currents on the string's worldsheet and derive an exact analytic result for the onset of instability, which we shall illustrate by numerical simulations obtained from our full 3D field theory code. ",
        "conclusions": "In this paper we have studied the behaviour and effects of the currents and charges on superconducting strings and loops. We have derived exact formulae for the energy of a loop, which exhibits a generic divergence at $\\omega^2 - k^2 = 0$, therefore proving that the chiral case is not an attractor, but rather a repeller.  Since for stable configurations we typically have $\\Sigma \\simeq \\Sigma_{QN}\\equiv Q/N $, from (\\ref{L2}) we can expect the final size of the loop to be very much smaller than its original size (for realistic initial conditions). Of course, we have not taken into account the vortex-antivortex interactions on scales small compared to the string width, so the inevitability of vorton formation is subject to this important caveat. In addition we conclude that springs and Q-loops are not allowed in this theory.  Studying the stability of the current on a straight string, we have also seen that the superconducting regime is unstable when the winding is too high. By Lorentz-invariance, it is easy to see that this will be the case if $m_o^2 v = k^2 - \\omega^2 > k_{\\rm inst}^2$, where $m_o^2 \\simeq k^2_c \\simeq \\omega_c^2$. (We believe that the gauged case will exhibit the same generic features as these global currents, though with the small quantitative differences already discussed.)  Now, from our analysis of the vorton state, we know that equilibrium typically will be achieved when $\\Sigma \\simeq \\Sigma_{QN}$. Using our ansatz (\\ref{Sigmaf}), this is equivalent to \\begin{equation} v \\simeq \\frac{\\Sigma_{QN} - \\Sigma_o}{\\Sigma_{QN} + \\Sigma_o} \\, . \\label{conc} \\end{equation} To ensure stability, we have to impose $v > -k_{\\rm inst}^2 / m_o^2$ which, with (\\ref{conc}) and some algebra, leads to the condition \\begin{equation} \\Sigma_{QN} > \\frac{m_o^2 - k_{\\rm inst}^2}{m_o^2 + k_{\\rm inst}^2} \\,\\Sigma_o \\, . \\label{conc2} \\end{equation} Loops that do not satisfy (\\ref{conc2}) will be unstable and lose quanta of winding. Hence, $\\Sigma_{QN}$ will increase, and the loop may reach a stable state. This process is associated with energy radiation, which would be interesting to quantify to determine possible observational signatures of this phenomenon.  Finally, the perturbative stability analysis carried out for superconducting currents in the magnetic regime has been extended to the chiral and electric cases, and appears to establish their stability. However, this issue is more subtle, and the analysis will be published separately \\cite{LS:DCS}.  \\begin{ack}  Y.L. and E.P.S. gratefully acknowledge very fruitful conversations with Jose Blanco-Pillado; Y.L. is especially grateful to him for pointing out ref.~\\cite{LA:CINST}. The numerical code employed here was originally developed with Jonathan Moore \\cite{JM:PHD,MSM:CS} to whom we are greatly indebted.  Y.L. is supported by EPSRC, the Cambridge European Trust and the Cambridge Newton Trust. This work was also supported by PPARC grant no. PPA/G/O/1999/00603.  Numerical simulations were performed on the COSMOS supercomputer, the Origin3800 owned by the UK Computational Cosmology Consortium, supported by Silicon Graphics Computer Systems, HEFCE and PPARC.  \\end{ack}"
    },
    "0207/astro-ph0207608_arXiv.txt": {
        "abstract": "\\vskip 24pt  Isolated neutron stars undergoing non-radial oscillations are expected to emit gravitational waves in the kilohertz frequency range.  To date, radio astronomers have located about 1,300 pulsars, and can estimate that there are about $2 \\times 10^8$ neutron stars in the galaxy.  Many of these are surely old and cold enough that their interiors will contain matter in the superfluid or superconducting state.  In fact, the so-called glitch phenomenon in pulsars (a sudden spin-up of the pulsar's crust) is best described by assuming the presence of superfluid neutrons and superconducting protons in the inner crusts and cores of the pulsars.  Recently there has been much progress on modelling the dynamics of superfluid neutron stars in both the Newtonian and general relativistic regimes.  We will discuss some of the main results of this recent work, perhaps the most important being that superfluidity should affect the gravitational waves from neutron stars (emitted, for instance, during a glitch) by modifying both the rotational properties of the background star and the modes of oscillation of the perturbed configuration.  Finally, we present an analysis of the so-called zero-frequency subspace (i.e.~the space of time-independent perturbations) and determine that it is spanned by two sets of polar (or spheroidal) and two sets of axial (or toroidal) degenerate perturbations for the general relativistic system. As in the Newtonian case, the polar perturbations are the g-modes which are missing from the pulsation spectrum of a non-rotating configuration, and the axial perturbations should lead to two sets of r-modes when the degeneracy of the frequencies is broken by having the background rotate. ",
        "introduction": "Jacob Bekenstein is one of those rare individuals who can make significant, original contributions to diverse areas of theoretical physics.  He is also a man of great integrity and, I believe, has a humility that serves him well in advising and supporting students and young scientists.  I am profoundly grateful that fate allowed me to be one of those young scientists and now lets me participate in this celebration of his career. One of the areas of theoretical physics that Jacob has worked on is relativistic fluid dynamics.  This is an important component of my current area of research, which is to develop models of Newtonian and general relativistic superfluid neutron stars.  My original interest in superfluids, appropriately enough, was sparked by Jacob, when he suggested that I look at superfluid analogs of effects predicted for quantum fields in curved spacetimes (the Hawking and Fulling-Davies-Unruh effects).  My current interest in superfluids is to determine how the dynamics of superfluid neutron stars differ from their ordinary, or perfect, fluid counterparts and if the different dynamics can lead to observable effects in gravitational waves.  In the remainder of this article, I will give an overview of what my collaborators and I have accomplished so far, including some new results (from work with Nils Andersson) on the structure of the so-called zero-frequency subspace (i.e.~the space of time-independent perturbations).  The main purpose is to show that superfluidity in neutron stars should affect their gravitational waves in two ways, by modifying the rotational properties of the background star and the modes of oscillation of the perturbed configuration.  While there are many mysteries about neutron stars that remain to be explained, we do have some significant observational facts to work with.  For instance, Lorimer \\cite{L01} reports that nearly 1300 pulsars (i.e.~rotating neutron stars) have now been observed.   By extrapolating the data on the local population, he can estimate that there are about $1.6 \\times 10^5$ normal pulsars and around $4 \\times 10^4$ millisecond pulsars in our galaxy.  Of course, there are also neutron stars that are no longer active pulsars.  To get a handle on their number Lorimer takes the observed supernova rate, which is about 1 per 60 years, and the age of the universe to find about $2 \\times 10^8$ neutron stars in the galaxy.  The overwhelming majority of these objects must be very cold in the sense that their (local) temperatures are much less than the (local) Fermi temperatures of the independent species of the matter.  One can estimate the Fermi temperature to be about $10^{12}~{\\rm K}$ for neutrons at supra-nuclear densities, and it is generally accepted that within the first year (and probably much sooner than that) nascent neutron stars should cool to temperatures less than $10^9~{\\rm K}$.  This is an interesting fact, in that nuclear physics calculations of the transition temperature for neutrons and protons to become superfluid and superconducting, respectively, consistently yield a value that is $10^9~{\\rm K}$ in order of magnitude (for recent reviews see \\cite{UL99,LS00}). Thus we can expect that a significant portion of the neutron stars in our galaxy will have at least two (and perhaps more) superfluids in their cores.  In addition to nuclear physics theory and experiment, the well-established glitch phenomenon in pulsars (e.g.~Vela and Crab) \\cite{RM69,L93} is perhaps the best piece of evidence that supports the existence of superfluids in neutron stars. A glitch is a sudden spin-up of the observed rotation rate of a neutron star, and can have a relaxation time of weeks to months \\cite{RD69}.  Baym et al \\cite{BPPR69} have noted that a relaxation mechanism based on ordinary fluid viscosity would be much too short to explain a weeks to months timescale and so they argue that this signals the presence of a neutron superfluid.  Now, a mainstay idea for explaining glitches is that of superfluids and their vortex dynamics, i.e.~how the vortices get pinned, unpinned, and then repinned \\cite{AI75,AAPS84a,AAPS84b} to nuclei in the inner crusts of the glitching pulsars.  This is known as the vortex creep model and in it glitches are a transfer of momentum via vortices from one angular momentum carrying component of the star to another.  The model has worked well to describe both the giant glitches in Vela and the smaller ones of the Crab. The vortex creep model can also be used to infer the internal temperature of a glitching pulsar, and for Vela it implies a temperature of $10^7~{\\rm K}$ \\cite{AAPS84b}.  It is also interesting to note the work of Tsakadze and Tsakedze \\cite{TT80} who have experimented with rotating superfluid Helium II and find behaviour very much like glitches in pulsars.  The classic description of superconductivity in ordinary condensed matter systems is based on the so-called ``BCS'' mechanism (see, for instance, \\cite{TT86} or \\cite{G85} for excellent presentations): the particles that become superconducting must be fermions, and below a certain transition temperature there must be an (usually effective) attractive interaction between them (at the Fermi surface with zero total momentum).  The interaction leads to so-called Cooper-pairing where a pair of fermions act like a single boson and a collection of them can behave as a condensate.  The mechanism is very robust, which is why it also forms the basis for discussion of nucleon superfluidity and superconductivity \\cite{S89}; i.e.~nucleons are fermions and the effective interaction between them at nuclear and supra-nuclear densities can be attractive.  For instance, it is known experimentally that the lowest excited states in even-even nuclei are systematically higher than other nuclei because of pairing between nucleons which must be broken \\cite{TT86}.  After many years of development, beginning with the work of Migdal \\cite{M59}, a consistent picture has emerged (in part, from gap calculations \\cite{UL99,LS00}): At long-range the nuclear force is attractive and leads to neutron ``Cooper'' pairing in ${}^1{\\rm S}_0$ states in the inner crust, but because of short-range repulsion in the nuclear force and the spin-orbit interaction neutrons pair into ${}^3{\\rm P}_2$ states in the more dense regions of the core \\cite{HGRR70}.  In the crust protons are locked inside of neutron rich nuclei embedded in a degenerate normal fluid of electrons.  In the inner crust the nuclei are also embedded in, and even penetrated by, the superfluid neutrons.  In the core, however, the nuclei have dissolved and the protons remain dilute enough that they feel only the long-range attractive part of the nuclear force and pair in ${}^1{\\rm S}_0$ states.  There is no pairing between neutrons and protons anywhere in the core since their respective Fermi energies are so different.  The core superfluid neutrons and superconducting protons are also embedded in a highly degenerate normal fluid of electrons.  Other possibilities, such as pion or hyperon condenstates, have been put forward but we will keep to the simplest scenario that considers only superfluid neutrons, superconducting protons, crust nuclei, and normal fluid electrons.  There are several ways in which the dynamics of a superfluid differ from its ordinary fluid counterpart, and each difference should have some impact on the gravitational waves that a superfluid neutron star emits.  One key difference is that a pure superfluid is locally irrotational. A superfluid, however, can mimic closely ordinary fluid rotation by forming a dense array of (quantized) vortices. In the core of each vortex the superfluidity is destroyed and the particles are in an ordinary fluid state, and can carry non-zero vorticity.  A second, very important difference is when there are several species of matter in a superfluid, or superconducting, state.  The superfluids of all the species will interpenetrate and each superfluid will be dynamically independent having its own unit four-vector and local particle number density.  Lastly, superfluids have zero viscosity, but when vortices and excitations are present, then dissipative mechanisms can exist.  For instance, the scattering of excitations off of the normal fluid in the vortex cores can lead to dissipative momentum exchange between the excitations and the superfluid, the net effect being that the superfluid motion becomes dissipative.  This form of dissipative mechanism is known as mutual friction.  In neutron stars there is a very efficient form of mutual friction \\cite{ALS84,AS88,S89} that depends on the entrainment effect \\cite{AB75,VS81}, which Sauls \\cite{S89} describes as follows: even though the neutrons are superfluid and the protons are superconducting both will still feel the long-range attractive component of the nuclear force.  In such a system of interacting fermions the resulting excitations are quasiparticles.  This means that the bare neutrons (or protons) are ``dressed'' by a polarization cloud of nucleons comprised of both neutrons and protons.  Since both types of nucleon contribute to the cloud the momentum of the neutrons, say, is modified so that it is a linear combination of the neutron and proton particle number density currents.  The same is true of the proton momentum.  Thus when one of the nucleon fluids starts to flow it will, through entrainment, induce a momentum in the other fluid.  Alpar et al \\cite{ALS84} have shown that the electrons track very closely the superconducting protons (because of electromagnetic attraction).  Around each vortex is a flow of the superfluid neutrons.  Because of entrainment, a portion of the protons, and thus electrons too, will be pulled along with the superfluid neutrons.  The motion of the plasma leads to magnetic fields being attached to the vortices. The mutual friction in this case is the dissipative scattering of the normal fluid electrons off of the magnetic fields attached to the vortices.  There has been much effort put forward to develop Newtonian \\cite{ML91,M91,LM94,LM95,P02} and general relativistic formalisms \\cite{C85,C89,CL93,CL94,CL95,CL98a,CL98b,LSC98,CLS00} for describing superfluid neutron stars.  In the simplest, but still physically interesting, formalism one has a system that consists of two interpenetrating fluids---the superfluid neutrons in the inner crust and core and the remaining charged constituents (i.e.~crust nuclei, core protons, and crust and core electrons) that will be loosely referred to as ``protons''---and the entrainment effect that acts between them. In principle the model can be expanded to have more than two interpenetrating fluids (see \\cite{P00} for instance).  As well a given superfluid can be confined to a distinct region in the star \\cite{ACL02}.  In this way the proton fluid, say, can be made to extend out farther than the superfluid neutrons.  This is a first approximation at incorporating the fact that the superfluid neutrons do not extend all the way to the surface of the star.  Our primary goal is to show that superfluidity will affect gravitational wave emission from neutron stars.  We will see that a suitably advanced, but plausible, detector will have enough sensitivity at high frequency to see modes excited during a glitch.  With such detections we will be able to place constraints, say, on the parameters that describe entrainment.  But in addition to studying glitches, we also need to analyze further the recently discovered instability in the r-modes of neutron stars \\cite{A98,FM98}. The instability is driven by gravitational wave emission (the CFS mechanism \\cite{C70,FS78,F78}) and the waves are potentially detectable by LIGO II \\cite{LOM98,OLCSVA98,AKS99a}. The conventional wisdom early on stated that mutual friction would act against the instability in a superfluid neutron star and thus effectively suppress the gravitational radiation. But Lindblom and Mendell \\cite{LM00} have found that mutual friction is largely ineffective at suppressing the r-mode instability.  However, there are many questions about the spectrum of oscillation modes allowed by a rotating superfluid neutron star and the analysis of instabilities is very likely to be much richer than the ordinary fluid case. Thus, another goal here is to lay some groundwork for a future detailed study of the CFS mechanism in superfluid neutron stars.  Specifically, we will demonstrate that the zero frequency subspace is spanned by two sets of polar (or spheroidal) and two sets of axial (or toroidal) degenerate perturbations for the general relativistic system. Like the Newtonian case \\cite{AC01b}, the polar perturbations are the g-modes which are missing from the pulsation spectrum of a non-rotating configuration, and the axial perturbations should lead to two sets of r-modes when the degeneracy of the frequencies is broken by having the background rotate.  Below we will alternate between discussions based on Newtonian gravity and those using general relativity.  Accuracy demands that a fully relativistic formalism be employed, however there are some questions of principle for which a Newtonian formalism can suffice.  For instance, in determining the number of different modes of oscillation that a superfluid neutron star can undergo it is much more tractable to use the Newtonian equations.  But, ulitmately there is the need for a general relativistic formalism. Newtonian gravity does not include gravitational waves and so one needs a fully relativistic formalism to get an accurate damping time of a mode of oscillation due to gravitational wave emission. Also, there is the well-known problem that Newtonian models do not produce reliable values for the mass and radius, in that, for a given central density, the predicted mass and radius in Newtonian models may differ considerably from those of general relativity. This is a crucial point since the superfluid phase transition in a neutron star is sensitive to density, as are the parameters relevant to entrainment, and the oscillation frequencies can depend sensitively on mass and radius. ",
        "conclusions": "\\label{con}  We have reviewed recent work to model the rotation and oscillation dynamics of Newtonian and general relativistic superfluid neutron stars.  We have seen that superfluidity affects both the background and the perturbation spectrum of neutron stars and both should therefore cause an imprint of superfluidity to be placed in the star's gravitational waves.  In particular our local analysis of the Newtonian mode equations indicates that the superfluid modes should have a sensitive dependence on entrainment parameters, something that is supported by the quasinormal mode calculations.  Given a suitably advanced detector, like the EURO configuration, we have seen that gravitational waves emitted during a Vela glitch, say, should be detectable. The key conclusion is that direct detection of gravitational waves from glitching pulsars can be used to greatly improve our understanding of the local state of matter in superfluid neutron stars.  This may become more important in the next few years because of indications of free precession in neutron stars \\cite{SLS00}, which if true means that the vortex creep model will have to be reconsidered.  We have also put in place some groundwork for a future analysis of the CFS mechanism in superfluid neutron stars.  We have done this by using the ``pull-back'' formalism to motivate fluid variations in terms of constrained Lagrangian displacements.  We have also used them as perturbations to help map out the zero frequency subspace, and found that it is spanned by two sets of polar and two sets of axial perturbations.  It remains to be seen if adding rotation will lift the degeneracy and yield two sets of polar and two sets of axial oscillation modes, or if the more general inertial hybrid modes result.  Finally, I would like to elaborate a little more on why I continue to treasure my two years in Israel with Jacob Bekenstein.  After I had completed working on the superfluid analogs of quantum field theory in curved spacetime effects I tried in vain to publish the work.  True to his character, Jacob had some very kind words of advise, which were to never worry that effort is lost when a project does not play out exactly as expected, because his own experience was that one, or many, pieces of it would eventually be of direct importance for something else.  For a young scientist, those were words of comfort and hope, and remain so for one that now has a little more experience."
    },
    "0207/astro-ph0207264_arXiv.txt": {
        "abstract": "The (\\ppp ) transition of singly--ionized carbon, [\\cii ], is the primary coolant of diffuse interstellar gas.  We describe observations of [\\cii ] emission towards nine high Galactic latitude translucent molecular clouds, made with the long wavelength spectrometer on board the Infrared Space Observatory.  To understand the role of dust grains in processing the interstellar radiation field (ISRF) and heating the gas, we compare the [\\cii ] integrated intensity with the far-infrared (FIR) integrated surface brightness for the 101 sampled lines of sight.  We find that [\\cii] is linearly correlated with FIR, and the average ratio is equal to that measured with the {\\sl COBE} satellite for all high-latitude Milky Way gas.  There is a significant decrease that was not detected with {\\sl COBE} in [\\cii ] emissivity at high values of FIR. Our sample splits naturally into two populations depending on the 60\\micron/100\\micron\\ surface brightness ratio, or color:  ``warm'' positions with $60/100 > 0.16$, and ``cold'' positions with $60/100 < 0.16$. A transition from sources with warm to those with cold 60/100 colors coincides approximately with the transition from constant to decreasing [\\cii ] emissivity.  We model the [\\cii ] and far--infrared emission under conditions of thermal equilibrium, using the simplifying assumptions that, in all regions heated by the ISRF, the most important source of gas heating is the photoelectric effect on grains and the most important source of gas cooling is [\\cii ] emission.  The model matches the data well, provided the ISRF incident flux is $\\chi_0 \\approx 1.6$ (in units of the nominal value near the Sun), and the photoelectric heating efficiency is $\\epsilon \\approx 4.3\\%$.  There are no statistically significant differences in the derived values of $\\chi_0$ and $\\epsilon$ for warm and cold sources.  The observed variations in the [\\cii ] emissivity and the 60/100 colors can be understood entirely in terms of the attenuation and softening of the ISRF by translucent clouds, not changes in dust properties. ",
        "introduction": "Diffuse interstellar gas clouds are heated by the absorption of starlight and cooled by the emission of spectral line radiation at far--infrared (FIR) and submillimeter wavelengths.  Energy can be transferred from the interstellar radiation field (ISRF) to the gas by the photoelectric effect on dust grains \\citep{wat72,dej77,dra78,bak94}, whereby far ultraviolet (FUV) (6\\,eV $< h\\nu <$ 13.6\\,eV) photons ionize grains and the ejected electrons heat the gas through collisions.  In diffuse \\hi\\ regions of low kinetic temperature ($\\tk \\sim 50-200\\K$), the ``cold neutral medium'' (CNM), this process is thought to be the dominant gas heating mechanism \\citep[see][]{wol95}.  The most abundant form of carbon gas in the CNM is the singly-ionized carbon atom, \\cii, because the first ionization potential of carbon is 11.3\\,eV and the CNM is permeated by FUV photons.  The 158\\micron\\ (\\ppp ) emission line transition of [\\cii ] should be the primary coolant of the CNM \\citep{wol95}.  There are two main reasons for this.  First, the \\cii\\ ion is relatively easy to excite by collisions given average CNM temperatures ($T\\sim 80\\K$;\\citealt{kul87}). After the hyperfine states of hydrogen, the $^2P_{3/2}$ fine structure state of \\cii\\ is the first excited state ($h\\nu/k = 91\\K$) of all gas-phase CNM constituents (the next state is the $^3P_1$ level of the \\oi\\ atom, which is 228\\K\\ above ground).  Second, once excited, the energy loss rate, given by the product $h\\nu\\,A_{ul}\\times \\ncii[^2P_{3/2}]$, is much higher for \\cii\\ than for other CNM species.  Here $A_{ul}$ is the Einstein spontaneous emission coefficient for the \\ppp\\ transition and $\\ncii[^2P_{3/2}]$ is the column density of \\cii\\ in the excited state .  Since the FUV radiation field that ionizes dust grains (and thus heats the gas) can also heat the grains, one expects a connection between the [\\cii ] 158\\micron\\ line emission from photoelectron-heated gas and the far-infrared (FIR) emission from FUV-heated dust grains.  Large-scale surveys of the Galaxy indeed reveal a link between the \\cii\\ and FIR emission properties of the CNM. An unbiased survey of spectral line and continuum emission in the Milky Way was conducted with two instruments on board the Cosmic Background Explorer ({\\sl COBE}) satellite:  the Far Infrared Absolute Spectrometer (FIRAS) and the Diffuse Infrared Background Experiment (DIRBE).  The [\\cii ] emission was observed with FIRAS to be closely correlated with \\hi\\ emission at high Galactic latitudes, with an average \\cii\\ cooling rate of $(2.65\\pm 0.15)\\times 10^{-26}\\,$erg\\,s$^{-1}\\,$(H~atom)$^{-1}$ \\citep{ben94}.  Both DIRBE and FIRAS maps of the dust continuum show that the dust opacity and FIR surface brightness are also correlated with \\hi\\ column density at high latitude \\citep{bou96}.  So on the large scale (FIRAS has an angular resolution of 7\\arcdeg, DIRBE of 40\\arcmin), CNM \\hi\\ gas has constant FIR {\\it and} [\\cii ] emissivity.  It is therefore likely that the processing of FUV radiation by dust grains (1) to produce FIR dust emission and (2) to heat the gas---resulting in [\\cii ] emission---occurs in similar ways throughout the cold interstellar medium.  It is difficult to test this idea directly using the {\\sl COBE} data, mainly because of inadequate resolution.  The emission from gas with vastly different physical conditions, for example molecular and atomic gas, might be combined together in a single FIRAS or DIRBE measurement, making it impossible to understand the physical mechanisms responsible for the observed emission.  High resolution observations of the [\\cii ] and FIR emission from translucent molecular clouds can be used to test the limits of models for gas heating by the ISRF.  Molecular gas can only exist in shielded locales where the FUV portion of the ISRF has been attenuated, and thus where the photoelectric heating is much weaker than in atomic regions.  In these regions, most of the carbon has combined into CO, and most of the [\\cii ] emission has ``turned off.''  Although FUV radiation is weak in molecular zones, optical photons are still present to heat dust grains and produce FIR emission.  Indeed, many {\\sl IRAS} cirrus clouds \\citep{low84} first detected in FIR emission at 60 and 100\\micron\\ are associated with intermediate-extinction molecular clouds, or translucent clouds \\citep{van88}.   Translucent molecular clouds are most easily observed when nearby and at high Galactic latitudes ($b\\gtrsim 15\\arcdeg$) \\citep{mag85}.  Currently over 100 nearby ($\\langle d\\rangle \\approx 105\\,$pc) translucent high--latitude clouds (HLCs), have been cataloged using their CO emission \\citep*[see][]{mag96}.  In a study of 75 HLCs, it was shown that all of the clouds are associated with \\hi, and that the HLCs probably condensed out of diffuse CNM gas \\citep*{gir94}.  We examine here the [\\cii ] and FIR emission towards a sample of translucent HLCs using the {\\sl ISO} satellite.  Since the mean angular size of HLCs in CO emission maps is $\\sim 1\\arcdeg$ \\citep{mag96}, the clouds are easily resolved by the $\\sim 71\\arcsec$ beam of {\\sl ISO}.  This allows us to choose between lines of sight towards individual clouds with varying amounts of atomic and molecular gas, which was not possible with {\\sl COBE}.  As described above, the lines of sight dominated by molecular gas include regions from which FUV photons are excluded.  In this paper we summarize the {\\sl ISO} observations of [\\cii ] emission towards HLCs, all of which are available from the {\\sl ISO} archive (\\S2).  We discover that HLCs fall into two categories based on their 60\\micron/100\\micron\\ color, ``cold'' sources and ``warm'' sources. We find that the line-of-sight [\\cii ] emissivity is consistent with the average high--latitude {\\sl COBE} cloud emission for both cold and warm positions.  We surmise that the chief distinction between the two 60/100 color regimes is the column density (\\S3).  Assuming that the local heating and cooling depend only on the fixed properties of dust grains and the attenuated ISRF intensity, we develop a model for the [\\cii ] and FIR emission of translucent clouds.  The model enables us to estimate the average line of sight photoelectric heating efficiency, $\\epsilon$, as well as the intensity of the interstellar radiation field incident on HLCs, $\\chi_0$ (\\S4).  We discuss the significance of our results, and point out some issues unresolved by this study (\\S5).  Finally, we summarize our conclusions (\\S6). ",
        "conclusions": "A linear relationship is observed between the $FIR$ and [\\cii ] emission for high Galactic latitude molecular clouds observed with {\\sl ISO} that is indistinguishable from the prediction based on high--latitude {\\sl COBE} data, implying that the [\\cii ] emissivity is constant for all low-extinction gas.  At high extinction the [\\cii ] emissivity begins to decrease due to the attenuation of the FUV portion of the interstellar radiation field. In contrast to other work, we find that differences in the 60\\micron/100\\micron\\ color can be accounted for solely by extinction variations.  Sources with both warm and cold colors seem to be exposed to the same radiation field, equal to the mean field near the Sun; and have the same photoelectric heating efficiency, close to the value for neutral grains.  The transition from sources with warm to those with cold 60/100 colors coincides approximately with the transition from constant to decreasing [\\cii ] emissivity.  Such regions where the FUV spectrum is attenuated and softened, but is still important for heating the gas, define the translucent regime."
    },
    "0207/astro-ph0207052_arXiv.txt": {
        "abstract": "The observed delay in the arrival times between high and low energy photons in gamma-ray bursts (GRBs) has been shown by Norris et al. to be correlated to the absolute luminosity of a GRB. Despite the apparent importance of this spectral lag, there has yet to be a full explanation of its origin. We put forth that the lag is directly due to the evolution of the GRB spectra. In particular, as the energy at which the GRB's $\\nu F_{\\nu}$ spectra is a maximum ($E_{pk}$) decays through the four BATSE channels, the photon flux peak in each individual channel will inevitably be offset producing what we measure as lag.  We test this hypothesis by measuring the rate of $E_{pk}$ decay ($\\Phi_{o}$) for a sample of clean single peaked bursts with measured lag.  We find a direct correlation between the decay timescale and the spectral lag, demonstrating the relationship between time delay of the low energy photons and the decay of $E_{pk}$. This implies that the luminosity of a GRB is directly related to the burst's rate of spectral evolution, which we believe begins to reveal the underlying physics behind the lag-luminosity correlation.  We discuss several possible mechanisms that could cause the observed evolution and its connection to the luminosity of the burst. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207578_arXiv.txt": {
        "abstract": "{\\small We propose a model for gamma-ray emitting microblazars based on the Compton interaction of a relativistic electron-positron plasma, ejected in a jet feature, with the UV-photon field provided by a high-mass stellar companion. Taking into account the gravitational effects of the star upon the accretion disk, we predict a jet precession which results in a variable, periodic, high-energy gamma-ray source. The specific case of Cygnus X-1 is briefly discussed.} ",
        "introduction": "The third EGRET catalog \\cite{Hartman} contains approximately 170 unidentified gamma-ray sources. A significant number of them seem to be Population I objects \\cite{Romero}. Two main groups of sources are distinguished at mid and low galactic latitudes, one related to the Gould belt region \\cite{Grenier} and the other formed by brighter sources at lower latitudes \\cite{Gehrels}. These sources are suspected to be galactic and their possible counterparts include early-type stars (both isolated and in binary systems), accreting neutron stars, radio-quiet pulsars, interacting SNRs, and microquasars \\cite{Paredes}.  We will concentrate on the last possibility, proposing a model for high-energy emission in microquasars with jets forming a small angle with the line of sight. These \\emph{microblazars} are expected to have highly variable and enhanced non-thermal flux due to Doppler boosting \\cite{Mirabel}, \\cite{Georganopoulos 1}. ",
        "conclusions": ""
    },
    "0207/astro-ph0207028_arXiv.txt": {
        "abstract": "{ Most of  the emission coming from the solar corona is confined into closed magnetic structures in the form of arcs (loops). Very little is known about the structure of stellar coronae. The magnetic topology, however, can be inferred by studying the radio emission coming from electrons trapped in the magnetic loops. Evident morphological changes are  produced in fact  by stellar rotation. We have performed 4 VLBA+Effelsberg runs distributed in time so as to cover well the rotational period of 6.44 days of the active star UX\\,Arietis. We present here some preliminary results, from those observations. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207502_arXiv.txt": {
        "abstract": "There remain several definitive $\\gamma$-ray pulsars that are as yet undetected in the optical regime. A classic case is the pulsar PSR B1951+32, associated with the complex CTB 80 SNR. Previous ground based high speed 2-d optical studies have ruled out candidates to $m_{V}$ $\\sim$ 24.  Hester (2000a) has reported an analysis of archival HST/WFPC2 observations of the CTB 80 complex which suggest a compact synchrotron nebula coincident with the pulsar's radio position. Performing a similar analysis, we have identified a possible optical counterpart within this synchrotron nebula at $m_{V}$ $\\sim$ 25.5 - 26, and another optical counterpart candidate nearby at $m_{V}$ $\\sim$24.5.  We assess the reality of these counterpart candidates in the context of existing models of pulsar emission. ",
        "introduction": "The detection of nonthermal high energy magnetospheric emission from isolated pulsars has remained a non-trivial problem, despite great advances in instrumentation and technological expertise.  To date, only 7 optical pulsars have been detected with emission believed to be magnetospherically dominated, and despite considerable effort, only 8 $\\gamma$-ray pulsars.  In contrast to radio emission, which is generally believed to be generated in close proximity to the magnetic poles, no clear theoretical model construct exists as regards the higher energies. The two principal schools of thought place $\\gamma$-ray emission localised either to the magnetic poles (Daugherty \\& Harding 1996) or located further out in the magnetosphere (Romani 1996). Considerable problems remain with these two frameworks, in terms of predicted fluxes, spectral indices and light curve morphologies, and it is clear that further work is required on this subject.  This is all the more relevant when one attempts to address the growing empirical database of lower energy emission, in particular in the optical regime. A consequence of non-linear processes within the magnetosphere, this synchrotron emission forms a useful constraint with which one can attempt to comprehensively develop a self-consistent theoretical framework. Consequently it is important to try and acquire synchrotron (optical/X-ray) data of known $\\gamma$-ray pulsars, and so extend this empirical database.  The pulsar PSR B1951+32, located within the complex combination supernova remnant (SNR) CTB 80, was first identified as a steep-spectrum, point-like source in the radio (Strom 1987), and discovery of radio pulses with an unusually fast 39.5-s period quickly followed (Kulkarni et al. 1988).  Canonically, the pulsar's age and the estimated dynamical age of the SNR are consistent at $\\sim10^5$ yrs (Koo et al. 1990) and both have been determined to be at a distance of $\\sim$ 2.5 kpc.  There is thus general agreement that the association is valid.  Evidence for pulsed emission was subsequently found in $\\gamma$-rays (Ramanamurthy et al. 1995) and possibly in X-rays (Safi-Harb et al. 1995; Chang \\& Ho 1997), with upper limits in the infrared (Clifton et al. 1988). The ROSAT observations in the X-ray regime do indicate a complex light curve strongly dominated by the intense X-ray radiation of a pulsar-powered synchrotron nebula (Safi-Harb et al. 1995; Becker \\& Truemper 1996). The unambiguous double-peaked $\\gamma$-ray (EGRET) light curve obtained by Ramanamurthy et al.  (1995) at the appropriate spin-down ephemeris suggested that the pulsar had a conversion efficiency, in terms of rotational energy to $\\gamma$-rays, of $\\sim$ 0.004. Consequently there are strong grounds for the possibility of an optical detection.  However, the pulsar is quite distant and located within a rather complex SNR. Recent radio observations at 92cm (Strom \\& Stappers 2000) indicate that the pulsar is located at the edge of the flat radio spectrum 'core' of the SNR, and it is clear from X-ray observations that the pulsar is to some extent interacting with its environment.  From models such as Pacini \\& Salvati (1987) and Shearer \\& Golden (2001) we would expect emission in the $V$ magnitude range 24-26 depending upon the line of sight absorption, estimated from E(B-V) = 0.8-1.4 (Blair et al. 1984). Ground-based CCD observations by Blair \\& Schild (1985) and Fesen \\& Gull (1985) yielded a relatively crowded field with two possible counterparts at $m_v ~\\sim$ 20 and $m_v ~\\sim$ 21 respectively (known hereafter as counterparts 1 and 2).  Using a ground-based MAMA detector in the TRIFFID camera, we have previously examined the central field of CTB 80, but could find no evidence of pulsations in either $B$ or $V$ from these two counterparts (O'Sullivan et al., 1998).  It was noted that counterpart 1 had an extension (see Fig.~\\ref{mama_V_ext}) towards the mapped radio timing position given in Table~\\ref{psr_radio}, which implied an unresolved stellar combination and/or plerionic/remnant material. Using PSF-fitting and deconvolution techniques, the removal of counterpart 1 yielded a best-estimate imposed decomposition of the ``extension'' into two further point sources (numbered 4 and 5 in Fig.~\\ref{mama_V_ext}), while a further point source (numbered 6 in Fig.~\\ref{mama_V_ext}) was detected some distance further away.  These 3 sources had apparent magnitudes of $m_V$ $\\approx$ 22.1/$B-V$ $\\approx$ 0.8, $m_V$ $\\approx$ 22.6/$B-V$ $\\approx$ 0.6, and $m_V$ $\\approx$ 23.1/$B-V$ $\\approx$ 0.0. However, none of them exhibited pulsed emission at the 1\\% level.  As was noted, both the photometry and time series analysis were complicated by their low signal-to-noise and proximity to counterpart 1.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{MS1912f1.eps}}  \\caption{Contour plot of the TRIFFID/MAMA V-band 8173-second summed image of the central region of CTB 80. This is a re-presentation of the data first discussed by O'Sullivan et al. (1998); the image has been rotated and registered to the same coordinate system as the WFPC2 F547M image discussed in this paper. The main contours range from intensities of 200 to 252 in steps of 3 and were chosen to highlight the fainter structures. The pulsar optical counterpart candidates from Blair \\& Schild (1985), Fesen \\& Gull (1985), and O'Sullivan et al. (2000) are marked by their ID numbers (1,2,4,5,6). Other nearby detected stars are also marked (w, y). The bright stars identified as B and C by Blair \\& Schild (1985) are also marked and highlighted by a few extra contours between intensities of 300 and 600. }  \\label{mama_V_ext} \\end{figure}  It is clear that in order to unambiguously resolve the various apparent components within the radio error ellipse, diffraction-limited photometry must be obtained, which would facilitate further attempts to detect optical pulsations from the suspected optical counterpart(s) to PSR B1951+32. Consequently we obtained from the HST archive WFPC2 images of the CTB 80 SNR obtained with the F547W, F673N, F656N, and F502N passbands.  In the next section we detail the reduction of the various exposures and their astrometric and photometric analyses. From this we assess the feasibility that we have identified new plausible optical counterparts, and assess the implications - both in terms of follow-up high speed 2-d photometry, and for current theoretical models. ",
        "conclusions": "Hester (2000a) has reported on a similar analysis of the same archival HST data, and agrees that the knot of extended continuum emission is synchrotron dominated as a consequence of the pulsar wind.  Hester (2000b) places the radio counterpart 0.5'' to the East/SE of this nebula (see Fig.~\\ref{hst_hester}) - roughly between our proposed counterpart candidates 1$_{HST}$ and 4$_{HST}$, although closer to 1$_{HST}$ than to 4$_{HST}$. We have reproduced his results by mapping the Foster et al. (1994) radio timing position onto the {\\it original} GSC-based astrometric solution from the HST data pipeline.  However, we are more confident in our 2MASS-recalibrated and astrometry, which also benefits from our correction utilising the accurate proper motion rate, published some time after Hester's analysis by Migliazzo et al. (2002). As a result our mapped radio timing position differs from Hester's by 0$\\farcs$47. Hester also comments that the relative locations, orientation and luminosity are similar to the knot situated 0.5'' to the SE of the Crab pulsar. The agreement, in terms of synchrotron luminosity from the nebula, is compelling, regardless of which of the two astrometric solutions are used, since they overlap within their errors. However the location of the pulsar is still uncertain, given the considerable discrepancy between the two best independent radio positions.  We have outlined the evolution in these radio positions and optical astrometric solutions which, taken together, locate the pulsar in an error region which is still too large to definitively pin down the pulsar's optical counterpart; however the more accurate mapped radio position (in terms of formal quoted errors) lies to the East/SE of the nebula. This supports the geometric orientation suggested by Hester (2000a,b).   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{MS1912f5.eps}} \\caption{Location of the radio counterpart according to Hester (2000b), along an axis that bisects apparent Balmer-dominated lobes (which lie outside the field shown here) in relation to the synchrotron nebula. Figure adopted from Hester (2000b).} \\label{hst_hester} \\end{figure}  One of the two newly resolved sources which we have identified as 1$_{HST}$ and 4$_{HST}$ may therefore indeed be the $\\sim$ 24th-26th magnitude pulsar as predicted by the models of Pacini \\& Salvati (1987) and Shearer \\& Golden (2001). The other possibilities are that both are unrelated field stars, or that 1$_{HST}$ is a localised `knot' within the bigger `knot' of the synchrotron nebula while 4$_{HST}$ is a field star; both of these possibilities would imply that the true optical counterpart is fainter still. This would possibly reconcile one difficulty with the present analysis - the inconsistency of the two new candidates with the formally small error ellipse of the radio timing position; but it would create another difficulty - the inconsistency of candidates fainter than 26th magnitude with the phenomenological model predictions. Deeper HST and/or diffraction-limited adaptive optics imaging will be needed to determine which of these various possibilities is the correct interpretation, by putting some hard constraints on the colour and spectral index of both the synchrotron nebulosity and the candidate counterparts. Goldoni et al. (1995) noted that the optical spectral index {\\it steepened} with age for the known optically emitting pulsars, while Shearer \\& Golden (2001) showed that this is consistent with the flattening of the pulse-peak luminosity relationship with the outer field strength. Consequently, with its expected steep optical spectral index, one would expect that the optical counterpart to PSR B1951+32 will not be as distinctly blue as younger Crab-like pulsars, but nevertheless distinguishable from regular thermally-emitting stars. Additionally, one could test the candidates by looking for proper motion, which (since the epoch of these HST observations) should now amount to a possibly detectable $0^{\\prime\\prime}.12$. Follow-up timing studies with the upgraded TRIFFID imaging photon-counting camera would permit the search for optical pulsations from a securely identified counterpart, with much higher statistical significance than previously. Not only would this provide the essential confirmation of any optical counterpart, it would also test the hitherto successful optical models of Pacini \\& Salvati (1987) and Shearer \\& Golden (2001) in an age regime where few optical pulsars have been found."
    },
    "0207/astro-ph0207358_arXiv.txt": {
        "abstract": "{ The globally-averaged star formation rate in the Universe has been steadily declining since at least $z\\sim 1$.  This may be due either to very local processes operating within the average galaxy, or to external, environmental effects.  Specifically, the build-up of structure may be responsible for terminating star formation in some galaxies and thus decreasing the global average.  We summarize our previous and ongoing work to  distinguish between these possibilities, by determining the average star formation rate as a function of redshift and environment, out to $z=0.5$. } \\addkeyword{galaxies: clusters}  \\begin{document} ",
        "introduction": "There is good observational evidence that the total amount of star formation in the universe has declined substantially over the past few Gigayears \\cite{L96,Wilson+02}. This may be a reflection of local physics on galactic scales, whereby galaxies consume their gas supply as time progresses, and star formation gradually declines.  However, observations show that star formation is inhibited in dense environments \\cite{B+97,B+98}.  In hierarchical models of galaxy formation, the abundance of dense clusters increases with time; therefore, perhaps the growth of structure is partly driving the decline in global star formation.  However, this scenario is only viable if a suppression of star formation is observed in environments less extreme than rich clusters, since the latter are too rare to have a significant impact on the globally averaged star formation rate (SFR).  We can attempt to distinguish between these two interpretations by tracing the SFR as a function of environment at a series of redshifts.  If the SFR--local density correlation is independent of redshift, there will be evidence that the global decline is due to environmental effects, coupled with the hierarchical growth of structure. To address this, we have begun a large programme to measure SFRs in different environments out to $z\\sim 0.5$.  The focus is on relatively low-density environments, since these have not been studied in much detail, and are common enough to contribute significantly to the global average. ",
        "conclusions": "Our goal is to construct the star formation history of the universe as a function of galaxy environment. Figure~\\ref{fig-envt} shows how the mean EW([\\ion{O}{2}]) within the virial radius depends on environment and redshift, for our present sample of groups and clusters. Surprisingly, the amount of emission in clusters is approximately constant with redshift, so the difference between the cluster and field SFR {\\it increases} with redshift.  This suggests that, at least at cluster densities, the average SFR is determined by local environment, and not an internal galaxy clock.  On the other hand, galaxy groups at $z\\sim 0.4$ have SFRs only slightly lower than the expected global average at that redshift. If SFR is entirely environment-dependent, this means that the local analogues of these groups should have substantially higher SFRs than the average."
    },
    "0207/astro-ph0207672_arXiv.txt": {
        "abstract": "Studies of brown dwarfs as a distinct galactic population have been largely pioneered by dedicated, vast-area surveys at deep optical and near-infrared wavebands. The Two Micron All Sky Survey has, because of its full-sky coverage and depth, been the most successful at providing new targets for further study. The author briefly reviews some of the ways in which 2MASS data have been used by brown dwarf researchers. One of the by-products -- namely the 2MASS Team's follow-up spectroscopy of brown dwarf candidates -- is potentially valuable as a scientific resource itself and is also being released to the public. ",
        "introduction": "The Two Micron All Sky Survey (2MASS) has, along with the Sloan Digital Sky Survey (SDSS) and to a lesser extent the Deep Near-infrared Survey (DENIS), provided a treasure trove of information with which to study brown dwarfs both in the field and in nearby open clusters. As an illustration, Table 1 contains a breakdown of the number of known nearby L and T dwarfs so far discovered by each. In addition to providing a discovery tool for field brown dwarfs, 2MASS data in particular have also been used in a variety of other ways to study substellar objects. The 2MASS mission is briefly overviewed in \\S2 and some of its contributions to brown dwarf science are highlighted in \\S3.  \\begin{table} \\begin{center} \\caption{Numbers of L and T Dwarfs Currently Known} \\begin{tabular}{ccc} \\tableline & No.\\ of     & No.\\ of\\\\ Survey   & L Dwarf     & T Dwarf\\\\ & Discoveries & Discoveries\\\\ \\tableline 2MASS    & 155 & 18\\\\ SDSS     &  72 & 11\\\\ DENIS    &   9 &  0\\\\ Others   &   9 &  3\\\\ \\tableline \\tableline \\end{tabular} \\end{center} \\end{table}  Furthermore, Keck Observatory follow-up of 2MASS-selected brown dwarf candidates has resulted in a wealth of high-quality spectroscopic information which the author is releasing to the general public for the first time. In tandem with the Keck data, the author is also releasing libraries of other dwarf spectra that he and his collaborators have amassed over the period 1989-2002. These libraries are being paired with the 2MASS final release dataset to produce tables of uniform spectral types and near-infrared photometry for all dwarf spectral types from K5 through T8. These libraries are further described in \\S4. ",
        "conclusions": ""
    },
    "0207/astro-ph0207391_arXiv.txt": {
        "abstract": "{We present a $\\rm 4.5-4.85~\\mu m$ $R=5\\,000$ spectrum of the low mass class I young stellar object GSS 30 IRS1 ($L=\\rm 25~L_{\\odot}$) in the $\\rho$ Ophiuchus core, observed with the infrared spectrometer (ISAAC) on the {\\it Very Large Telescope} (VLT-UT1). Strong line emission from the ro-vibrational transitions of $\\rm ^{12}CO$ and $\\rm ^{13}CO$ is detected. In total more than 40 distinct lines are seen in the covered region. The line emission is spatially extended and detected up to $\\rm 2\\arcsec = 320~AU$ from the central source but is spectrally unresolved ($\\Delta v < 30~\\rm km~s^{-1}$). This is the first time strong emission in the fundamental ro-vibrational band from CO has been observed from an embedded young stellar object. The line fluxes were modeled using a 1-dimensional full radiative transfer code, which shows that the emission is fully consistent with a gas in LTE at a single well constrained temperature ($T=515\\pm5~\\rm K$). Furthermore, the ratios between lines from the two detected isotopic species of CO show that the $\\rm ^{12}CO$ lines must be optically thick. However, this is inconsistent with the observed spatial extent of the emission, since this implies such low CO column densities that the lines are optically thin. A likely solution to the discrepancy is that the lines are emitted by a smaller more dense region and then scattered in the bipolar cavity present around the central star. This gives a rough estimate of the total molecular gas mass of $\\rm 1-100~M_{\\oplus}$ and a physical extent of $\\rm\\sim20-100~AU$. We propose that the most likely origin of the line emission is post-shocked gas in a dense dissociative accretion shock from the inner $\\rm 10-50~AU$ of a circumstellar disk. The presence of a shock capable of dissociating molecules in the disk will have implications for the chemical evolution of disks around young low mass stars. ",
        "introduction": "The innermost $\\rm 50~AU$ of the circumstellar environments around embedded young stellar objects is poorly constrained observationally due to the large extinction through the embedding material ($A_V>20~\\rm mag$) and the small angular size ($<1\\arcsec$) of the region. The usual molecular probes in the millimeter-submillimeter region are not effective for the temperatures ($T>200~\\rm K$) and densities ($n\\rm_{H_2}>10^7~cm^{-3}$) thought to be present. However, an understanding of the processes taking place in this regime is essential to obtain a complete picture of the process of accretion and the driving of outflows from low mass protostars as well as the early chemical and physical evolution of circumstellar disks.  One of the most effective probes of warm dense gas is through emission in molecular ro-vibrational transitions. The most common bands readily available from ground-based facilities are the CO overtones and $\\rm H_2$ fundamental bands around $\\rm 2.2~\\mu m$ \\citep[e.g.][]{Reipurth} and the fundamental transitions of CO in the $M$-band around $\\rm 4.7~\\mu m$. Also emission from very hot water gas ($\\sim 2000~\\rm K$) near $2.29~\\rm \\mu m$ has been reported toward a few sources \\citep[e.g.][]{Najita2}. Since the upper levels of these transitions lie at temperatures of up to a few thousand Kelvin, they probe hot gas with temperatures between 100 and $\\rm 1000~K$, making them ideal to study the region of interaction between disk, protostar, accretion and outflow. With the new generation of sensitive ground-based high resolution spectrometers for the $\\rm 3-5~\\mu m$ region an efficient window has been opened for the detailed study of the CO ro-vibrational lines toward embedded young stellar objects (YSO). The past generation of instruments was suitable to either observe a low resolution spectrum with a fairly wide spectral range \\citep[e.g.][]{teixeira} or a high resolution echelle spectrum with a very narrow spectral range \\citep[e.g.][]{carr}. The main exception are the Fourier Transform Spectroscopy (FTS) observations of \\cite{mitchell,mitchell4}, who observed a number of high-mass stars in the entire M band at high spectral resolution ($R>10^6$), but such studies are limited to the brightest objects. VLT-ISAAC has a large instantaneous spectral range in the $M$-band ($0.237~\\mu m$), which combined with a medium resolution of $R=5\\,000 - 10\\,000$ and a limiting magnitude of $M\\sim 10$ allows the entire fundamental band of gaseous CO of a wide range of young stars to be observed in a short time, including low mass stars down to a few tenths of a solar mass in the nearest star-forming clouds.  The embedded stars studied so far in CO ro-vibrational bands have showed mostly lines in absorption \\citep{mitchell4,adwinElias29,adwinL1489} implying that the warm CO gas is seen in front of a bright infrared continuum, produced by hot dust close to the central object. The fundamental CO lines are usually only seen in emission towards more evolved sources characterized by a class II type spectrum where a circumstellar disk is directly visible \\citep{carr,Blake}. CO overtone bandhead emission at $\\rm2~\\mu m$ has been observed in emission toward a few intermediate mass pre-main sequence stars \\citep{Thompson, Najita} and T Tauri stars \\citep{carr1989}, and has been associated with hot gas ($T\\sim 1\\,500-5\\,000~\\rm K$) located within a fraction of an AU in a Keplerian disk.  We present here the peculiar $\\rm 4.5-4.8~\\mu m$ spectrum of the illuminating source IRS1=\\object{Elias 21} of the bipolar reflection nebula \\object{GSS 30} located in the core of the $\\rm \\rho$ Ophiuchus molecular cloud at a distance of $\\rm 160~pc$. It has been classified as a low mass class I YSO from its spectral energy distribution (SED) \\citep{GSS, Elias, WLY} and low bolometric luminosity \\citep[$\\rm L_{bol}=21-26~L_{\\sun}$, ][]{GWAY,Bontemps}. Extensive polarimetric studies in the H and K band of the reflection nebula have shown that the source is surrounded by a large ($\\rm\\sim2\\,000~AU$) disk-like envelope and a smaller circumstellar disk of $\\rm\\sim150~AU$, which are inclined about $25\\degr$ away from the plane of the sky, i.e. close to edge-on \\citep{Chrysostomou,Chrysostomou2}. The high degree of linear polarization (up to 50\\%) as well as the presence of circular polarization imply that the light coming from the reflection nebula must have been multiply scattered, before heading into the line of sight.  Two other sources (IRS2 and IRS3) are present toward the $K$-band reflection nebula (within 20\\arcsec  of IRS1). IRS2 has a class III SED and is probably a more evolved star \\citep{Andre}. IRS3 has a class I SED, but is much fainter than IRS1 in the near-infrared. It has a bolometric luminosity of $0.13~L_{\\odot}$, \\citep{Bontemps}. IRS3 shows strong $\\rm 6~cm$ emission \\citep[see][where IRS3 is designated \\object{LFAM 1}]{Leous}.  The presence of a molecular outflow has not been firmly established. High velocity red- and blue-shifted CO millimeter emission to the south of GSS 30 has previously been reported by \\citet{Tamura}, but since both lobes are located to the south of the infrared source, the gas is likely to be associated with the \\object{VLA 1623} jet, which is passing only $30\\arcsec$ to the SW of IRS1. Using millimeter interferometric line data, \\citet{Zhang} find evidence for a spherical expansion of the core surrounding the three sources in GSS 30. In addition, the inner region of the core seems to be cleared, which is indicative of a young outflow. This is additionally supported by the presence of variable unbound water maser emission from within $0\\farcs3$ of IRS1 \\citep{Claussen}. ",
        "conclusions": "In summary, to determine the origin of the CO gas emission it must be explained why only a single well-defined temperature is needed. Which mechanism can heat up to $\\rm 100~M_{\\oplus}$ of gas to a unique temperature of $\\rm 500~K$, yet keep the intrinsic line width less than $\\rm 30 ~km~s^{-1}$? Finally, why is no other similar embedded source from the literature showing the same strong CO rovibrational emission as GSS 30 IRS1?   \\subsection{Outflow?}  If the warm CO gas resides in the wind/outflow component of the circumstellar environment then an estimate of the time needed to create it can be found from the mass loss rate derived in Section \\ref{massloss} under the assumption that the wind is predominantly ionized. A mass between 3 and 100$~M_{\\oplus}$  corresponds to a production time of minimally 5 years and maximally a few centuries, while the most likely time is about one century. It seems unlikely that hot, ionized material emitted over a longer period should thermalize at a single temperature. Also, an outflow origin of the line emission would produce broad wings in the lines, which is not observed. If the wind is predominantly neutral, the resulting mass loss rate would be so high that a clear outflow should have been detected.  \\subsection{Inner disk?}  Another possibility is that the emission is produced by warm thermalized gas in the disk itself. Since more evolved T Tauri disks are known to exhibit similar CO emission, although with smaller intensities, it is conceivable that we are seeing an equivalent process in the case of a younger disk. A typical circumstellar disk around a low mass young star is often observed to have a mass of a few times the minimum solar nebula, i.e. $\\rm M_{disk} \\sim a~few~ 0.01~M_{\\odot}$ \\citep{Andre,Osterloh}. The disk mass for GSS 30 IRS1 inferred by the 1.3 mm continuum emission is $M_{\\rm disk} = 0.03~M_{\\odot}$ \\citep{Andre}. In hydrostatic equilibrium, a disk irradiated by the central star has a surface density $\\Sigma\\sim R^{-3/2}$ \\citep{CG}. If viscous dissipation in the disk is taken into account the surface density attains a flatter $R$-dependency, $\\Sigma\\sim R^{-1}$. In both cases the accumulated disk mass reaches the observed $\\rm 10-100~M_{\\oplus}$ between 2 and 10~AU from the central star. A single temperature of $\\rm 515~K$ is however not consistent with the disk models, which prescribe a large range in temperatures with a small 1000-2000~K component within a few stellar radii to an extended 100-200 K component within a few AU of the central star. However, in no models for passive circumstellar disks is it expected that the disk exhibits a single temperature, indeed quite the contrary is the case. Furthermore, Keplerian rotation within a few AU would produce observable line broadening. The narrow lines could only be explained if we are seeing light emitted vertically from the plane of the disk before being scattered into the line of sight, effectively removing the effects of rotation.  \\subsection{Accretion shock?}  A final option is that the emission lines are cooling lines from post-shocked dense gas. Vibrational $\\rm H_2$ emission is one of the principal tracers of low density ($n\\rm_{H_2}\\lesssim 10^7~cm^{-3}$) shocked gas. As mentioned in Section \\ref{HotGas} the (0--0) $\\rm S(9)$ line of $\\rm H_2$ at $\\rm 4.695~\\mu m$ is seen in our $M$-band spectrum, while no $\\rm H_2$ lines are seen in the $K$-band spectrum of GSS 30 IRS1 by \\cite{GL} although no upper limit is given. Assuming that the observed $\\rm H_2$ (0--0) $\\rm S(9)$ line is thermalized at the CO temperature, the required molecular gas mass is $5~ \\rm M_{\\oplus}$. This is assuming that the $\\rm H_2$ emission is seen directly and is not corrected for extinction, since the extinction is hard to estimate for embedded stars. A typical Ophiuchus extinction of $A_{\\rm V}=25~\\rm mag$ \\citep{teixeira2} will increase the molecular gas mass by a factor of 2. If the emission in the $\\rm H_2$ line is scattered in the same way as the CO emission all values must be multiplied by $\\epsilon$. If the $\\rm H_2$ emission is thermalized at $\\rm 515~K$ and directly observed then the integrated line flux expected for the $\\rm 2.12~\\mu m$ (1--0) S(1) line is $\\rm 7.5\\times 10^{-13}~erg~cm^{-2}~s^{-1}$, which should be observable. If the $\\rm H_2$ lines are scattered into the line of sight with an efficiency of $\\epsilon$, then no $\\rm H_2$ lines should be visible in the $M$-band. Consequently, the $\\rm H_2$ line does not seem to be directly associated with the same gas emitting the CO lines since it is at least a factor of 10 too bright, but it may be an indication of a warmer component or the line may be pumped. Sensitive observations of $\\rm H_2$ lines in the K and L band are needed to unambigously determine the origin of the molecular hydrogen emission.  \\cite{Neufeld} show that if a dense gas of $10^{7.5}<n\\rm_{H_2}<10^{12}~cm^{-3}$ is shocked, the vibrational emission from $\\rm H_2$ is significantly suppressed partly because the shock dissociates the $\\rm H_2$ molecules and partly because the reforming molecules can easily be collisionally de-excited in the dense post-shock gas. The main coolants will be rotational and vibrational lines from $\\rm H_2O$, $\\rm OH$ and $\\rm CO$. This is valid for a wide range of shock velocities $5<v_{\\rm s}<100~\\rm km~s^{-1}$ as long as the shock is of jump-type. A J-shock occurs in any case for $v_{\\rm s}\\gtrsim \\rm 30~km~s^{-1}$ and for smaller shock velocities if the magnetic field is weak or if the length scale over which the shock occurs is small compared to the length scale for acceleration of charged particles. Furthermore, the theoretical post-shock temperature structure gives an interesting prediction. A J-shock will dissociate the gas at the shock front. The gas then cools slowly  from $\\rm 10^5~K$ principally through the $\\rm Ly\\alpha$-line until $\\rm CO$ forms at around $\\rm 7000~K$ and $N\\rm(HI)\\sim 10^{21}~cm^{-2}$. Once the CO molecules are available the gas cools quickly to about $\\rm 500~K$ through the $\\rm CO$ vibrational lines, where a constant temperature plateau is maintained due to the release of chemical potential energy as $\\rm H_2$ is formed which is balanced against the thermalized vibrational emission from molecules. The plateau occurs at an atomic hydrogen column density of $\\rm 10^{21}-10^{22}~cm^{-2}$, which in the model by \\cite{Neufeld} corresponds to a $\\rm CO$ column density of $\\rm 3\\times 10^{17}-3\\times 10^{18}~cm^{-2}$. It is stressed that the slower C-shock scenario was not treated in such high density models, and it may produce similar cooling lines. Further modelling is needed to exclude a slow shock. It is evident that the model values are remarkably similar to the values observed in GSS 30 IRS1 with temperature $T_{\\rm gas}\\rm =515~K$ and column density $N\\rm (CO)=2\\times 10^{18}~cm^{-2}$. However, it is unclear why no hot $\\rm 5000-7000~K$ component in the fundamental CO band is observed, which is also expected from the shock. Also no overtone emission is seen in the $K$-band spectrum of \\cite{GL}. One plausible explanation is that the accretion activity is highly variable and that the hot component is only seen as the shock is occurring while the warm thermalized gas is seen in the post-shock medium. This would be supported by the variable water maser emission. The single observed $\\rm H_2$ line is not inconsistent with an accretion shock although it is too bright to be associated with a dense shocked gas and may only be an indication of a warmer less dense component.  As described in \\cite{Neufeld}, the most likely site for a dense shock in an embedded YSO is in a circumstellar disk which is accreting matter. When infalling matter is colliding with the disk at supersonic velocities, a J-shock is expected to occur at the surface of the disk. If the accretion velocities are too low to produce a J-shock, the molecules will not be dissociated and the observed temperature plateau will not be present. The main uncertainty is if the high infall velocities required to create a dissociative shock are possible. However, infall velocities of $\\rm 10~km~s^{-1}$ within $\\rm 10~AU$ are reasonable \\citep{Cassen}. Therefore, if the dissociative accretion shock scenario holds, the intrinsic line widths are {\\it predicted} to be larger than or very close to $\\rm 10~km~s^{-1}$. An accretion shock at a large distance ($\\rm \\gtrsim 10~AU$) from the central star is consistent with the observation that the line emitting region is physically separated from the continuum emitting region. A shock will not heat a large amount of dust and the continuum emission, which has a colour temperature of about $\\rm 700~K$, is likely to be produced by irradiated dust within $\\rm 1~AU$ of the star.    Other observable tracers of a shock in a dense medium are the rovibrational lines from the $\\rm H_2O~\\nu_2$ bending mode around $\\rm 6~\\mu m$ as well as the vibrational transitions of $\\rm OH$, all of which are unavailable from ground-based instruments. However, they may be detected by the upcoming Stratospheric Observatory For Infrared Astronomy (SOFIA) mission.  The possible dissociation of the molecules has important consequences for the chemistry in the inner disk. If the emission is indeed coming from an accretion shock, and if a significant amount of matter passes through the shock, it will significantly affect the models of chemical evolution of the innermost parts of disks around low mass YSOs, by evaporating ice mantles and inducing high temperature chemistry."
    },
    "0207/astro-ph0207444_arXiv.txt": {
        "abstract": "Low spectral resolution ($\\Delta\\lambda$/$\\lambda\\sim$50) mid-infrared observations with high angular resolution (0.3--0.5$''$) using the Long Wavelength Spectrometer on the Keck I Telescope are reported of the nuclei of five highly luminous infrared bright galaxies. Spectra of eight distinct nuclei, ranging in luminosity from $\\sim 10^{11}$L$_ {\\odot}$ to $>10^{12} $L$_{\\odot}$ have been obtained.  Four of the nuclei show the characteristic PAH emission features, i.e. 11.3$\\mu$m emission as well as the 8.6$\\mu$m shoulder of the 7.7$\\mu$m band.  The other nuclei show either weak PAH emission bands or no evidence for these bands.  The high spatial resolution of the observations reveals extended emission in the 11.3$\\mu$m PAH band associated with several of the compact nuclear sources.  When proper account is taken of the diffuse PAH emission, most of the compact sources show little or no directly associated PAH emission. The diffuse PAH emission is extended over spatial scales of 100--500 pc; its presence shows that there is significant circumnuclear UV/optical emission exciting the aromatic bands, most likely associated with circumnuclear starbursts.  After the spectra of the nuclear sources are corrected for the spectrum of the diffuse PAH emission, the peak apparent silicate optical depth at 9.7$\\mu$m can be as great as 15, corresponding to $>$ 150 magnitudes of visible light extinction.  Because of the large silicate optical depths, mid-infrared spectra are not probing  the nature of the true nuclei in the most opaque compact sources. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0207/astro-ph0207499_arXiv.txt": {
        "abstract": "This is the first in a series of three papers presenting a new calculation of the mass of the Galaxy based on radial velocities and distances measured for a sample of some 100 faint $16<B<20$ field blue horizontal-branch (BHB) stars. This study aims to reduce the uncertainty in the measured mass of the Galaxy by increasing the number of halo objects at Galactocentric distances $r>30$ kpc with measured radial velocities by a factor five. Faint A-type stars in the Galactic halo have been identified from $UB_JR$ photometry in six UK Schmidt fields. These samples include field BHB stars as well as less luminous stars of main-sequence surface gravity, which are predominantly field blue stragglers. We obtain accurate CCD photometry and spectra to classify these stars. This paper describes our methods for separating out clean samples of BHB stars in a way that is efficient in terms of telescope time required. We use the high signal--to--noise ratio (S/N) spectra of A--type stars of Kinman, Suntzeff \\& Kraft (published in 1994), and their definitive spectrophotometric $\\Lambda$ classifications, to assess the reliability of two methods, and to quantify the S/N requirements. First we revisit, refine and extend the hydrogen line width {\\em versus} colour relation as a classifier (here called the {\\em $D_{0.15}$--Colour} method). The second method is new and compares the {\\em shapes} of the Balmer lines.  With this method (here called the {\\em Scale width--Shape} method) there is no need for colours or spectrophotometry. Using the equivalent width of the Ca II K line as an additional filter we find we can reproduce Kinman, Suntzeff \\& Kraft's $\\Lambda$ classifications with both methods.  In a sample of stars with strong Balmer lines, EW H$\\gamma>13$\\AA\\, (equivalent to the colour range $0 \\leq (B-V)_0 \\leq 0.2$), halo BHB stars can be separated from halo blue stragglers reliably. For the spectroscopy (i.e. both classification methods) the minimum required continuum S/N is $15\\,{\\mathrm\\AA}^{-1}$. For the {\\em $D_{0.15}$--Colour} method $(B-V)_0$ colours accurate to 0.03 mag. are needed. ",
        "introduction": "The mass of the Galaxy is a key quantity for our understanding of the nature and distribution of dark matter. Although we know a great deal about the contents and properties of the halo of the Galaxy our current estimate of the total mass is subject to a large uncertainty. In part this is due to the Sun's location within the disk of the Galaxy which means it is difficult to determine the rotation curve accurately outside the solar radius. Beyond $\\sim 20\\,$kpc measures of the enclosed mass of the halo have relied principally on the kinematics of satellite galaxies and globular clusters. The current state of the art is the analysis by Wilkinson \\& Evans (1999, hereafter, WE99) who calculate the mass within $50\\,$kpc to be $5.4^{+0.2}_{-3.6} \\times 10^{11} M_{\\odot}$, and the total mass to be $1.9^{+3.6}_{-1.7} \\times 10^{12} M_{\\odot}$. The quoted uncertainties are larger than in earlier studies (e.g. Little \\& Tremaine, 1987, Zaritsky et al. 1989, Kochanek, 1996) and were determined by Monte Carlo simulation of artificial data sets. The two principal sources of error are i) the large uncertainties in the proper motions for the six satellites for which measures exist, and ii) the small size of the dataset \\---\\ there are currently only 27 known satellites at Galactocentric distances $> 20\\,$kpc.  Therefore, we are in the peculiar situation where the mass profile of the Galaxy is less well determined than for some nearby spiral galaxies. There is no possibility of substantially increasing the number of known satellite galaxies and globular clusters and the best prospect for improving the measurement of the mass is through isolating large numbers of another distant halo--tracer. Field blue horizontal branch (BHB) stars should provide just such a tracer and WE99 calculate that to reduce the uncertainty on the total mass to $20\\%$ requires a sample of 200 distant BHB stars.  This is the first in a series of three papers presenting a new calculation of the mass of the Galaxy using radial velocities of BHB stars. Field BHB stars are luminous standard candles that are abundant in the Galactic halo (e.g. Yanny et al. 2000), and for nearly twenty years, since the study of Pier (1983), have presented an under--exploited resource with which to measure the density profile and phase space structure of the Galaxy halo out to large distances, $\\sim 100\\,$kpc. A number of dynamical analyses of rather small samples of BHB stars have been published (e.g. Sommer--Larsen, Christensen \\& Carter, 1989, Norris \\& Hawkins, 1991, Arnold \\& Gilmore, 1992). Unfortunately, samples of field A--type stars in the halo include not only BHB stars but also stars of main sequence surface gravity, field blue stragglers, that are some 2 magnitudes less luminous. Progress towards the goal of acquiring a large sample of distant BHB stars has been slow because of the difficulty of separating out the BHB stars without the investment of large amounts of telescope time. In this first paper we describe our procedures for classifying samples of halo A--type stars, and present a new efficient method that requires only spectroscopic observations of intermediate signal--to--noise ratio. In Paper II we will present photometry and spectroscopy of faint $16<B<20$ candidate BHB stars in two northern high Galactic latitude fields and four southern fields. Paper III will contain the dynamical analysis of the new sample of confirmed BHB stars and a new estimate of the mass of the Galaxy.  The structure of this paper is as follows. In \\S2 we provide a brief outline of the basic parameters of our new survey for distant field BHB stars in the halo. \\S3 contains a summary of previous methods employed to sieve out BHB stars from samples of halo A--type stars. In \\S4 we present the details of the two classification methods we employ, called the {\\em $D_{0.15}$--Colour} method, and the {\\em Scale width--Shape} method. We describe the procedures used to measure the Balmer line profiles and the equivalent width of the Ca II K line, and quantify the random and systematic errors. In \\S5 we use the high signal--to--noise ratio (S/N) spectra of Kinman, Suntzeff \\& Kraft (1994; hereafter KSK) and their spectrophotometric classifications to quantify the S/N requirements for applying these methods to samples of faint halo stars. In \\S6 we discuss the use of the Ca II K line as a metallicity indicator. Finally, \\S7 provides a summary of the main conclusions of the paper. ",
        "conclusions": "\\begin{figure} \\rotatebox{0}{ \\centering{ \\scalebox{0.40}{ \\includegraphics*[20,140][700,700]{fig9.ps} }}} \\caption{Ca II K line ($3933\\,$\\AA) EW for selected stars in KSK, listed in Table 4. The curves represent lines of metallicity for [Fe/H] = -1.0, -2.0 and -3.0 taken from Wilhelm et al. (1999a). The straight line represents a best fit to stars in the Pleiades and Coma clusters (open triangles) assumed to be of solar metallicity. Filled squares are stars from M92 ([Fe/H] = -2.2) and filled triangles are from M3 ([Fe/H] = -1.5). Field BHB stars are shown by open circles labeled with the KSK adopted values.} \\label{cak_plot} \\end{figure}  The purpose of this paper has been to set out efficient methods for classifying reliably halo A-type stars into the classes BHB and blue straggler, and to quantify the S/N requirements. We disregarded the spectrophotometric $\\Lambda$ method and Stromgren $u$ photometry as possible methods of classification because they require too much telescope time, bearing in mind that additional spectroscopy is required to measure radial velocities. Instead we have concentrated on measurements of the Balmer lines H$\\gamma$ and H$\\delta$ from intermediate-resolution ($\\sim 3$\\AA) spectroscopy (not spectrophotometry). The {\\em $D_{0.15}$--Colour} method plots the width of the Balmer lines against $(B-V)_0$ colour. The main improvements we have introduced over previous applications of this method are profile fitting of the Sersic function, and quantification of the errors. We find that with spectra of S/N$=15\\,$\\AA$^{-1}$ and colours accurate to $3\\%$ samples of halo BHB stars will be $\\sim87\\%$ complete with only $\\sim7\\%$ contamination by blue stragglers. Spectra of this S/N provide radial velocities accurate to 15 km s$^{-1}$. The small contamination of the sample can be further reduced by identifying stars with anomalously strong Ca II K absorption for their temperature.  \\begin{table} \\centering \\begin{tabular}{@{}lccccc} \\hline \\\\[-11pt] Name & $EW_{Ca}$ & $\\sigma_{Ca}$ & $(B-V)_0$ & [Fe/H]$_K$ & [Fe/H]$_C$\\\\ \\hline \\\\[-11pt] M92 II-23 & 0.50 & 0.07 & 0.06 & -2.2 & -2.1 \t\\\\ M92 IV-27 & 1.03 & 0.07 & 0.15 & -2.2 & -1.9 \t\\\\ M92 S-20 & 0.65 & 0.07 & 0.13 & -2.2 & -2.3 \t\\\\ M92 S-24 & 0.52 & 0.05 & 0.01 & -2.2 & -1.7 \t\\\\ M92 VI-10 & 0.64 & 0.07 & 0.11 & -2.2 & -2.2 \t\\\\ M92 XII-01 & 1.05 & 0.05 & 0.17 & -2.2 & -2.0 \t\\\\ M92 XII-09 & 0.99 & 0.07 & 0.13 & -2.2 & -1.8 \t\\\\ M92 XII-10 & 0.53 & 0.06 & 0.05 & -2.2 & -1.9 \t\\\\ M3 182 & 1.07 & 0.07 & 0.08 & -1.5 & -1.5 \t\\\\ M3 II-11 & 4.05 & 0.13 & 0.42 & -1.5 & -1.8 \\\\ M3 VI-18 & 1.63 & 0.06 & 0.16 & -1.5 & -1.5 \\\\ HD 2857 & 1.76 & 0.06 & 0.17 & -1.7 & -1.5 \t\\\\ HD 14829 & 0.41 & 0.05 & 0.01 & -2.2 & -1.9 \t\\\\ HD 60778 & 1.11 & 0.02 & 0.08 & -1.4 & -1.4 \t\\\\ HD 74721 & 0.55 & 0.01 & 0.02 & -1.3 & -1.7 \t\\\\ HD 86986 & 0.92 & 0.02 & 0.09 & -1.7 & -1.7 \t\\\\ HD 109995 & 0.64 & 0.01 & 0.05 & -1.7 & -1.8 \t\\\\ HD 161817 & 1.61 & 0.02 & 0.13 & -1.8 & -1.4 \t\\\\ \\hline \\label{tab_cak} \\end{tabular} \\caption{Ca II K line EW and the calculated [Fe/H] ([Fe/H]$_C$) compared to those in KSK ([Fe/H]$_{K}$) for stars shown in Fig. \\ref{cak_plot}.} \\end{table}   A second classification technique, the {\\em Scale width--Shape} method, compares the shapes of the Balmer lines by plotting two of the parameters of the profile fit. We find that with spectra of S/N$=15\\,$\\AA$^{-1}$ samples of halo BHB stars will be $\\sim83\\%$ complete with only $\\sim12\\%$ contamination by blue stragglers. In other words without the need for colours the samples are almost as clean and complete as with the {\\em $D_{0.15}$--Colour} method. This argues that if accurate colours $<3\\%$ are available for the sample the best method is the {\\em $D_{0.15}$--Colour} method. However where the accuracy of the photometry is worse the most efficient method for studying the dynamics of the Galactic halo would be the {\\em Scale width--Shape} method, saving the need to obtain accurate photometry. This may be the best strategy for studying the outer halo of the Galaxy beyond $r=100$ kpc, with candidate BHB stars selected from the SDSS database."
    },
    "0207/astro-ph0207450_arXiv.txt": {
        "abstract": "Recent work by several groups has established the properties of the dwarf satellites to M31. We reexamine the reported kinematics of this group employing a fresh technique we have developed previously. By calculating the distribution of a $\\chi$ statistic (which we define in the paper) for the M31 system, we conclude that the total mass (disk plus halo) of the primary is unlikely to be as great as that of our own Milky Way. In fact the $\\chi$ distribution for M31 indicates that, like NGC 3992, it does not have a massive halo. In contrast, the analysis of the satellites of NGC 1961 and NGC 5084 provides strong evidence for massive halos surrounding both spiral galaxies. ",
        "introduction": "According to Mateo (1998) more dwarf galaxies have been discovered as Local Group members since 1971 than during the preceding 222 years. A large fraction of these systems are bound to M31. Recently, Evans and Wilkinson (2000, EW), Evans et al. (2000) and C\\^{o}t\\'{e} et al. (2000) have independently estimated the total mass associated with the M31 (disk plus halo) and found that it is probably less massive than our Milky Way. Here we present an alternative way of estimating the total mass of M31 using these same data but employing a novel technique we have developed (Erickson, Gottesman and Hunter 1987, EGH1; and 1999, EGH2).  The use of the kinematics of dwarf satellite galaxies to investigate the masses and extents of galaxy halos is a very powerful technique (see EGH1, EGH2, Zaritsky et al. 1993, ZSFW1 and 1997, ZSFW2). Almost all studies have been limited to describing the average properties of a generic galaxy because very few individual systems posses enough satellites to allow unique inferences. Owing to the limitations imposed by projection effects, comparisons between observations and models (see EGH2, and \\S3.1 below) indicate that about seven satellites are required before the signature of a massive halo is likely to be observed. We will discuss this further below. In EGH2 we studied the statistics \\linebreak of over 71 satellite-primary pairs. Unfortunately, very few of the primaries had as many as three satellites. However, NGC 1961 (Gottesman, Hunter and Shostak, 1983; hereafter GHS) and NGC 5084 (Carignan et al., 1997), as individual galaxies, meet the requirements for a substantial number of satellites (at least eight for NGC 1961 and seven for NGC 5084). In the Local Group, M31 is surrounded by a cloud of dwarf systems (see EW, for example) and it is our purpose to examine the kinematics of its satellites using the \\linebreak methods we have developed in our earlier papers, and to compare the properties of M31 with those of NGC 1961 and NGC 5084. In the sections that follow, we will briefly review our methodology, discuss the data for the three galaxies being compared and draw some conclusions about the properties of M31. ",
        "conclusions": ""
    },
    "0207/astro-ph0207385_arXiv.txt": {
        "abstract": "It is widely believed that the optical/UV continuum of quasars (the \"Big Blue Bump\") represents optically thick thermal emission from accretion onto a black hole. Narrow line radio galaxies don't show such a component directly, and were historically thought for that reason to be rotation-powered, with large kinetic luminosity in the radio jets but very little accretion or optical radiation. When the Unified Model came along, identifying at least some narrow line radio galaxies as hidden quasars, the compelling observational motivation for this radio galaxy scenario lost some of its force. However, it is far from clear that all narrow line radio galaxies contain hidden quasar nuclei.  The clear sign of a hidden quasar inside a radio galaxy is the appearance of quasar spectral features in its polarized (scattered) light. However that observational test requires suitably placed scattering material to act as a mirror, allowing us to see the nuclear light. A rather robust and more general test for a hidden quasar is to look for the predicted high mid-IR luminosity from the nuclear obscuring matter.  The nuclear waste heat is detected and well isolated in the nearest narrow line radio galaxy, Cen A. This confirms other indications that Cen A {\\em does} contain a modest quasar-like nucleus. However we show here that M87 {\\em does not}: at high spatial resolution, the mid-IR nucleus is seen to be very weak, and consistent with simple synchrotron emission from the base of the radio jet. This fairly robustly establishes that there are \"real\" narrow line radio galaxies, without the putative accretion power, and with essentially all the luminosity in kinetic form.  Next we show the intriguing mid-IR morphology of Cygnus A, reported previously by us and later discussed in detail by Radomski et al.\\ (2002). All of this mid-IR emission is consistent with reprocessing by a hidden quasar, known to exist from spectropolarimetry by Ogle et al.\\ (1997) and other evidence. ",
        "introduction": "Radio lobe emission is fairly isotropic, so it's easy to make lists of double radio sources that are nearly unbiased with respect to orientation. The visibility of optical point sources in {\\em most}\\/ of the optically dull (weak or low ionization emission line) galaxies show that there are no $\\gtwid$ parsec scale tori present which are able to obscure the unresolved optical sources in those cases (Chiaberge et al.\\ 1999, 2000). Whatever the nature of the M87 optical point source, we know it's small from the variability on year timescales. The same would apply to the other optically dull nuclei if they vary as M87 does.  This would strengthen the Chiaberge et al.\\ argument that we have unobscured sightlines nearly all the way to the central engines.  Since the M87 optical/UV flux is quite variable (e.g., Tsvetanov et al.\\ 1998) jet synchrotron emission is a possibility. By correlating the radio synchrotron core fluxes and the optical point source fluxes in FR I radio galaxies generally, Chiaberge et al.\\ (1999) infer that the latter are in fact likely to be beamed synchrotron sources. Crucial tests of the nature of the optical point sources can be made with spectroscopy and polarimetry. We hope to do this with adaptive optics, excluding most of the starlight that dominates in arcsec apertures.  However, a large minority of low ionization FR I and FR II radio galaxies show no point source, and are similar in this way to AGN with hidden nuclei. In fact the closest FR I, Cen A, {\\em does}\\/ have a big molecular torus (see Fig.~2 of Rydbeck et al.\\ 1993!), and substantial evidence for a hidden nucleus as well (see references cited earlier). As a working hypothesis we might suppose that the same is true for all those without detectable optical pointlike nuclei. (Of course sensitivity of the optical/UV observations also enters in.)  Thus the FR I family is heterogeneous: some contain hidden optical/UV nuclei and some do not. It's been difficult to find FR I objects with strong evidence for hidden AGN, but in fact at least a few are known to be quasar-like from direct spectroscopy (3C120 is well known; see also Lara et al.\\ 1999 and Sarazin et al.\\ 1999). The infrared peaks in some blazars such as 3C 66A are clearly distinct from the radio synchrotron, providing further evidence for FR I objects with hidden thermal emission. Therefore the ``nonthermal\" model {\\em does not} apply to all optically dull or low ionization radio galaxies, or to all FR I galaxies. The FR class has no apparent {\\em direct} relation to the mode of energy production, consistent with much recent evidence that FR Is behave very much like FR IIs at VLBI scales. We will understand this better after the completion of our mid-IR program.  Finally we reiterate an assumption in our program: we assume that any mid-IR dust emission is isotropic to within a factor of a few. ",
        "conclusions": "Radio lobe emission is fairly isotropic, so it's easy to make lists of double radio sources that are nearly unbiased with respect to orientation. The visibility of optical point sources in {\\em most}\\/ of the optically dull (weak or low ionization emission line) galaxies show that there are no $\\gtwid$ parsec scale tori present which are able to obscure the unresolved optical sources in those cases (Chiaberge et al.\\ 1999, 2000). Whatever the nature of the M87 optical point source, we know it's small from the variability on year timescales. The same would apply to the other optically dull nuclei if they vary as M87 does.  This would strengthen the Chiaberge et al.\\ argument that we have unobscured sightlines nearly all the way to the central engines.  Since the M87 optical/UV flux is quite variable (e.g., Tsvetanov et al.\\ 1998) jet synchrotron emission is a possibility. By correlating the radio synchrotron core fluxes and the optical point source fluxes in FR I radio galaxies generally, Chiaberge et al.\\ (1999) infer that the latter are in fact likely to be beamed synchrotron sources. Crucial tests of the nature of the optical point sources can be made with spectroscopy and polarimetry. We hope to do this with adaptive optics, excluding most of the starlight that dominates in arcsec apertures.  However, a large minority of low ionization FR I and FR II radio galaxies show no point source, and are similar in this way to AGN with hidden nuclei. In fact the closest FR I, Cen A, {\\em does}\\/ have a big molecular torus (see Fig.~2 of Rydbeck et al.\\ 1993!), and substantial evidence for a hidden nucleus as well (see references cited earlier). As a working hypothesis we might suppose that the same is true for all those without detectable optical pointlike nuclei. (Of course sensitivity of the optical/UV observations also enters in.)  Thus the FR I family is heterogeneous: some contain hidden optical/UV nuclei and some do not. It's been difficult to find FR I objects with strong evidence for hidden AGN, but in fact at least a few are known to be quasar-like from direct spectroscopy (3C120 is well known; see also Lara et al.\\ 1999 and Sarazin et al.\\ 1999). The infrared peaks in some blazars such as 3C 66A are clearly distinct from the radio synchrotron, providing further evidence for FR I objects with hidden thermal emission. Therefore the ``nonthermal\" model {\\em does not} apply to all optically dull or low ionization radio galaxies, or to all FR I galaxies. The FR class has no apparent {\\em direct} relation to the mode of energy production, consistent with much recent evidence that FR Is behave very much like FR IIs at VLBI scales. We will understand this better after the completion of our mid-IR program.  Finally we reiterate an assumption in our program: we assume that any mid-IR dust emission is isotropic to within a factor of a few."
    },
    "0207/astro-ph0207666_arXiv.txt": {
        "abstract": "We present the results of a search for submillimetre-luminous host galaxies of optically dark gamma-ray bursts (GRBs) using the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT).  We made photometry measurements of the 850-$\\mu \\rm{m}$ flux at the location of four `dark bursts', which are those with no detected optical afterglow despite rapid deep searches, and which may therefore be within galaxies containing substantial amounts of dust. We were unable to detect any individual source significantly. Our results are consistent with predictions for the host galaxy population as a whole, rather than for a subset of dusty hosts.  This indicates that optically dark GRBs are not especially associated with very submillimetre-luminous galaxies and so cannot be used as reliable indicators of dust-enshrouded massive star-formation activity.  Further observations are required to establish the relationship between the wider GRB host galaxy population and SCUBA galaxies. ",
        "introduction": "\\label{sec:intro}   \\subsection{Gamma-ray Bursts} \\label{subsec:grbtheory}  It is generally now accepted that the afterglow emission resulting from (long-duration, soft-spectrum) gamma-ray bursts (GRBs) can be explained by an ultra-relativistic shock-wave expanding into a surrounding medium \\citep{mesrees,GRBreview,mesrev}.  The precise nature of the progenitor systems is not a settled issue. The two most popular theories both involve stellar remnants: the collapsar/hypernova model, in which a single massive progenitor star undergoes core collapse \\citep{woosley,microquasar}; and the binary merger theory, in which two massive stellar remnants, such as neutron stars, merge \\citep{lattimer,binary}. Both of these scenarios may be able to explain the energetics of the explosion that produces the GRBs, particularly if they are beamed, since the energies of GRBs are comparable to those involved in the formation of typical stellar-progenitor black holes.  Recently, the collapsar/hypernova model has gained support from three sets of observations: firstly, it was confirmed that the positions of some GRBs accurately localized by the observation of optical and/or radio afterglows, were found to be within star-forming regions of their host galaxies \\citep{offset}, which themselves are frequently starburst galaxies \\citep{sokolov}.    This is supported by X-ray determinations of the HI column density along the line of sight to GRBs, which is consistent with their residing in GRBs \\citep{galama}.  During their creation, it is likely that the massive stellar remnants required in the binary merger scenario would receive a substantial `kick' velocity, so that the merger event causing the GRB would take place outside of the star-forming region of the host galaxy \\citep{galama}. Also, the delay required between formation of the remnants and their merger may well be long enough for star formation to have ceased in the host, suggesting that the host galaxy would no longer be luminous.  Hence the position measurements of \\citet{offset} support the single massive progenitor theory. Note though, that one should be wary of a potential selection effect: the detection of a GRB afterglow requires a certain minimum density in the surrounding interstellar medium (ISM), and so GRBs may only be identifiable when they occur inside galaxies \\citep{chevalier}.  Secondly, in a few cases, optical afterglows of GRBs have been seen to contain a secondary brightening in flux after a few weeks \\citep{sn980425,sn980326,reichart, galamaII,000911SN}. This flux increase, though often not very significant, has been attributed to a supernova occurring simultaneously with the GRB, although alternative ideas have been postulated, such as dust echoes \\citep{esin} or interactions of the GRB shock wave with wind-driven density structures in the surrounding ISM \\citep{EnricoIII}. Supernovae associated with GRBs are compatible with the single progenitor collapsar/hypernova model rather than the merger hypothesis.  Finally, the detection of iron features in the X-ray afterglows suggest the presence of an iron-enriched ISM surrounding the GRB progenitor \\citep{amati,piro,yoshida}. This high iron mass is consistent with ejecta from a massive stellar progenitor.  If the hypernova theory is the correct explanation for long-duration GRBs, then there should be a direct link between GRBs and high-mass star formation activity. Since the gamma-ray emission from the initial explosion is not attenuated by dust, and can be detected from high redshifts, GRBs should be unbiased tracers pointing to star formation activity wherever massive stars are living and dying \\citep{krumholz,BlainPriya,b32,010222,EnricoII}.  \\subsection{SCUBA galaxies} \\label{subsec:SCUBAgals}  Since the commissioning of the SCUBA instrument \\citep{wayne} in 1997 on the JCMT on Mauna Kea, a new era of submillimetre (submm) cosmology has been possible (see, e.g. \\citealt{bsik,s38}).  In particular, the discovery of a substantial population of dust-enshrouded `SCUBA galaxies' has launched a debate regarding observational estimates of the global star formation rate in the Universe.  SCUBA traces the interstellar dust in galaxies (with a temperature of the order of tens of degrees Kelvin), which may be heated by the UV light emitted by OB stars and/or a hosted AGN. The relative contributions  of starlight and AGN to dust-heating in a particular galaxy are hard to determine precisely, but studies using \\emph{Chandra} to observe hard X-rays from submm-selected galaxies suggest that about 20 per cent of the sample contain a detectable AGN \\citep{bautz,fabian,chandra8mjy}.  Therefore, it is likely that the luminous dust emission from SCUBA galaxies is powered predominantly by star formation. In this case, because of the large bolometric luminosities of the SCUBA galaxies (at least $10^{12} L_{\\sun}$), they make a substantial contribution to overall star formation activity in the high-redshift universe, comparable to or greater than that of optically-selected galaxies \\citep{bsik}.  Thus, they are good candidates to be hosts of many GRBs, if GRBs are indeed associated with the death of high-mass stars.  The optical properties of SCUBA galaxies, where they are well-located, indicate that most are very faint, with $R\\,>\\,25$ \\citep{s38}, not dissimilar to the GRB host galaxy population as a whole \\citep{970828}. It is likely that about 10--20 per cent of the star-formation activity in the high-redshift universe takes place in submm galaxies brighter than 2\\,mJy.  As a result, perhaps 1 in 5 GRBs should reside in such objects.  This point is examined further in Section \\ref{sec:results}.   \\subsection{Submm observations of GRBs}  This paper describes our submm SCUBA observations of the host galaxies of four GRBs.  To increase the chances of finding dusty SCUBA galaxies in this pilot program, we selected GRBs which were `dark' in the sense that their afterglows were undetected in the optical despite deep, rapid searches.  Estimates vary, but optical afterglows have been searched for and not found in roughly 30-50 per cent of GRBs with X-ray afterglows. In many cases the non-detection may simply be because the searches were not deep or rapid enough (Galama \\& Wijers 2001), but in some instances it is clear that the optical afterglows are genuinely underluminous (Groot et al. 1998, van Paradijs et al. 2000, Lazzati et al. 2001a, Ramirez-Ruiz et al. 2002). An obvious possibility is that these bursts are heavily obscured by dust. Since GRBs are expected to destroy any dust in their vicinity \\citep{waxdrain,fruchter,reichyost,bram}, the obscuration would have to be due to dust elsewhere along the line of sight, as might well be expected in very dusty submm-bright galaxies.  Note that we did not attempt to observe GRB afterglows in the submm. Target of Opportunity programs to observe GRB afterglows are underway at the JCMT and IRAM 30-m telescopes \\citep{SmithI,SmithII,010222}. Our observations took place long enough after the initial explosion for only the host galaxy emission to be detectable.  There have been two previous SCUBA detections of a GRB host galaxy.  The first was found serendipitously during a submm afterglow search \\citep{010222} for a GRB that had an optical afterglow (GRB 010222). This suggests that the GRB was located in a luminous dusty galaxy, but either in a relatively dust-free region or near the edge of the galaxy, so that the explosion could clear out any dust along the line of sight.  Also, \\citet{b32} reported the results of a successful targeted search for submm emission from  GRB 000418.  This too had an optical transient identification.  Another relevant detection was that of a GRB host in the radio (GRB 980703; \\citealt{BergerIII}), whose flux was claimed to be caused by a large star-formation rate rather than AGN activity, based on optical spectroscopy and the absence of radio variability. However, a deeply embedded AGN within a dust- and gas-rich galaxy could plausibly contribute the emission from this galaxy, especially as the radio source is located very near to the nucleus of the host galaxy. ",
        "conclusions": "\\label{sec:conc}  From our small sample of four reliably-identified dark GRBs, we find that the `dark bursts' do not preferentially select dusty host galaxies with very significant amounts of star formation. Looking at some members of our sample, we can explain the lack of optical transients for other reasons. It seems likely that to characterise the optically dark bursts as a physically distinct population of GRBs would be misleading. In each case different circumstances due to a combination of the observing conditions and the physical conditions at the location of the GRB could give rise to the lack of an optical afterglow.  If, instead, we view the GRBs included here as four examples of the overall population, we find that the results agree with the distribution predicted assuming that GRBs tracew high-mass star formation.  Our lack of a single strong detection implies that no more than 20 per cent of GRB hosts are submm galaxies detectable at a flux density brighter than 5\\,mJy using SCUBA. In interpreting the results, however, the small sample size inevitably makes it hard to draw any solid conclusions.  Combining our three most significant results with the observations of \\citet{SmithI,SmithII}, we find that 11 GRB locations have been observed with SCUBA to noise levels $<$ 2mJy, and in all but two (GRB 980329, \\citealt{SmithI} and GRB 000210, this paper) detection of submm emission from a host galaxy was ruled out. The two GRB hosts detected by SCUBA \\citep{010222,b32} had afterglows with optical transients.  Therefore, while some GRBs definitely are located in dusty galaxies, the route to selecting these GRBs on the basis of their afterglow data is not yet clear.  Since GRBs can only clear dust out to less than 100 pc distance \\citep{fruchter}, it would be surprising if optically-selected samples of GRBs are not generally biased against dusty hosts.  However, as noted above, there may be alternative explanations for the optical faintness of some afterglows beyond location in a dusty host galaxy.  Our results indicate that radio-located optically-dark bursts seem not to be reliable indicators of luminous, dusty host galaxies.  It may be that the physical conditions of the ISM in the densest star forming regions are incompatible with the generation of intense radio emission from GRB shocks.  In that regard, it would ultimately be interesting, and certainly possible in the SWIFT era, to study a sample of hosts of purely X-ray selected GRBs, with the hope that the prompter X-ray emission would be less affected by the wider environment of the progenitor.  From the SCUBA observations alone, we cannot separate the possibility that there is little dust in these systems from the possibility that there is a lot of dust but insufficient UV photons to make it glow brightly. Alternatively, the dust may be heated to high temperatures and so cannot be detected by SCUBA at all. The predicted flux of a $5 \\times 10^{12} \\rm{L}_{\\sun}$ galaxy at $z = 1$ decreases from 20\\,mJy to 0.08\\,mJy as the dust temperature increases from 20\\,K to 80\\,K (see Fig.\\,5 in \\citet{blain_rev}). Typical dust temperatures for submm-selected objects are thought to be 40\\,K \\citep{ivison2000}, but a temperature $\\sim 60$\\,K would be hot enough to prevent such a galaxy being detected above the 2\\,mJy confusion limit of SCUBA.  However, optical colours of the host galaxies of GRBs 970828, 981226 and 000210 \\citep{970828,981226VLA,p15} show only modest reddening, to the optical depths that can be probed, and thus perhaps imply little dust overall in the hosts. Ideally, high-resolution observations to locate dust-enshrouded star-formation activity with respect to the location of GRBs need to be made.  The results presented here suggest that a more extensive deep submm survey of GRB hosts is necessary to investigate their far-infrared and submm properties in detail, and this is now underway at the JCMT (Barnard et al., in prep.)."
    },
    "0207/astro-ph0207516_arXiv.txt": {
        "abstract": " ",
        "introduction": "Danescu et al.\\ (1999) have compiled space motions for 38 globular clusters, and 35 of these have [Fe/H] values less than $-1.0$. Since that time 3 more halo clusters (NGC6522, NGC7006, \\& Pal 13) have been added to the compilation. This sample of 38 clusters is the basis for the following kinematical investigation.  The tabulated UVW components of space motion (corrected solar peculiar motion and for Galactic rotation of 220 km~sec$^{-1}$) are referred to to the Galactic center where $V_{\\pi}$ is the radial component in the Galactic plane directed positive away from the center, $V_{\\phi}$ is the tangential component positive in the direction of Galactic rotation and $V_{z}$ is the vertical component directed positive in the direction of the north Galactic pole.  The average of each of these components for the entire sample was removed before forming the following six moments: $\\overline{V_{\\pi}V_{\\pi}}, \\overline{V_{\\phi}V_{\\phi}}, \\overline{V_{z}V_{z}}, \\overline{V_{\\pi}V_{\\phi}}, \\overline{V_{\\pi}V_{z}}$, and $\\overline{V_{\\phi}V_{z}}$.  These six moments are the components of the symmetrical velocity dispersion tensor.  The eigenvalues of the matrix were then determined in the standard way and are given below.  The units of the eigenvalues denoted by $\\sigma$ are km~sec$^{-1}$ and the one sigma uncertainties throughout were computed by the bootstrap method.  \\begin{center} Clusters with [Fe/H]$<-1.00$\\ (n=38) \\end{center} \\vspace{-0.3 in} \\begin{eqnarray} \\sigma_{a} = 161^{+18.3}_{-15.0} & l = 338^{+10.5}_{-6.65} & b = -5.65^{+13.6}_{-15.0} \\nonumber \\\\ \\sigma_{b} = 120^{+9.44}_{-15.5} & l = 87.2^{+155.}_{-30.0.} & b = -73.5^{+27.9}_{-6.63} \\\\ \\sigma_{c} = 95.6^{+2.66}_{-17.6} & l = 66.1^{+11.2}_{-10.8} & b = 15.4^{+22.2}_{-21.1} \\nonumber \\\\ & \\overline{V_{\\phi}} = 59.6^{+13.5}_{-17.9} \\nonumber \\end{eqnarray}  \\noindent As a check on the sensitivity of the above result to the [Fe/H] cutoff, we repeated the procedure for the 26 clusters with [Fe/H]$<-1.5$ and find:  \\begin{center} Clusters with [Fe/H]$<-1.50$\\ (n=26) \\end{center} \\vspace{-0.3 in} \\begin{eqnarray} \\sigma_{a} = 172^{+17.9}_{-19.2} & l = 330^{+11.4}_{-8.11} & b = -6.31^{+14.9}_{-14.5} \\nonumber \\\\ \\sigma_{b} = 122^{+8.45}_{-27.7} & l = 179^{+91.7}_{-112.} & b = -82.8^{+32.6}_{+6.58} \\\\ \\sigma_{c} = 94.7^{+5.76}_{-27.8} & l = 60.8^{+12.1}_{-12.1} & b = -3.48^{+30.1}_{-21.3} \\nonumber \\\\ & \\overline{V_{\\phi}} = 49.3^{+23.7}_{-28.1} \\nonumber \\end{eqnarray}  In both of the above solutions the mean streaming motions in the $\\pi$ and $z$ directions differed from zero by less than 1.5 sigma. Also the above values of $\\overline{V_{\\phi}}$ are independent of distance from the Galactic center.  Asymmetric error ellipsoids in the components of space motion can be expected due to relatively small errors in radial velocity and relatively large errors in proper motion. In order to ensure that such effects were not unduly influencing solution (1) the errors in radial velocity, proper motion and distance given by Dinescu et al.\\ (1999) were used as dispersions in assumed gaussian distributions and solutions were performed on space motions determined by the propagation of randomly chosen uncertainties in each of the above quantities. The mean and standard deviation of the major axis from 100 trials with 38 clusters is $\\sigma_{a} = 167 \\pm 7.98$ towards $l = 332 \\pm 6.28$ and $b = -5.94 \\pm 7.64$. Results for the other two axes are $\\sigma_ {b} = 122 \\pm 10.9$ and $\\sigma_{c} = 108 \\pm 10.1$. This independent assessment of uncertainties thus strengthens the case for the generally $prolate$ velocity ellipsoid given by solution (1).  An examination of the spatial distribution of the clusters from solution (1) shows that only 7 of the 38 clusters were in the hemisphere beyond the Galactic center. In order to assess the effect of this spatial imbalance on our result, kinematic solutions were made using the above 7 clusters combined with 7 independently chosen clusters from the hemisphere containing the sun. Thirty solutions were made with 14 clusters each and the median results for each aspect of the major axis are: $\\sigma_{a} = 166^{+20}_{-9},\\ \\ l = 348^{+8}_{-7},\\ \\ $ and $b = -13^{+29}_{-8}$. While the one sigma uncertainties of solution (1) do overlap with the above result, perhaps the most important implication of this exercise is a clear need for more cluster space motion determinations. ",
        "conclusions": "Both the velocity ellipsoid and the spatial distribution of the most metal poor Galactic globular clusters appear to be anisotropic.  It is interesting that the major axis of the velocity ellipsoid is directed towards low Galactic latitude, consistent with the direction of the short axis of the spatial distribution and not far from the pole of the outlying satellite galaxy distribution.  A quantitative comparison for the different samples is given in Table 1 where for reference we take the major axis of the kinematic solution (1).  \\begin{deluxetable}{ccccccc} \\footnotesize \\tablenum{1} \\tablecolumns{7} \\tablecaption{Comparison of the Axes of the Various Samples} \\tablehead{ \\colhead{Sample} & \\colhead{Number} & \\colhead{Solution} & \\colhead{ Axis } & \\multicolumn{2}{c}{Galactic Coordinates} & \\colhead{$\\Delta\\theta$\\tablenotemark{a}} \\\\ \\colhead{} & \\colhead{n} & \\colhead{number} & \\colhead{Major/Minor} & \\colhead{~~~~~$l$} & \\colhead{~~~~~$b$} & \\colhead{} } \\startdata Velocity ellipsoid & 38 & 1 & Major &~$338^{+10.5}_{-6.65}$&~$-5.65^{+13.6}_{-15.0}$& $0^{\\circ}$ \\nl [Fe/H]$\\leq-1.0$ & & & & & & \\nl \\\\ Velocity ellipsoid & 26 & 2 & Major &~$330^{+11.4}_{-8.11}$&~$-6.31^{+14.9}_{-14.5}$&$8.0^{\\circ}\\pm13$ \\nl [Fe/H]$\\leq-1.5$ & & & & & & \\nl \\\\ Clusters & 94 & 3 & Minor &~$326^{+13.9}_{-12.2}$&~$-10.8^{+7.22}_{-13.5}$&$13^{\\circ}\\pm16$ \\nl [Fe/H]$\\leq-1.0$ & & & & & & \\nl \\\\ Clusters & 60 & 5 & Minor &~$337^{+10.6}_{-39.4}$&~$-22.1^{+12.7}_{-10.3}$&$16^{\\circ}\\pm18$ \\nl [Fe/H]$\\leq-1.5$ & & & & & & \\nl \\\\ Clusters & 66 & 7 & Minor &~$327^{+14.4}_{-11.3}$&~$-9.65^{+6.99}_{-12.6}$&$11^{\\circ}\\pm16$ \\nl [Fe/H]$\\leq-1.0$, $|b|>10^{\\circ}$ &   & &  & & & \\nl \\\\ Outlying Satellites & 10 & H2000 & Minor &~$336^{+8.28}_{-3.79}$&~$11.1^{+3.89}_{-5.05}$&$17^{\\circ}\\pm15$ \\nl \\\\ Clusters & 33 & 9 & Major &~$338^{+6.19}_{-29.0}$&~$-6.09^{+4.63}_{-2.92}$&$0.44^{\\circ}\\pm15$ \\nl [Fe/H]$\\geq-0.7$ & & & & & & \\nl \\enddata  \\tablenotetext{a}{$\\Delta\\theta$ is the angular difference in degrees between axes of the first sample and each of the subsequent samples. } \\end{deluxetable}  There would appear to be no straight forward explanation for the observations above. If the gravitational potential was spherical, one can argue that the direction of highest velocity dispersion would also be that of the longest spatial axis and vice versa. The above result then apparently rules out a spherical potential but would be consistent with one that was flattened in the direction of highest velocity dispersion. Recalling that the age-metallicity relation for the most metal poor globular clusters is extremely steep (vandenBerg 2000), a naive picture is that an initially flattened filament underwent a rapid collapse tranverse to the long axis during which time satellite galaxies, the metal poor clusters, and the triaxial structure discussed in H2000 were formed. Simultaneously the Galaxy was being assembled by a slower collapse along the filament during which time most of its angular momentum was being induced and/or redistributed. While triaxial structures are formed in cold dark matter simulations, the scenario outlined above may be more consistent with a warm dark matter theory for structure formation where the lowest mass objects form from the fragmentation of caustics (c.f. Bode, Ostriker, \\& Turok 2001)."
    },
    "0207/hep-ph0207090_arXiv.txt": {
        "abstract": "I review recent constraints on and claimed detections of a time-varying fine-structure constant $\\alpha$. Our results are consistent with no variation in $\\alpha$ from the epoch of recombination to the present day, and restrict any such variation to be less than about $4\\%$. The forthcoming MAP and Planck experiments will be able to break most of the currently existing degeneracies between $\\alpha$ and other cosmological parameters, and measure $\\alpha$ in the early universe with an accuracy comparable to current claimed recent-universe detections. ",
        "introduction": "The search for observational evidence of variations in the `fundamental' constants that can be measured in our four-dimensional world is an extremely exciting area of current research, with several independent claims of detections in different contexts emerging in the past year or so, together with other improved constraints. See \\cite{intro} for a non-technical introduction, and \\cite{uzan} for an extensive compilation of existing constraints.  The importance of this search can not be over-emphasized. Varying non-gravitational constants are forbidden by General Relativity and all metric theories of gravity. A varying $\\alpha$ violates the Equivalence Principle, signaling the breakdown of gravity as a purely geometric phenomenon and proving the existence of additional gravitational fields (apart from the metric) in the universe. Indeed, many models proposed as explanations for varying $\\alpha$ are {\\it already} ruled out by the standard gravitational tests alone, regardless of any predictions and experimental results on $\\alpha$ itself. ",
        "conclusions": "We have shown that the currently available data is consistent with no variation of $\\alpha$ from the epoch of recombination to the present day, though interestingly enough the CMB and LSS datasets seem to prefer, on their own, variations of $\\alpha$ with opposite signs. Whether or not this statement has any physical relevance (beyond the results of the statistical analysis) is something that remains to be investigated in more detail. In any case, any such (relative) variation is constrained to be less than about $4\\%$.  The prospects for the future are definitely bright. The forthcoming satellite experiments (including data on polarization and temperature-polarization cross-correlation) will provide a dramatic improvement on these results. These tools, together with other measurements coming from BBN \\cite{Avelino:2000ea} and quasar and related data \\cite{Murphy:2000pz,Webb:2000mn} offer the exciting prospect of being able to map the value of $\\alpha$ at very many different cosmological epochs, including early-universe constraints with an accuracy comparable to the currently existing recent-universe ones \\cite{graca}. This should allow us to impose very tight constraints on higher-dimensional models where these variations are ubiquitous."
    },
    "0207/astro-ph0207046_arXiv.txt": {
        "abstract": "In the lenticular galaxy NGC~1023 a new population of star clusters (``faint fuzzies'') was recently discovered by Larsen \\& Brodie.  These clusters are found inside the disc and are faint ($23 \\leq V \\leq 24$~mag) and extended with effective radii of $r_{\\rm eff} \\approx 7$ to $15$~pc.  \\\\ We present here N-body calculations of a likely formation-scenario through merging star clusters in clusters of star clusters (super-clusters).  Such super-clusters are observed to form in interacting galaxies.  \\\\ The resulting merger objects have masses comparable to the ``faint fuzzies'' and show large effective radii ($r_{\\rm eff} > 7$~pc).  Even though these objects are suffering from strong tidal forces they are able to survive and reach the estimated ages of the extended star clusters in NGC~1023. ",
        "introduction": "\\label{sec:intro}  During a recent search for globular clusters (GC) in NGC~1023, a lenticular galaxy at approximately $9$~Mpc distance, Larsen \\& Brodie (2000) found faint objects ($23 \\leq V \\leq 24$~mag) that have GC luminosities but effective radii ($=$ half-light radii) of $7$ to $15$~pc which is much larger than ordinary GCs that have effective radii of $2$ to $3$~pc.  Using spectra obtained with the Keck telescope they (Brodie \\& Larsen 2002) confirmed that these objects belong to NGC~1023 and have ages of about $\\geq 7$ to $8$~Gyr, are moderately metal rich ($[\\rm Fe/H] = -0.58 \\pm 0.24$) while showing a super-solar $\\alpha$-element abundance ($[\\alpha/{\\rm Fe}] = +0.3$ to $+0.6$).  While GC are usually connected with the spherical component of a galaxy, these new objects appear to be situated in the disc of NGC~1023.  The distribution of these objects shows rotation around the galaxy. The authors conclude that they have found a new class of star clusters which they call ``faint fuzzies''. \\\\  NGC~1023 has a dwarf companion (NGC~1023A) with which it is steadily interacting, and due to its lenticular appearance it probably suffered from interactions and mergers in the past leading to strong star formation events. \\\\  Interactions between gas-rich galaxies lead to strong star formation events.  For example, HST-images of the Antennae reveal (Whitmore et al.\\ 1999) that the knots of intense star formation produce clusters of massive young star clusters.  These ``super-clusters'' ($=$ clusters of star clusters; not to be confused with super stellar clusters (SSC), which are individual massive star clusters) appear to contain dozens to hundreds of massive star clusters within a region spanning a few hundred pc to a kpc in diameter. \\\\  Here we report N-body results concerning the dynamical evolution of such super-clusters.  The resulting merger objects have similar properties as the ``faint fuzzies''. ",
        "conclusions": "\\label{sec:conclus}  S0 (lenticular) galaxies are commonly believed to have suffered from interactions with other galaxies which transformed their morphology from a late-type spiral into their early-type form (Kennicutt 1998; Abraham \\& van den Bergh 2001).  Therefore the type of NGC~1023 tells us that this galaxy suffered maybe several but at least one strong interaction (merger) with another galaxy. As is seen in interacting galaxies these events lead to an enhanced star formation activity and to an enhanced rate of super nova explosions of massive stars which accounts for the super-solar enhancement of $\\alpha$-elements.  Many of the new stars mostly form in star clusters which are also grouped in super-clusters, as is demonstrated by the Antennae and Stephan's Quintet (Gallagher et al.\\ 2001). These super-clusters are bound entities and form, through successive mergers of their constituent star clusters, new objects which are larger than ordinary GCs.  Our simulations show that even strong tidal fields are not able to destroy these objects and they survive for a Hubble-time.  \\\\  The ``faint fuzzies'' found in NGC~1023 may have therefore formed as super-clusters during the interactions and the resulting star bursts that shaped NGC~1023.  Like in the Antennae the super-clusters were born within the galaxy and not in the outer tidal tails.  Therefore the ``faint fuzzies'', which one can call extended globular clusters or small ultra-compact dwarf galaxies by virtue of their long two-body relaxation time, stay within the optically visible galaxy and show a similar rotation pattern like the disc.  These objects can be old because they are able to survive until the present, they may be rotating and have different populations of stars if old field stars were captured by the young super-cluster. \\\\  Given the ubiquity of young super-clusters in the Antennae and Stephan's Quintet, it appears reasonable to expect that such objects form whenever strong star bursts occur.  It is therefore exciting, but retrospectively not very surprising, that ``faint fuzzies'' were discovered in a lenticular galaxy. \\\\  Super-clusters that form in outer tidal tails may leave the mother galaxy either by being ejected by the merging galaxies or become unbound due to a tidal field from the galaxy cluster.  Such merger objects will appear similar to the UCDs in Fornax (Fig.~\\ref{fig:evol}).  A related object may be $\\omega$-Cen in our Milky Way, which is quite different and more massive than the other GCs.  It incorporates at least two different populations of stars and shows a rotation-law resembling that of our merger object (Freeman 2001). $\\omega$-Cen is quite old and could therefore be a relic of a star-burst during the formation of the Milky Way.  We are investigating the hypothesis that $\\omega$-Cen may also be related to super-clusters.  Fig.~\\ref{fig:evol} summarises in a schematic way the possible evolutionary paths and resulting objects that may form under this hypothesis."
    },
    "0207/astro-ph0207270_arXiv.txt": {
        "abstract": "{Recently, several ways of obtaining observational proof of the existence of black-hole horizons have been proposed. We argue here that such proof is fundamentally impossible: observations can provide arguments, sometimes very strong ones, in favour of the existence of the event horizon, but they cannot prove it. This applies also to future observations, which will trace very accurately the details of the spacetime metric of a body suspected of being a black hole. } ",
        "introduction": "It is generally believed that the compact component in X-ray binary systems is either a star possessing a material surface (a neutron star or a quark star) or a black hole, i.e. an object whose surface is formed by an event horizon. The evidence of the presence of a material surface is obtained from two types of observations. First, stable periodic pulsations of the X-ray emission indicate the presence of a strong ($10^{9} - 10^{15}$ G) magnetic field, which by the virtue of the ``no-hair\" theorem excludes the presence of a black hole in the system. Second, observations of X-ray bursts - thermonuclear explosions occurring in matter accumulated at the surface of the compact object are an obvious proof of the absence of an event horizon.  Although X-ray pulsations or X-ray bursts indicate the presence of a solid surface, their absence does not prove the presence of an event horizon. ``Absence of evidence is not evidence of absence\". However, there is direct evidence that compact bodies in X-ray binaries form (at least) two types of objects: their masses show a bimodal distribution (see Miller, Shahbaz \\& Nolan \\cite{msn98}). Neutron star masses are all concentrated around the ``canonical\" value of 1.4~\\msol\\ whereas the second class of bodies, usually called ``black hole candidates\" have higher masses in the range of $\\sim$ 5 to 18 \\msol (see e.g. Narayan, Garcia \\& McClintock \\cite{ngm01} and Greiner, Cuby \\& McCaughrean \\cite{gcmc01}).  The reason for suspecting the more massive bodies of being black holes is that their masses are higher than the maximum mass of a neutron (or quark) star, which is never larger than $\\sim 3$ \\msol\\  (see e.g. Salgado et al. \\cite{setal94}).  In general, the maximum mass of a compact body can be expressed as $8.4\\left(\\varepsilon_0/10^{14} {\\rm g\\ cm^{-3}} \\right)^{-1/2}$\\msol, where $\\varepsilon_0$ is the fiducial density above which the equation of state is taken to be described by a causality-limit equation of state (Rhoades \\& Ruffini \\cite{rr74}; Friedmann \\& Ipser \\cite{fi87}). Bahcall, Lynn \\& Selipsky (\\cite{bls90}) showed that stars with a material surface can have masses as high as 10 \\msol, if one is willing to entertain configurations of sub-nuclear density. These are the so-called ``Q-stars.\" The mean-field description of nuclear interactions given by Bahcall et al. (\\cite{bls90}) allows baryonic matter to have densities this low.   Although, as shown by Miller et al. (\\cite{msn98}), it is unlikely that bodies with masses larger than 10 \\msol\\ are Q-stars---because this would require unrealistically low densities at which hadronic bulk matter would persist---``unlikely\" is not a very satisfactory argument in favour of the black-hole existence. One would rather wish a ``positive'' proof of the event-horizon's existence. This has been attempted by Ramesh Narayan and collaborators. The claim is that properties of Advection Dominated Accretion Flows (ADAFs; Abramowicz et al. \\cite{aetal95}; Narayan \\& Yi \\cite{ny94}, \\cite{ny95a},\\cite{ny95b}) can be used to prove the existence of event-horizons. ",
        "conclusions": "We have shown that it is fundamentally impossible to give an observational proof for the existence of a black-hole horizon. One could argue that it is not necessary to give such a proof -- a black hole is a specific space-time metric, whose properties can, in principle, be determined through observations. If so, no `direct' proof would be necessary, assuming one could determine that the spacetime around a compact object corresponds to the Kerr solution of the Einstein equations.  One way to distinguish a black hole from a rotating star is through the study of orbital and other frequencies (e.g., epicyclic) of accreting matter moving in strong-field gravity (Kato \\cite{sk01}; Wagoner \\cite{wso}; Abramowicz \\& Klu{\\' z}niak \\cite{ak01} \\cite{ak02}). Another method of determining the space-time geometry is by observation of the energy spectrum reflected from an accretion disc deep in the gravitational well of a compact object (Fabian et al. \\cite{frsw}, \\cite{firy}). Finally, the capture of stellar-mass compact objects by supermassive black holes in galactic nuclei produces gravitational radiation whose properties reflect the structure of black-hole space-time (Ryan \\cite{r95}; Hughes \\cite{sc01}). Such gravitational radiation could be observed by the gravitational-wave antenna LISA.  The last method is the most powerful since it uncovers the compact body's multipolar structure. Like the other methods, however, it suffers from a fundamental weakness because it assumes that only a black hole can be the `source' of the ($a\\neq 0$) Kerr metric. Although it was shown that it is very unlikely that other sources exist (Abramowicz, Lasota \\& Muchotrzeb \\cite{alm}), such a possibility cannot be excluded.  Nevertheless, the case for the existence of black holes in the Universe is very strong and the evidence very convincing. We think, however, that a shadow of doubt will always cast its pall on our certainty in this matter. But it is a fertile doubt: it has already inspired new ideas and will surely continue to do so."
    },
    "0207/astro-ph0207336_arXiv.txt": {
        "abstract": "I present the latest results from a search for hot molecular cores at mid-infrared wavelengths from the largest optical telescopes available at the present time. Three well-observed hot molecular cores were imaged, G29.96-0.02, G19.61-0.23, and G34.26+0.15. Even though mid-infrared sources have been claimed to be detected previously at the hot molecular core locations of both G19.61-0.23 and G34.26+0.15, only the hot molecular core in G29.96-0.02 resulted in a detection. New upper limits on mid-infrared emission are given for the hot molecular cores that were not detected. ",
        "introduction": "Hot molecular cores (HMCs) are believed to represent the earliest stages of massive stellar birth, however there are very few observations of these objects to date. The basic characteristics of hot molecular cores are: 1) they are compact sources seen in radio wavelength ammonia (or molecular line) images but have little or no radio continuum emission; 2) they lie in massive star forming regions near ultracompact HII (UC HII) regions; 3) they are too young and/or embedded to be seen in the optical or near infrared; and 4) they are often coincident with water maser emission. Only a small number of sources exist that have had such a holistic set of observations performed.  Cesaroni et al. (1994) argue that the dust and gas in these HMCs are well-mixed and that the gas temperatures derived from ammonia line imaging should be an approximation to the dust temperatures. Since emission from a 160 K blackbody peaks at 18 $\\micron$ and these HMCs have gas temperatures of 50-200 K, the mid-infrared is a good wavelength regime to try to detect these types of sources.  I present here sub-arcsecond 18 $\\micron$ images of the sites of three hot molecular cores. All three sites are certified massive star forming regions due to the presence of UC HII regions. All sites contain water masers that are isolated from centimeter radio emission sources. Furthermore, all three sites have been found to contain ammonia emission coincident with the water maser locations, as observed through high-resolution molecular line imaging.  The 18 $\\micron$ observations of these HMCs were performed at the Gemini North and W.M. Keck II telescopes using OSCIR. The reasons for using such large telescopes are two-fold: 1) to separate the mid-infrared emission of the HMC from that of any other nearby sources, like extended UC HII regions; and 2) to exploit the higher sensitivity of the larger telescopes so that detections are made more readily, or so that stringent upper limits can be placed on the mid-infrared emission coming from those sources not detected. ",
        "conclusions": "Three well-observed hot molecular cores were imaged at 18 $\\micron$, G29.96-0.02, G19.61-0.23, and G34.26+0.15. Only the hot molecular core in G29.96-0.02 resulted in a detection. This was a surprise considering the previous ``detections'' of HMCs in the sites of both G19.61-0.23 and G34.26+0.15. This small research program has shown that the mid-infrared can be useful in observing and studying some HMCs. It has also shown that the resolution and sensitivity afforded by the 8 to 10-m class telescopes are important to perform this research effectively."
    }
}